{
    "0808/0808.3221_arXiv.txt": {
        "abstract": "The precise localization of short/hard (Type I) gamma-ray bursts (GRBs) in recent years has answered many questions but raised even more. I present some results of a systematic study of the optical afterglows of long/soft (Type II) and short/hard (Type I) GRBs, focusing on the optical luminosity as another puzzle piece in the classification of GRBs. ",
        "introduction": "In the last decade, since the precise localization via the discovery of their optical \\citep{vanParadijs970228}, X-ray \\citep{Costa970228} and radio \\citep{Frail970508} afterglows, much knowledge has been gained about the high-energy transients known as gamma-ray bursts, which we now know are extremely energetic \\citep{Kulkarni990123} explosions of massive stars \\citep{WoosleyBloom} in cosmological distances \\citep{Metzger970508}. But there are two classes of GRBs \\citep{Kouveliotou1993}, and those of short duration remained enigmatic until the advent of the dedicated GRB mission \\emph{Swift} \\citep{GehrelsSwift}. This satellite made it possible to also pinpoint short GRBs via their afterglows in the X-ray \\citep{Gehrels050509B}, optical \\citep{Hjorth050709, Fox050709} and radio \\citep{Berger050724} regimes, establishing not only that they too originate at cosmological distances, but also that at least some of them are not associated with recent star formation \\citep{Berger050724, Barthelmy050724}, implying that they must originate from different progenitors than long/soft GRBs. The most favored model is the coalescence of two compact objects of which at least one is a neutron star \\citep[for reviews, see][]{NakarReview, LeeRR}.  The discovery of two temporally long GRBs with no associated supernova emission down to deep limits \\citep{Fynbo060614, Gal-Yam060614, DellaValle060614, Ofek060505} exacerbated the debate on GRB classification \\citep{Gehrels060614} which had been triggered by GRB 050724, a GRB in an elliptical galaxy which was temporally long as seen by \\emph{Swift}, but would have been classified as short if detected by BATSE, due to the presence of a long, soft, X-ray-bright tail of emission. This led \\cite{Zhang060614}, in the light of classifying GRBs according to the underlying progenitor physics \\citep{BloomPhysics}, to create a new classification based on several observed quantities and phenomena, such as duration, spectral lag \\citep{NorrisBonnell}, light curve shape and host galaxy characteristics. GRBs that originate in the explosions of massive stars are called Type II events, whereas those unassociated with recent star formation (probably due to compact object coalescence) are Type I GRBs.  By the end of 2007, the \\emph{Swift} mission had resulted in a plethora of (optical) afterglow discoveries, allowing the creation of a large sample of Type II GRB afterglows, as well as a first moderately large sample of Type I GRB afterglows. Since several years, I and my collaborators have been systematically studying large samples of optical afterglows. In this article, I report results from two recent papers, the first one \\citep{KannPaperIV} comparing the afterglows of \\emph{Swift}-era Type II GRBs with those of the pre-\\emph{Swift} era \\citep{KannPaperIII}, the second one comparing the afterglows of Type I GRBs to the complete Type II GRB afterglow sample \\citep{KannPaperV}. I focus especially on three enigmatic events, GRBs 060121, 060505 and 060614, which point to problems in the Type I/II classification scheme, in the light of the luminosities of their optical afterglows. ",
        "conclusions": "I have presented the results of two studies on Type I and Type II GRB afterglows in the \\emph{Swift} era, with a special focus on three enigmatic and hard-to-classify GRBs in the light of their optical afterglow luminosities. In none of the three cases is the afterglow luminosity a decisive factor in determining the GRB progenitor class, but they still offer a further puzzle piece toward a deeper understanding of GRBs and their progenitors.  \\begin{theacknowledgments} D.A.K. acknowledges B. Zhang for the invitation to the Nanjing conference as well as comments, S. Klose for all the financial support, and the contribution of the many authors who had a part in creating the two papers that this proceeding is based upon. Also thanks to Y. F. Huang for many an excellent information both before and after the conference. \\end{theacknowledgments}   \\doingARLO["
    },
    "0808/0808.3401_arXiv.txt": {
        "abstract": "The abundance and structure of dark matter subhalos has been analyzed extensively in recent studies of dark matter-only simulations, but comparatively little is known about the impact of baryonic physics on halo substructures. We here extend the {\\small SUBFIND} algorithm for substructure identification such that it can be reliably applied to dissipative hydrodynamical simulations that include star formation. This allows, in particular, the identification of galaxies as substructures in simulations of clusters of galaxies, and a determination of their content of gravitationally bound stars, dark matter, and hot and cold gas. Using a large set of cosmological cluster simulations, we present a detailed analysis of halo substructures in hydrodynamical simulations of galaxy clusters, focusing in particular on the influence both of radiative and non-radiative gas physics, and of non-standard physics such as thermal conduction and feedback by galactic outflows. We also examine the impact of numerical nuisance parameters such as artificial viscosity parameterizations.  We find that diffuse hot gas is efficiently stripped from subhalos when they enter the highly pressurized cluster atmosphere. This has the effect of decreasing the subhalo mass function relative to a corresponding dark matter-only simulation. These effects are mitigated in radiative runs, where baryons condense in the central subhalo regions and form compact stellar cores. However, in all cases, only a very small fraction, of the order of one percent, of subhalos within the cluster virial radii preserve a gravitationally bound hot gaseous atmosphere. The fraction of mass contributed by gas in subhalos is found to increase with the cluster-centric distance. Interestingly, this trend extends well beyond the virial radii, thus showing that galaxies feel the environment of the pressurized cluster gas over fairly large distances. The compact stellar cores (i.e.~galaxies) are generally more resistant against tidal disruption than pure dark matter subhalos. Still, the fraction of star-dominated substructures within our simulated clusters is only $\\sim 10$ per cent. We expect that the finite resolution in our simulations makes the galaxies overly susceptible to tidal disruption, hence the above fraction of star--dominated galaxies should represent a lower limit for the actual fraction of galaxies surviving the disruption of their host dark matter subhalo. ",
        "introduction": "\\label{sec:intro}  The hierarchical cold dark matter (CDM) model has been established as the standard model of cosmic structure formation and its parameters are now tightly constrained by a variety of observations \\citep[e.g.][\\ \\ and references therein]{2008arXiv0803.0547K, 2006Natur.440.1137S}. Within the CDM scenario, galaxy clusters are the largest gravitationally bound objects in the universe, and they exhibit rich information about the process of structure formation. As a result, they attract great interest both from observational and theoretical points of view. Due to their recent formation, clusters of galaxies also represent the ideal places where the hierarchical assembly of structures is `caught in the act'. Indeed, the complexity of their internal structure reflects the infall of hundreds, if not thousands, of smaller objects and their subsequent destruction or survival within the cluster potential. The dynamical processes determining the fate of the accreting structures also provide the link between the internal dynamics of clusters and the properties of their galaxy population.  Modern cosmological simulations provide an ideal tool to study the non-linear dynamics of these processes in a cosmological environment.  Thanks to the high resolution reached by dark matter (DM) only simulations, a consistent picture of the population of subhalos within galaxy clusters has now emerged \\citep[e.g., ][]{1999ApJ...524L..19M,2000ApJ...544..616G,2001MNRAS.328..726S,2002MNRAS.335L..84S, 2003MNRAS.345.1313S,2004MNRAS.352..535D,2004MNRAS.352L...1G,2004MNRAS.348..333D,2004ApJ...609...35K}. Although these simulations give a highly detailed description of the evolution of DM substructures, they can not provide information on how gas-dynamical processes affect the properties and the dynamics of substructures. Yet such an understanding is needed to make a more direct link of the properties of dark matter substructures with those of the galaxies that orbit in clusters. For example, we would like to know the timescale over which subhalos can retain their diffuse gas once they infall into a cluster, as this also determines the time interval over which gas can continue to cool and to refuel ongoing star formation.  Being highly pressurized, the intra--cluster medium (ICM) is quite efficient in stripping gas from merging structures (\\citealt{1972ApJ...176....1G}; see also \\citealt{2008MNRAS.383..593M} and references therein), thereby altering both the mass of the subhalos and their survival time. The stripping also governs the timescale over which galaxy morphology and colours are modified inside galaxy clusters \\citep{1999MNRAS.308..947A,2004AJ....127.3361K}.  Finally, although baryons are expected to only play a sub-dominant role for the dynamics of substructures due to their limited mass fraction, they may nevertheless alter the structure of subhalos in important ways, affecting their survival, mass loss rate, abundance and radial distribution. Quantifying these differences relative to dark matter-only simulations is an important challenge for the present generation of cosmological hydrodynamical simulations.  So far there have only been a limited number of investigations of the properties of sub-halos within hydrodynamical simulations.  Based on cluster simulationsw with the adaptive mesh refinement code ART, \\citet{2005ApJ...618..557N} have demonstrated that the stellar mass of infalling galaxies is approximately conserved and the survival of the galaxies is slightly enhanced. Similar studies using the Tree-SPH code GASOLINE have been performed by \\citet{2006MNRAS.366.1529M}, where twice as many sub-halos within the central part of a Galaxy sized halo were found when compared with pure dark matter simulations. Finally \\citet{2008ApJ...678....6W} used a TreeSPH code to investigate the halo occupation and sub-halo correlation functions for one group-sized halo, also pointing out the enhanced survival of sub-structures in the hydrodynamical runs.  Here we present a systematic analysis of a large sample of galaxy clusters, with an emphasis on the role of non-standard physics and the details of feedback prescriptions on the survival of galaxies. Additionally, we also present an analysis of the evolution of the diffuse gas within the galaxies. In general our findings are in broad agreement with those in previous studies, but we stress that our star dominant systems are still not as numerous as expected from semi-analytic modelling to reproduce the cluster luminosity functions (e.g., \\citet{2007MNRAS.375....2D}; see also \\citet{2006RPPh...69.3101B} for a review on semi-analytic models of galaxy formation).  Thanks to improvements in numerical codes and higher computer performance, cosmological hydrodynamical simulations of galaxy clusters can now model a number of the most important physical processes in galaxy formation, such as radiative cooling, star formation, and feedback in energy and metals, while at the same time reaching sufficient resolution to treat these processes in a numerically robust and physically meaningful way \\citep[e.g., ][\\ , for a recent review]{2008SSRv..134..269B}. This allows more realistic structure formation calculations of the coupled system of dark matter and baryonic gas than possible with dark matter only simulations, far into the highly non-linear regime.  For instance, radiative cooling has the effect of letting a significant fraction of the baryons cool at the halo centres, which alters the shape and concentration of halos through the mechanism of adiabatic contraction \\citep[e.g.,][]{2004ApJ...616...16G}. At the same time, the condensation of cooled baryons and their subsequent conversion into collisionless stars is expected to largely protect them from stripping, reducing the rate of disruption of substructures and increasing their survival times. Indeed, assuming that galaxies at least survive for a while the disruption of their host DM halos has been shown to be crucial for semi--analytical models of galaxy formation to provide a successful description of the observed galaxy population in clusters \\citep[e.g., ][]{2004MNRAS.352L...1G}.  This discussion highlights the importance of understanding the properties and the evolution of cluster substructures as predicted by modern cosmological hydrodynamical simulations. Evidently, an important technical prerequisite for such an analysis is the availability of suitable numerical algorithms to reliably identify substructure in dissipative simulations. To this end we introduce a modified version of the {\\small SUBFIND} algorithm, and test its robustness.  We then apply it in a detailed analysis of the properties of substructures in a large ensemble of galaxy clusters, which have been simulated both at different resolutions in their DM-only version, and in several further versions where different physical processes where included, such as non-radiative and radiative hydrodynamics, star formation, energy feedback, gas viscosity and thermal conduction. The aim of this analysis is to quantify the numerical robustness of the measured properties of substructures, and to identify the quantitative influence of these physical processes on substructure statistics in comparison with dark matter only models.  The paper is organized as follows. We describe in Section~\\ref{sec:sim} the sample of galaxy clusters and the simulations performed.  Section~\\ref{sec:detection} provides a description of the algorithm used to identify gravitationally bound substructures. This algorithm is a modified version of SUBFIND \\citep{2001MNRAS.328..726S}, which we suitably changed to allow extraction of bound structures from N-Body/SPH simulations also in the presence of gas and star particles.  In Section~\\ref{sec:sub}, we present our results about the properties of the substructures in both DM and hydrodynamical simulations. We summarize our results and outline our main conclusion in Section \\ref{sec:conc}. An Appendix is devoted to the presentation of tests of resolution and numerical stability of our analysis. ",
        "conclusions": "\\label{sec:conc}  In this study, we presented an analysis of the basic substructure properties in high--resolution hydrodynamical simulations of galaxy clusters carried out with the TreePM-SPH code {\\small GADGET}. We used a fairly large set of 25 clusters, with virial masses in the range $1-33 \\times 10^{14}\\msun$, out of which eight clusters have masses above $10^{15}\\msun$. The identification of the substructures has been carried out with the {\\small SUBFIND} algorithm \\citep{2001MNRAS.328..726S}, which we suitably modified to account for the presence of gas and star particles in hydrodynamical simulations. The primary aim of our analysis was to quantify the impact that gas physics has on the properties of subhalos, especially on their abundance. Note that previous work has so far almost exclusively analyzed substructures in dark matter-only simulations, so that our study is one of the first attempts to directly compare the DM-only results to those of hydrodynamical simulations that also account for baryonic physics.   We analyzed non--radiative hydrodynamical simulations as well as radiative simulations, including star formation and different efficiencies for energy feedback associated with galactic outflows. Also, we examined the sensitivity of our results with respect to different parameterizations for the artificial viscosity needed in SPH. The main results of our analysis can be summarized as follows.  \\begin{itemize} \\item Consistent with previous studies \\citep{2001MNRAS.328..726S,2004MNRAS.348..333D}, we find that the subhalo mass function in DM-only runs converges well for different numerical resolution over the mass range where substructures are resolved with at least $\\simeq 30$ gravitationally bound DM particles. We extend this result to show that the same convergence also holds well for the total subhalo mass function in hydrodynamical runs, irrespective of whether cooling and star formation are included in the simulations.  \\item In non--radiative hydrodynamical simulations, the presence of a gas component leads to a reduction of the total subhalo mass function with respect to DM-only runs. On average, the reduction in mass of the substructures is slightly larger than expected from the complete stripping of the baryonic component. This is due to a combination of the modification of the subhalo orbits arising from the pressure force exerted by the intra--cluster medium. Furthermore, we see indications for an increased susceptibility of the surviving DM halos to tidal disruption.  \\item In general we find that only a very small fraction of subhalos, of order of one percent, within the cluster virial radii maintain a self-bound atmosphere of hot gas. This indicates a high efficiency for the stripping of the hot gas component in our simulations. The radial dependence of the mean gas fraction within subhalos indicates that gas stripping starts already at large cluster--centric distances, beyond the virial radius.  Using a scheme for reduced artificial viscosity has the effect of reducing viscous stripping in the cluster periphery, while increasing the stripping within the virialized high-density region.  \\item In radiative runs that include star formation, some of the baryons are protected from stripping due to their concentration at the centres of the subhalos.  In this case, substructures have masses which are comparable to, or are slightly larger than in the DM-only runs.  The stellar mass fraction is also found to strongly increase for subhalos close to the centre. This confirms that compact clumps of star particles are indeed more resistant against tidal destruction than DM subhalos. Despite this, we find that only a small fraction of subhalos are dominated by their stellar component. As a result, the number of star--dominated galaxies is smaller than expected when, like usually assumed in semi--analytical models of galaxy formation, they are allowed to survive for a dynamical friction time after the disruption of their DM halo. This suggests that our hydrodynamical simulations at their current resolution may not be able yet to accurately follow the survival of galaxies within clusters after the destruction of their DM subhalos. \\end{itemize}  Our study shows that the technical improvements we implemented in {\\small SUBFIND} allow a reliable identification of substructures in hydrodynamical simulations that include radiative cooling and star formation. This allows a direct identification of satellite galaxies not only in clusters of galaxies, but also in future high-resolution hydrodynamical simulations of galaxy-sized halos. This is essential to make direct contact with observational data, and also allows studies of the structural impact of baryonic physics on substructure. As our result show here, gas physics {\\em does alter} substructure statistics in interesting ways, but the effects depend on physical processes such as radiative cooling and feedback. This also means that the approximate self-similarity of substructure properties that has been found in DM-only simulations of clusters and galaxy-size objects \\citep{1999ApJ...524L..19M} is not expected to hold nearly as well in simulations with radiative cooling. For example, in galaxy-size halos, many subhalos will be so small that they do not experience atomic cooling; as a result their behaviour should be close to our non-radiative results, and we would here expect a reduction of their abundance relative to the DM-only case.   In future work, it will be important to further improve the numerical description and resolution of our simulations, in order to be able to more accurately follow in particular the star--dominated substructures (i.e., galaxies) within galaxy clusters. This will then also provide important input for the semi--analytical models of galaxy formation, allowing a check and calibration of their dynamical friction and gas stripping prescriptions."
    },
    "0808/0808.3771_arXiv.txt": {
        "abstract": "\\noindent We present analyses of a 50\\,ks observation of the supergiant X-ray binary system \\ifApJ{Cygnus\\,X-1 (\\mbox{Cyg\\,X-1})/\\linebreak}{Cygnus\\,X-1/}HDE\\,226868 taken~with the \\Chandra{} High Energy Transmission Grating Spectrometer (HETGS). Cyg\\,X-1 was in its spectrally hard state and the observation was performed during superior conjunction of the black hole, allowing for the spectroscopic analysis of the accreted stellar wind along the line of sight. A significant part of the observation covers X-ray dips as commonly observed for \\Cyg{} at this orbital phase, however, here we analyze only the high count rate nondip spectrum. The full 0.5--10\\,keV continuum can be described by a single model consisting of a disk, a narrow and a relativistically broadened Fe K$\\alpha$ line, and a power-law component, which is consistent with simultaneous \\ifApJ{\\textsl{Rossi X-Ray Timing Explorer}}{\\RXTE} broad band data. We detect absorption edges from overabundant neutral O, Ne, and Fe, and absorption line series from highly ionized ions and infer column densities and Doppler shifts. With emission lines of He-like Mg\\,\\textsc{xi}, we detect two plasma components with velocities and densities consistent with the base of the spherical wind and a focused wind. A simple simulation of the photoionization zone suggests that large parts of the spherical wind outside of the focused stream are completely ionized, which is consistent with the low velocities ($<$200\\,km\\,s$^{-1}$) observed in the absorption lines, as the position of absorbers in a spherical wind at low projected velocity is well constrained. Our observations provide input for models that couple the wind activity of HDE\\,226868 to the properties of the accretion flow onto the black hole. ",
        "introduction": "\\label{sec:intro} \\object[Cygnus X-1]{Cygnus\\,X-1\\ifApJ{ (Cyg X-1)}{}} was discovered in 1964 \\citep{Bowyer1965} and soon identified as a high-mass X-ray binary system (HMXB) with an orbital period of 5.6\\,d \\citep{MurdinWebster1971,WebsterMurdin1972,Bolton1972}. It consists of the supergiant O9.7 star \\object[HDE 226868]{HDE\\,226868} \\citep{Walborn1973,Humphreys1978} and a compact object, which is dynamically constrained to be a black hole \\citep{GiesBolton1982}. The detailed spectroscopic analysis of HDE\\,226868 by \\citet{Herrero1995} gives a stellar mass $M_\\star\\approx18\\,M_\\odot$, leading to a mass of $M_\\mathrm{BH}\\sim$10\\,$M_\\odot$ for the black hole, if an inclination $i\\approx35^\\circ$ is assumed. Note that \\citet{Ziolkowski2005} derives a mass of $M_\\star=(40\\pm5)\\,M_\\odot$ from the evolutionary state of HDE\\,226868, corresponding to $M_\\mathrm{BH}=(20\\pm5)\\,M_\\odot$, while \\citet{Shaposhnikov2007} claim $M_\\mathrm{BH}=(8.7\\pm0.8)\\,M_\\odot$ from X-ray spectral-timing relations.  \\Cyg{} is usually found in one of the two states that are distinguished by the soft X-ray luminosity and spectral shape, the timing properties, and the radio flux \\citep[see, e.g.,][]{Pottschmidt2003,Gleissner2004_III,Gleissner2004_II,Wilms2006}: the low/hard state is characterized by a lower luminosity below 10\\,keV, a hard Comptonization power-law spectrum (photon index $\\Gamma\\sim1.7$) with a cutoff at high energies (folding energy $E_\\mathrm{fold}\\sim150\\,$keV) and strong variability of $\\sim$30\\% root mean square (rms). Radio emission is detected at the $\\sim$15\\,mJy level. In the high/soft state, the soft X-ray spectrum is dominated by a bright and much less variable (only few~\\% rms) thermal disk component, and the source is invisible in the radio. Within the classification of \\citet{RemillardMcClintock2006}, the high/soft state of \\Cyg{} corresponds to the steep power-law state rather than to the thermal state, as a power-law spectrum with photon index $\\Gamma\\sim2.5$ may extend up to $\\sim$10\\,MeV \\citep{Zhang1997,McConnell2002,CadolleBel2006}. Most of the time, \\Cyg{} is found in the hard state, but transitions to the soft state and back after a few weeks or months are common every few years. Transitional or intermediate states \\citep{Belloni1996} are often accompanied by radio and/or X-ray flares. Similar to a transition to the soft state, the spectrum softens during these flares and the variability is reduced. This behavior is called a ``failed state transition'' if the true soft state is not reached \\citep{Pottschmidt2000,Pottschmidt2003}. Transitional states have occurred more frequently since mid-1999 than before \\citep{Wilms2006}, which might indicate changes in the mass-accretion rate due to a slight expansion of HDE\\,226868 \\citep{Karitskaya2006}.  HMXBs are believed to be powered by accretion from the stellar wind. The accretion rate and therefore X-ray luminosity and spectral state are thus very sensitive to the wind's detailed properties such as velocity, density, and ionization. For HDE\\,226868, \\citet{Gies2003} found an anticorrelation between the H$\\alpha$ equivalent width (an indicator for the wind mass loss rate $\\dot M_\\star$) and the X-ray flux. Considering the photoionization of the wind would allow for a self-consistent explanation \\citep[see, e.g.,][]{Blondin1994}: a \\emph{lower} mass loss gives a \\emph{lower} wind density and therefore \\emph{higher} degree of ionization due to the irradiation of hard X-rays, which reduces the driving force of HDE\\,226868's UV photons on the wind and results in a \\emph{lower} wind velocity $v$, leading finally to a \\emph{higher} accretion rate \\citep[$\\propto\\dot M_\\star/v^4$,][]{BondiHoyle1944}. However, \\citet{Gies2008} find suggestions that the photoionization and velocity of the wind might be similar during both hard and soft states. UV observations allow the photoionization in the HDE\\,226868\\,/\\,\\Cyg{} system to be probed: \\citet{Vrtilek2008} reported P\\,Cygni profiles of \\Ion{N}{v}, \\Ion{C}{iv}, and \\Ion{Si}{iv} with weaker absorption components at orbital phase $\\phi_\\mathrm{orb}\\approx0.5$, i.e., when the black hole is in the foreground of the supergiant. This reduced absorption, which was already found by \\citet{Treves1980}, is due to the \\citet{HatchettMcCray1977} effect, showing that those ions become superionized by the X-ray source. \\citet{Gies2008} model the orbital variations of the UV lines assuming that the wind of HDE\\,226868 is restricted to the shadow wind from the shielded side of the stellar surface \\citep{Blondin1994}, i.e., the \\citet{Stromgren1939} zone of \\Cyg{} extends to the donor star. However, this assumption applies only to the spherical part of the wind, which might therefore hardly contribute to the mass accretion of \\Cyg{}.  As HDE\\,226868 is close to filling its Roche lobe \\citep{Conti1978,GiesBolton1986_II,GiesBolton1986_III}, the wind is not spherically symmetrical as for isolated stars, but strongly enhanced toward the black hole \\citep[``focused wind\\ifApJ{;''}{'';}][]{FriendCastor1982}. The strongest wind absorption lines in the optical are therefore observed at the conjunction phases \\citep{Gies2003}. Similarly, X-ray absorption dips occur preferentially around $\\phi_\\mathrm{orb}=0$, i.e., during superior conjunction of the black hole \\citep{BalucinskaChurch2000}. These dips are probably caused by dense, neutral clumps, formed in the focused wind where the photoionization is reduced, although recent analyses have also suggested that part of the dipping activity may result from the interaction of the focused wind with the edge of the accretion disk \\citep{Poutanen2008}.  The photoionization and dynamics of both the spherical and focused winds can also be investigated with the high-resolution grating spectrometers of the modern X-ray observatories \\Chandra{} or \\textsl{XMM-Newton}. As none of the previously reported observations of \\Cyg{} was performed at orbital phase $\\phi_\\mathrm{orb}=0$ and in the hard state, when the wind is probably denser and less ionized than in the soft state, the \\Chandra{} observation presented here allows for the most detailed investigation of the focused wind to date.  The remainder of this paper is organized as follows: in \\Sect{sec:obs}, we describe our observations of \\Cyg{} with \\Chandra{} and the \\textsl{Rossi X-Ray Timing Explorer} (\\RXTE), and how we model CCD pile-up for the \\Chandra-% \\ifApJ{High Energy Transmission Grating Spectrometer (HETGS)}{HETGS} data. We present our investigations in \\Sect{sec:analysis}: after investigating the light curves, we model the nondip continuum and analyze neutral absorption edges and absorption lines from the highly ionized stellar wind -- and the few emission lines from He-like ions, which indicate two plasma components. In \\Sect{sec:discuss}, we discuss models for the stellar wind and the photoionization zone. We summarize our results after comparing them with those of the previous \\Chandra{} observations of \\Cyg. ",
        "conclusions": "\\label{sec:discuss} In the following, we discuss our results and derive constraints on the stellar wind in the accretion region.  \\subsection{Velocity and Density of the Stellar Wind}\\label{sec:winddensity} While a spherically symmetric model for the stellar wind in the HDE\\,226868/\\Cyg{} system can be excluded by observations \\cite[see, e.g.,][]{GiesBolton1986_III,Gies2003,Miller2005,Gies2008}, a symmetric velocity law \\begin{equation} v(r) \\;\\;=\\;\\; v_\\infty \\;\\cdot\\; f(r/R_\\star) \\label{eq:vr} \\end{equation} is usually assumed to obtain a first estimate of the particle density in the wind. The fraction $f$ of the terminal velocity $v_\\infty$ is often parameterized by \\citep[\\Equation~3]{LamersLeitherer1993} \\begin{equation} f(x) \\;\\;=\\;\\; f_0 \\:\\;+\\:\\; \\left(1-f_0\\right) \\cdot \\left(1-1/x\\right)^\\beta \\quad (\\mbox{for } x\\ge1) \\label{eq:f} \\end{equation} with $f_0:=v_0/v_\\infty$, where $v_0$ is the velocity at the base of the wind ($x=1$). The simple model for the radiatively driven wind of isolated stars by \\citet{CastorAbbottKlein1975} is obtained for $\\beta=1/2$. The photoionization of the wind, however, suppresses its acceleration \\citep{Blondin1994}, such that a smaller $f$ (e.g., a larger $\\beta$ within the same model) is required in the Str\\\"omgren zone. \\citet{GiesBolton1986_III} have explained the orbital variation of the 4686\\,\\AA{} He\\,\\textsc{ii} emission line profile of \\Cyg{} with a similar model for the focused wind, where $v_\\infty$, $R_\\star$, and $\\beta$ depend on the angle~$\\theta$ from the binary axis. $\\beta$~was interpolated between 1.60 and 1.05 for $\\theta$ between $0$ and $20^\\circ$. The value $\\beta(\\theta\\!=\\!20^\\circ)=1.05$ is, however, often used for a spherically symmetrical wind as well \\citep{Lachowicz2006,Szostek2007,Poutanen2008}. \\citet{Vrtilek2008} use $f_0=0.01$ to avoid numerical singularities and fit the (relatively low) value $\\beta\\approx0.75$ to their models for UV lines. $v_0$ is likely to be of the order of the thermal velocity of H atoms, which is $(2kT/m_\\mathrm{H})^{1/2}=23\\,$km\\,s$^{-1}$ and corresponds to $f_0=0.011$ for HDE\\,226868's effective temperature $T_\\mathrm{eff}=32$\\ifApJ{,}{\\,}000\\,K and $v_\\infty=2100\\,$km\\,s$^{-1}$ \\citep{Herrero1995}.  \\begin{table}\\centering \\caption{Solutions to \\Eq{eq:nr_numbers} for the Model of \\Eq{eq:f}} \\label{tab:density_solutions} \\begin{tabular}{cccccc} \\hline \\hline $n_\\mathrm{H}(x)$      & $d(x)$ &  $f_0$  & $\\beta$  & $x$      & $2100 \\cdot f(x)$ \\\\ (10$^{10}$\\,cm$^{-3}$) &        & $=v_0/v_\\infty$  & & $=r/R_\\star$ & for $\\beta=\\beta_\\mathrm{max}$ \\\\ \\hline \\eVS 440                    & 200  & 0.005 & $<\\infty$ & 1 & 11\\\\ \\hline \\eVS 110                    &  50  & 0.011 & $\\le1$    & $\\le1.01$ & $\\le41$\\\\ 110                    &  50  & 0.011 & $\\le2$    & $\\le1.08$ & $\\le36$\\\\ 110                    &  50  & 0.011 & $\\le3$    & $\\le1.18$ & $\\le30$\\\\ \\hline \\eVS 22                     &  10  & 0.011 & $\\le1$    & $\\le1.08$ & $\\le180$\\\\ 22                     &  10  & 0.011 & $\\le2$    & $\\le1.29$ & $\\le127$\\\\ 22                     &  10  & 0.011 & $\\le3$    & $\\le1.48$ & $\\le\\phantom{1}95$\\\\ \\hline 4.4                    &   2  & 0.011 & $\\le1$    & $\\le1.36$ & $\\le570$\\\\ 4.4                    &   2  & 0.011 & $\\le2$    & $\\le1.69$ & $\\le368$\\\\ 4.4                    &   2  & 0.011 & $\\le3$    & $\\le1.97$ & $\\le271$\\\\ \\hline \\end{tabular}\\\\ {\\bfseries Note.} \\mbox{The binary separation $a=41\\,R_\\odot$ corresponds to $x_a=a/R_\\star=2.4$.} \\end{table} With the mean molecular weight per H atom, $\\mu\\approx1.4$,\\footnote{% The helium abundance per H atom is $\\approx$10\\% \\citep{Wilms2000}. \\label{footnote:HeAbund}} the~continuity equation $\\dot M_\\star \\;=\\; \\mu\\, m_\\mathrm{H}\\, n_\\mathrm{H}(r) \\:\\cdot\\: 4\\pi r^2\\, v(r)$ gives the following estimate for the hydrogen density profile: \\begin{equation} n_\\mathrm{H}(r) \\;=\\; \\frac{\\dot M_\\star/(\\mu\\,m_\\mathrm{H})}{4\\pi R_\\star^2 v_\\infty} \\:\\cdot d(r/R_\\star) \\;\\;\\mathrm{with}\\;\\; d(x) \\;=\\; x^{-2}/f(x) \\label{eq:nr} \\end{equation} Using the parameters of HDE\\,226868 ($\\dot M_\\star=3\\!\\times\\!10^{-6}\\,M_\\odot\\,\\mathrm{yr}^{-1}$, $R_\\star=17\\,R_\\odot$; \\citealp{Herrero1995}; also summarized by \\citealt{Nowak1999_II}, Table~1), \\Eq{eq:nr} predicts \\begin{equation} n_\\mathrm{H}(r) \\;\\;=\\;\\; 2.2\\!\\times\\!10^{10}\\,\\mathrm{cm}^{-3} \\;\\cdot\\; d(r/R_\\star) \\;\\;. \\label{eq:nr_numbers} \\end{equation} \\Eq{eq:nr} can only be solved within the model of \\Eq{eq:f}% \\mbox{\\ifApJ{---}{ --~}}or~any other model for $f$ in \\Eq{eq:vr} with $f(x)\\ge f_0$\\ifApJ{---}{ -- }for \\begin{equation} d \\;\\;<\\;\\; f_0^{-1} \\quad\\quad\\mbox{and}\\quad\\quad 1 \\;\\;\\le\\;\\; x \\;\\;\\le\\;\\; \\sqrt{1/(f_0d)\\;} \\;\\;. \\label{eq:density_solution} \\end{equation} For $f_0\\approx0.01$, the value $d\\approx190$\\ifApJ{---}{ -- }% which would be required to explain the density\\footnote{% A plasma of fully ionized H and He  contains $\\approx$1.2 electrons per H atom. } $n_\\mathrm{H,1} \\approx 4.2\\!\\times\\!10^{12}\\,$cm$^{-3}$ obtained from the unshifted Mg\\,\\textsc{xi} triplet (\\Sect{sec:HeTriplets})% \\ifApJ{---}{ -- } can never be reached within our model for the continuous spherical wind. This result shows that the density is likely to be overestimated by an $R$-ratio analysis which ignores the strong UV-flux of the O9.7 star. Table~\\ref{tab:density_solutions} lists therefore some solutions to \\Eq{eq:nr_numbers} for much lower densities as well. Due to the additional constraint that the radial velocity is less than 100\\,km\\,s$^{-1}$ (Table~\\ref{tab:MgXIlines}), we suggest that the first emission component stems from close to the stellar surface. Although the simple wind model of \\Eq{eq:f} may not be appropriate in this region, the results of Table~\\ref{tab:density_solutions} are rather insensitive to the assumptions on the velocity law, as a wide range of $\\beta$ values was considered, but mostly depend on the wind's initial velocity $v_0$, for which we have used a reasonable estimate. Note that $f_0$ and $d$ also depend on $v_\\infty$, but their product, which is important in \\Eq{eq:density_solution}, does not.  In spite of the systematic errors of the (absolute) density analysis with the $R$-ratio, we infer that the second plasma component is much denser relative to the first one. As it is seen at a larger redshift of 400--1000\\,km\\,s$^{-1}$, we favor its identification with the focused wind. The two emission components could also be caused coincidentally by dense clumps in the stellar wind \\citep[which are common for O-stars; e.g.,][]{Oskinova2007,LepineMoffat2008}, but the interpretation as a slow base of the wind close to the stellar surface and a focused wind between the accreting black hole and its donor star provides a consistent description of both emission components: an undisturbed wind (with $v_\\infty$ and $f_0$ as above, and $\\beta=1.05\\rightarrow1.6$) would reach a velocity of $v(a)\\approx900\\leftarrow1200$\\,km\\,s$^{-1}$ in the distance of the black hole. The focused wind in the orbital plane would then be detected with a projected velocity $v(a)\\sin i=500\\leftarrow700$\\,km\\,s$^{-1}$ at $\\phi_\\mathrm{orb}=0$. While photoionization reduces the efficiency of acceleration for the spherical wind, the denser focused wind is less strongly affected due to self-shielding and reaches the expected velocity. It is an observational fact that the focused wind has another ionization structure: optical emission lines (H$\\alpha$, He\\,\\textsc{ii} $\\lambda4686$) from ions which only exist at a low ionization parameter have been observed in the focused stream \\citep[see, e.g.,][]{GiesBolton1986_III,Ninkov1987b,Gies2003}. \\mbox{For the spherical wind,} however, \\citet{Gies2008} have conjectured that the part between \\Cyg{} and the donor star might be completely ionized in the soft state. Our observations show that the situation in the hard state may be similar.  \\subsection{Modeling of the Photoionization Zone} \\label{sec:XSTAR} We use the photoionization code \\textsc{XSTAR}~2.1ln7b\\footnote{% See \\url{http://heasarc.gsfc.nasa.gov/lheasoft/xstar/xstar.html}\\ifApJ{}{.}} \\citep{Kallman2001} to model the photoionization zone. As the latter is quite complex (due to the inhomogeneous wind density, which is strongly entangled with the X-ray flux), we do not claim to describe the photoionized wind self-consistently, but only want to derive a first approximation. For an optically thin plasma, the relative population of a given atom's ions is merely a function of the ionization parameters \\begin{equation} \\xi(r) \\;\\;=\\;\\; \\frac{L}{n_\\mathrm{H}\\,r^2} \\;\\;=\\;\\;  \\frac{L_{37}}{n_{13}\\,r_{12}^2} \\;\\;\\mathrm{erg}\\;\\mathrm{cm}\\;\\mathrm{s}^{-1} \\;\\;, \\label{eq:ionizationParameter} \\end{equation} where $L_{37}$ is the ionizing source luminosity above 13.6\\,eV in~$10^{37}\\,\\mathrm{erg}\\;\\mathrm{s}^{-1}$, $n_{13}$ is the hydrogen density in $10^{13}\\,\\mathrm{cm}^{-3}$ and $r_{12}$~is the distance from the source in $10^{12}\\,\\mathrm{cm}$. \\begin{figure}\\centering \\includegraphics[width=0.95\\columnwidth]{f13_color.eps} \\caption{Relative population of the ionization stages of iron as a function of the ionization parameter $\\xi=L/(n_\\mathrm{H}\\,r^2)$, as calculated with \\textsc{XSTAR} for $L_{37}=3.5$ and $n_{13}=0.01$.} \\label{fig:IonizationBalance} \\end{figure} \\begin{table}{\\centering \\ifApJ{}{\\vskip-5mm} \\caption{Ionization Parameters for Peak Ion Populations} \\label{tab:ionParameter} \\begin{tabular}{cccc} \\hline \\hline Ion & log$\\:\\xi$ & FWHM(log$\\:\\xi$) & \\eVS $r_{12}=\\left(L_{37}/n_{13}/10^{\\log\\xi}\\right)^{1/2}$\\\\ \\hline Fe\\,\\textsc{xxvi}  &  2.91  &  0.27  &  0.6--0.8 \\\\ Fe\\,\\textsc{xxv}  &  2.75  &  0.32  &  0.7--1.0 \\\\ Fe\\,\\textsc{xxiv}  &  2.65  &  0.30  &  0.8--1.1 \\\\ Fe\\,\\textsc{xxiii}  &  2.58  &  0.28  &  0.8--1.1 \\\\ Fe\\,\\textsc{xxii}  &  2.51  &  0.31  &  0.9--1.3 \\\\ Fe\\,\\textsc{xxi}  &  2.44  &  0.36  &  0.9--1.4 \\\\ Fe\\,\\textsc{xx}  &  2.32  &  0.44  &  1.0--1.7 \\\\ Fe\\,\\textsc{xix}  &  2.12  &  0.40  &  1.2--1.9 \\\\ Fe\\,\\textsc{xviii}  &  2.01  &  0.23  &  1.6--2.0 \\\\ Fe\\,\\textsc{xvii}  &  1.97  &  0.14  &  1.8--2.1 \\\\ \\hline Ca\\,\\textsc{xx}  &  2.69  &  0.56  &  0.7--1.2 \\\\ Ca\\,\\textsc{xix}  &  2.39  &  0.54  &  0.9--1.6 \\\\ \\hline Ar\\,\\textsc{xviii}  &  2.55  &  0.65  &  0.7--1.5 \\\\ Ar\\,\\textsc{xvii}  &  2.22  &  0.50  &  1.0--1.8 \\\\ \\hline S\\,\\textsc{xvi}  &  2.38  &  0.69  &  0.8--1.7 \\\\ S\\,\\textsc{xv}  &  2.08  &  0.43  &  1.2--2.0 \\\\ \\hline Si\\,\\textsc{xiv}  &  2.21  &  0.65  &  0.9--2.0 \\\\ Si\\,\\textsc{xiii}  &  1.96  &  0.39  &  1.5--2.4 \\\\ Si\\,\\textsc{iv}  &  1.09  &  0.18  &  4.7--5.7 \\\\ \\hline Mg\\,\\textsc{xii}  &  2.06  &  0.58  &  1.2--2.3 \\\\ Mg\\,\\textsc{xi}  &  1.82  &  0.31  &  1.8--2.6 \\\\ \\hline Ne\\,\\textsc{x}  &  1.94  &  0.40  &  1.6--2.5 \\\\ Ne\\,\\textsc{ix}  &  1.74  &  0.24  &  2.1--2.8 \\\\ \\hline O\\,\\textsc{viii}  &  1.80  &  0.26  &  1.9--2.6 \\\\ O\\,\\textsc{vii}  &  1.62  &  0.22  &  2.5--3.2 \\\\ \\hline N\\,\\textsc{v}  &  1.51  &  0.01  &  3.25--3.30 \\\\ \\hline C\\,\\textsc{iv}  &  1.48  &  0.19  &  3.3--4.1 \\\\ \\hline \\end{tabular}\\\\} ~\\\\ {\\bfseries Notes.} We list the range in ionization parameter and distance for the maximum ion population (\\Fig{fig:IonizationBalance}) obtained in an \\textsc{XSTAR} simulation with $L_{37}=3.5$ and $n_{13}=0.01$. Note that the binary separation in units of $10^{12}\\,$cm is $a_{12}=2.9$. \\end{table}  For \\Cyg{} as observed in the hard state, extrapolation of the unabsorbed model obtained by our analysis (Table~\\ref{tab:continuum}) gives $L_{37}\\approx3.5$. Although there are obviously strong variations in the density, the \\textsc{XSTAR} calculation has been performed with constant $n_{13}=0.01$, which is the average of $n_\\mathrm{H}(r)$ from \\Eq{eq:nr_numbers} for $R_\\star\\le r\\le a$. \\Fig{fig:IonizationBalance} shows the resulting population of Fe ions. The ionization parameters at which the population distributions of some ions of interest peak, as well as the corresponding full width at half maximum (FWHM), are presented in Table~\\ref{tab:ionParameter}. We have also included Si\\,\\textsc{iv}, N\\,\\textsc{v}, and C\\,\\textsc{iv} for a comparison with the work of \\citet{Vrtilek2008} and \\citet{Gies2008}. The H- and He-like ions detected with \\Chandra{} only appear at considerable distance from the X-ray source. Taking into account that the photons emanating from \\Cyg{} first propagate through a lower density wind and that the (eventually stalled) wind of higher density is only reached in the vicinity of the star, the actual distances will be even larger than the $r_{12}$ values quoted in Table~\\ref{tab:ionParameter}.  \\subsection{Origin of Redshifted X-Ray Absorption Lines} \\begin{figure}\\centering \\includegraphics[width=0.95\\columnwidth]{f14.eps} \\caption{Gray-scale image of $v_\\mathrm{rad,BH}(\\vecr)$ of \\Eq{eq:proj_velocity}, i.e., projected wind velocity against the black hole. A spherically symmetric velocity law given by \\Eqs{eq:vr} and (\\ref{eq:f}) with $v_\\infty=2100\\,\\mathrm{km\\,s}^{-1}$, $f_0=0.01$, and $\\beta=1.05$ is assumed. The star and the black hole are shown as filled black circles; the size of the latter is the accretion radius $2GM_\\mathrm{BH}/v(a)^2$. On the black circle passing through the center of the star and through the black hole $v_\\mathrm{rad,BH}(\\vecr)$ is $0$ since $\\alpha(\\vecr)=90^\\circ$. \\emph{\\mbox{Positive} velocities} (+) are seen only \\emph{within this circle}; a lighter gray means a larger redshift. \\emph{Negative velocities} ($-$) can likewise only occur \\emph{outside of this circle}; a lighter gray here means a larger blueshift. \\ifApJ{Lines of {\\color{red}the}}{The dashed gray lines of} constant \\ifApJ{$v$}{$v_\\mathrm{rad,BH}$} are shown from $-1400\\,$km\\,s$^{-1}$ to $+1000$\\,km\\,s$^{-1}$ in steps of 200\\,km\\,s$^{-1}$. The two highlighted sectors contain all observable lines of sight toward the black hole (after rotation in this plane) for the inclination $i=35^\\circ$. The labels show the corresponding orbital phases. } \\label{fig:proj_velocity} \\end{figure} Additional constraints on the accretion flow can be derived by considering the Doppler shifts observed in absorption lines (Sections~\\ref{sec:absorptionLines} and \\ref{sec:absorptionLineSeries}). Models for the wind velocity like that of \\Eqs{eq:vr} and (\\ref{eq:f}) predict a velocity $\\vecv(\\vecr)$ at the position $\\vecr$ in the stellar wind. But only the projection of $\\vecv$ against the black hole, $v_\\mathrm{rad,BH}$, can eventually be observed as radial velocity in an X-ray absorption line. With the angle $\\alpha(\\vecr)$ between $\\vecv(\\vecr)$ and the direction from $\\vecr$ toward the black hole, $v_\\mathrm{rad,BH}(\\vecr)$ is \\begin{equation} v_\\mathrm{rad,BH}(\\vecr) \\;\\;=\\;\\; \\cos\\alpha(\\vecr) \\;\\cdot\\; |\\vecv(\\vecr)| \\;\\;. \\label{eq:proj_velocity} \\end{equation} Assuming a velocity field with radial symmetry with respect to the star, $\\alpha(\\vecr)$ is just the angle at $\\vecr$ between the star and the black hole. For example, $\\alpha$ is $90^\\circ$ (and $v_\\mathrm{rad,BH}$ is therefore 0) on the sphere containing the center of the star and the black hole diametrically opposed. Redshifted absorption lines can only be observed from wind material inside this sphere, while the part of the wind outside of it is always seen at a blueshift -- a fact which is independent of the assumed velocity field as long as it is directed radially away from the star.  For a spherically symmetrical wind model, any line of sight can be rotated to an equivalent one in a half-plane limited by the binary axis. Only a sector with half-opening angle~$i$ can be observed unless the system's inclination is $i=90^\\circ$. \\Fig{fig:proj_velocity} shows the projected velocity $v_\\mathrm{rad,BH}(\\vecr)$ for a simple wind model and two sectors containing all observable lines of sight toward \\Cyg{} for an inclination of $i=35^\\circ$.  We now investigate the region where the projected wind velocity (\\Fig{fig:proj_velocity}) is compatible with the observed Doppler shifts (Tables~\\ref{tab:H-He-like_velocityShifts} and \\ref{tab:lineSeries}). We are confident that this method allows for the identification of the absorption regions, as the low observed velocities are always found close to the $\\alpha=90^\\circ$ sphere -- independent of the wind model. For most of the investigated ions (\\Ion{Ne}{x}, \\Ion{Mg}{xii}, \\Ion{Si}{xiv} and \\textsc{xiii}, \\Ion{S}{xv}, \\Ion{Ar}{xvii}, \\Ion{Fe}{xix} and \\textsc{xviii}) the inferred distance from the black hole agrees with the predictions of Table~\\ref{tab:ionParameter}, e.g., the projected velocity $-117\\,$km\\,s$^{-1}\\le v_\\mathrm{Ne\\,X}\\le -69\\,$km\\,s$^{-1}$ measured for \\Ion{Ne}{x} is (during $0.93\\le\\phi_\\mathrm{orb}\\le1$) obtained at $r_{12}=1.78\\pm0.07$. For many other ions, both results are still very similar: e.g., \\Ion{Ne}{ix} with $-214\\,$km\\,s$^{-1}\\le v_\\mathrm{Ne\\,IX}\\le -123\\,$km\\,s$^{-1}$ is expected at $r_{12}=1.95\\pm0.07$ from the $v_\\mathrm{rad,BH}$ model and at $r_{12}=2.1$--2.8 from the photoionization model. Only for the highly ionized iron lines, the small distances are -- as already anticipated in \\Sect{sec:XSTAR} -- underestimated by the \\textsc{XSTAR} simulation run with constant average density, which overestimates the wind density close to the black hole. For example, the velocity range $-9\\,$km\\,s$^{-1}\\le v_\\mathrm{Fe\\,XXIV}\\le52\\,$km\\,s$^{-1}$ measured for \\Ion{Fe}{xxiv} corresponds to $r_{12}=1.51\\pm0.07$, while the population of this ion peaks at $r_{12}=0.8$--1.1 within the model presented in Table~\\ref{tab:ionParameter}.  \\subsection{Previous \\Chandra{} Observations}\\label{sec:prevObs} \\begin{figure}\\centering \\includegraphics[width=0.95\\columnwidth]{f15.eps} \\caption{Orbital phase and mean ASM count rate of the previously reported \\Chandra{} observations of \\Cyg{} (ObsID~3814 is presented in this paper). } \\label{fig:observations} \\end{figure} In previous \\Chandra-HETGS observations of \\Cyg, the stellar wind was seen under different viewing angles, or the X-ray source was in other spectral states (see \\Fig{fig:observations}), which probably means that the properties of the wind were also different. \\citet{Schulz2002} have analyzed the 14\\,ks \\ifApJ{}{observation }ObsID~107 performed in 1999 October (at $\\phi_\\mathrm{orb}\\approx0.74$),\\footnote{% Note that \\citet{Schulz2002} and \\citet{Miller2005} quote an erroneous date and thus the wrong orbital phase $\\phi_\\mathrm{orb}=0.93\\ne0.74$ for this observation.} when \\Cyg{} was in a transitional state. They derive a neutral column density of $N_\\mathrm{H}=6.2\\!\\times\\!10^{21}\\:\\mathrm{cm}^{-2}$ from prominent absorption edges (see also Table~\\ref{tab:neutralAbundances}) and detect some emission and absorption lines with indications of P~Cygni profiles. \\citet{Miller2005} have investigated the focused wind with the 32\\,ks \\ifApJ{}{observation }ObsID~2415 of \\Cyg{} in an intermediate state, performed in 2001 January (at $\\phi_\\mathrm{orb}\\approx0.77$). They report absorption and emission lines of H- and He-like resonance lines of Ne, Na, Mg, and Si with a mean redshift of $\\approx$100\\,km\\,s$^{-1}$, as well as some lines of highly ionized Fe and Ni. The column density is $6.2\\!\\times\\!10^{21}\\:\\mathrm{cm}^{-2}$ \\citep{Miller2002}. \\citet{Marshall2001} describe Ly\\,$\\alpha$ and He\\,$\\alpha$ absorption lines, redshifted by $(450\\pm150)\\:\\mathrm{km\\,s}^{-1}$ from the 12.6\\,ks \\ifApJ{}{observation }ObsID~1511 of \\Cyg{} in the hard state, performed in 2001 January (at $\\phi_\\mathrm{orb}\\approx0.83$). \\citet{ChangCui2007} report dramatic variability in the 30\\,ks \\ifApJ{}{observation }ObsID~3407 performed in 2001 October (at $\\phi_\\mathrm{orb}\\approx0.88$), when \\Cyg{} reached its soft state. While a large number of absorption lines (mostly redshifted, but not by a consistent velocity) is identified in the first part of their observation, most of them weaken significantly or cannot be detected at all during the second part. The complete ionization of the wind due to a sudden density decrement is given as a possible explanation. \\citet{Feng2003} detect asymmetric absorption lines with the 26\\,ks \\ifApJ{}{observation }ObsID~3724 of \\Cyg{} in the soft state, performed in 2002 July (at $\\phi_\\mathrm{orb}=0$). The line centers are almost at their rest wavelengths, but the red wings are more extended, especially for the transitions of highest ionized ions, which they explain by the inflowing focused wind reaching both the highest redshift and ionization parameter closest to the black hole.  The interpretation of the absorption lines presented in the previous section describes our observation consistent with wind and photoionization models. It can, however, not be applied to all of the other observations; the model of \\Fig{fig:proj_velocity} neither predicts a conspicuously high redshift at $\\phi_\\mathrm{orb}\\approx0.83$ and at a higher soft X-ray flux (ObsID~1511), nor a positive redshift at $\\phi_\\mathrm{orb}\\approx0.77$ at a still higher flux (ObsID~2415, see \\Fig{fig:observations}). Inhomogeneities (e.g., density enhancements in shielding clumps) and asymmetries of the wind (e.g., due to noninertial forces in the binary system) may therefore play an important role in these cases.  \\subsection{Summary} In this paper, we have presented a \\Chandra-HETGS observation of \\Cyg{} during superior conjunction of the black hole, which allows us to detect the X-ray absorption signatures of the stellar wind of HDE\\,226868. The light curve near $\\phi_\\mathrm{orb}=0$ is shaped by absorption dips; these are, however, excluded from this analysis by selecting nondip times of high count rate only. At a flux of $\\sim$0.25\\,Crab, we have to deal with moderate pile-up in the grating spectra, for which can be accounted very well with the \\texttt{simple\\_gpile2} model. The continuum of both \\Chandra's soft and \\RXTE's broadband X-ray spectrum has been described consistently by a single model, consisting of an empirical broken power-law spectrum with high-energy cutoff (which is typical for the hard state) and a subordinated disk component with an inner radius close to the ISCO. The joint modeling of both spectra reveals the presence of both a narrow and a relativistically broadened fluorescence Fe K$\\alpha$ line. \\Chandra{} has resolved absorption edges of neutral atoms with an overabundance of metals, suggesting an origin not only in the ISM, but also in the stellar wind of the evolved supergiant. The previously suspected anticorrelation of neutral column densities and X-ray flux, which is due to photoionization, is confirmed. Absorption lines of highly ionized ions are produced where the wind becomes extremely photoionized. For H- and He-like ions, Lyman and He-series are detected up to the Ly or He$\\:\\epsilon$ line, which we use to measure the column density and velocity of absorbing ions. For the wealth of Fe L-shell transitions, column densities can best be obtained by directly using a physical model for the complete line series of an ion. The nondip spectrum shows almost no Doppler shifts, probably indicating that we have excluded the focused wind by our selection of nondip times. We have also detected two plasma components in emission by two pairs of i- and f-lines of He-like Mg\\,\\textsc{xi}. The first one is roughly at rest and we identify it with the base of the spherical wind close to the stellar surface. The second plasma component is denser and observed at a redshift that is compatible with the focused wind. A simple \\textsc{XSTAR} simulation indicates that most of the observed ions only exist in a distance of the black hole, where the velocity of a spherical wind, projected against the black hole, is low -- which is a consistent explanation of the small Doppler shifts observed in absorption lines. We review the previously reported \\Chandra{} observations of \\Cyg{} and find that not all of them can be described in the same picture, as the wind may be affected by asymmetries or inhomogeneities. A detailed spectroscopy of the absorption dips, which might shed more light on the focused wind, will be presented in a subsequent paper."
    },
    "0808/0808.1890_arXiv.txt": {
        "abstract": "The orbital parameters of extra-solar planets have a significant impact on the probability that the planet will transit the host star. This was recently demonstrated by the transit detection of HD 17156b whose favourable eccentricity and argument of periastron dramatically increased its transit likelihood. We present a study which provides a quantitative analysis of how these two orbital parameters affect the geometric transit probability as a function of period. Further, we apply these results to known radial velocity planets and show that there are unexpectedly high transit probabilities for planets at relatively long periods. For a photometric monitoring campaign which aims to determine if the planet indeed transits, we calculate the expected transiting planet yield and the significance of a potential null result, as well as the subsequent constraints that may be applied to orbital parameters. ",
        "introduction": "With the number of known extra-solar planets exceeding 300, statistical interpretations of the distribution of orbital parameters are becoming increasingly significant. These parameter distributions help us unlock the mysteries surrounding the planet formation process to which many challanges have been presented, not the least of which contains the mechanisms that drive planetary migration \\citep{arm07}. \\citet*{for08b} showed that transit light curves in particular can be used to characterize orbital eccentricities and hence give further insight into the global eccentricity distribution.  In terms of the sheer number of transit light curves, the major contributors have been the shallow wide-field surveys such as the Transatlantic Exoplanet Survey (TrES) \\citep{man07}, the XO project \\citep{joh08}, the Hungarian Automated Telescope Network (HATNet) \\citep{pal08}, and SuperWASP \\citep{and08}. In addition, there have been at least five cases in which planetary transits were detected through photometric follow-up of planets already known via their radial velocity (RV) discoveries. These five planets are HD 209458b \\citep{cha00,hen00}, HD 149026b \\citep{sat05}, HD 189733b \\citep{bou05}, GJ 436b \\citep{gil07}, and HD 17156b \\citep{bar07a}. The case of HD 17156b is of particular interest since it is a 21.2 day period planet which happens to have a large eccentricity ($e = 0.67$) and an argument of periastron which places the periapsis of its orbit in the direction toward the observer and close to parallel to the line of sight, resulting in an increased transit probability.  Conversely, the dominant sources of RV planet discoveries have been the California \\& Carnegie Planet Search \\citep{mar97} and the High Accuracy Radial velocity Planet Searcher (HARPS) \\citep{pep04} teams. However, in the near future we can expect to see larger-scale surveys \\citep*{kan07b} and new instruments \\citep{li08} which will increase both the number and diversity of known planets. There have been suggestions regarding the strategy for photometric follow-up of these radial velocity planets at predicted transit times \\citep{kan07a} and the instruments that could be used for such surveys \\citep{lop06a}. Some attempts have been made to detect these possible transits \\citep{lop06b,sha06} which have thus far been unsuccessful.  This paper discusses the effect of orbital parameters on the geometric transit probability of planets. We calculate orbital constraints that may be applied, particularly in the absence of transit signatures in photometric follow-up observations. Section 2 describes how the eccentricity and argument of periastron of known planetary orbits affect transit probability. It further presents applications of this effect to known RV planets and discusses how uncertainties in the orbital parameter values affect the reliability of the ephemeris calculations. In Section 3, we show how orbital constraints can be applied in the absence of a photometrically detected transit signal, and we discuss the potential transit yield and statistical significance of a scenario in which no transits are found in a large sample of RV planets. We summarize and conclude in Section 4. ",
        "conclusions": "It is still uncertain at this stage how many of the known radial velocity planets transit their parent stars. What is clear is that the eccentricity distribution of the known exoplanets will increase the transit likelihood, making detections for long-period planets, such as HD 17156b, feasible. We have shown in this paper that there is enough potential amongst longer period planets for transit detections to motivate a photometric monitoring campaign at the predicted times of transit for these targets. \\citet{fle08} have shown that long-period transiting planets may yet be discovered through ground-based transit surveys, particularly if data sets from different surveys are combined.  As pointed out by \\citet{bar07b}, eccentric planets that have a periastron oriented away from the observer are far more likely to exhibit a secondary than a primary eclipse. The detection of such a secondary eclipse is considerably more challenging than for a primary eclipse since it relies on a minimum level of planetary flux and is best pursued at infrared wavelengths. The discussion in \\S \\ref{i_and_e} shows that even an assumption of $\\omega = 3\\pi/2$ can place constraints on the orbital inclination. A prime candidate for such a study is HD 80606b \\citep{nae01} which has a period of 111.87 days and an eccentricity of 0.927. Scaling Figure \\ref{fig1} to this period and eccentricity yields a secondary transit probability of $\\sim 15$\\%.  Many of the results presented in this paper can easily be applied to any system since the results generally scale linearly with the sum of the stellar and planetary radii. Through applying these results to current and future radial velocity planet discoveries, one can choose targets for an efficient observing campaign which may help to discover long-period transiting planets and hence add invaluable information to planetary structure and formation theories."
    },
    "0808/0808.4151_arXiv.txt": {
        "abstract": "We consider the prospects for testing the dark matter interpretation of the DAMA/LIBRA signal with the Super-Kamiokande experiment.  The DAMA/LIBRA signal favors dark matter with low mass and high scattering cross section.  We show that these characteristics imply that the scattering cross section that enters the DAMA/LIBRA event rate determines the annihilation rate probed by Super-Kamiokande.  Current limits from Super-Kamiokande through-going events do not test the DAMA/LIBRA favored region.  We show, however, that upcoming analyses including fully-contained events with sensitivity to dark matter masses from 5 to 10 GeV may corroborate the DAMA/LIBRA signal.  We conclude by considering three specific dark matter candidates, neutralinos, WIMPless dark matter, and mirror dark matter, which illustrate the various model-dependent assumptions entering our analysis. ",
        "introduction": "The DAMA/LIBRA experiment has seen, with $8.2\\sigma$ significance~\\cite{Bernabei:2008yi}, an annual modulation~\\cite{Drukier:1986tm} in the rate of scattering events, which could be consistent with dark matter-nucleon scattering.  Much of the region of dark matter parameter space that is favored by DAMA is excluded by null results from other direct detection experiments, including CRESST~\\cite{Angloher:2002in}, CDMS~\\cite{Akerib:2005kh}, XENON10~\\cite{Angle:2007uj}, TEXONO~\\cite{Lin:2007ka,Avignone:2008xc}, and CoGeNT~\\cite{Aalseth:2008rx}.  On the other hand, astrophysical uncertainties~\\cite{Brhlik:1999tt,Gondolo:2005hh} and detector effects~\\cite{Bernabei:2007hw} act to open up regions that may simultaneously accommodate the results from DAMA and these other experiments.  Following DAMA's latest results, several recent studies~\\cite{Feng:2008dz,Petriello:2008jj,Chang:2008xa,% Fairbairn:2008gz,Savage:2008er} have studied the consistency of DAMA with other direct detection experiments, with varying assumptions and varying conclusions.  What is clear, however, is that if DAMA is seeing dark matter, the preferred region of parameter space has dark matter mass in the range $m_X \\sim 1-10~\\gev$ and spin-independent proton scattering cross section $\\sigmaSI \\sim 10^{-5} - 10^{-2}~\\pb$. Although neutralinos have been proposed as a possible explanation~\\cite{Bottino:2003iu}, such low masses and high cross sections are not typical of weakly-interacting massive particles (WIMPs), and alternative dark matter candidates have been suggested to explain the DAMA signal~\\cite{Smith:2001hy,Feng:2008ya,Feng:2008mu,% Foot:2008nw,Feng:2008dz,Khlopov:2008ty,Andreas:2008xy,Dudas:2008eq}.  The current state of affairs also makes it abundantly clear that data from complementary experiments is likely required to sort out the true nature of this result and is certainly required to establish definitively the detection of dark matter.  Other direct detection experiments may play this role.  In this paper, we note that corroborating evidence may come from a very different source, namely, from the indirect detection of dark matter at Super-Kamiokande (Super-K).  In contrast to direct detection experiments, which rapidly lose sensitivity at low masses, given physical limitations on threshold energies, Super-K's limits remain strong for low masses. Super-K is therefore poised as one of the most promising experiments to either corroborate or exclude many dark matter interpretations of the DAMA/LIBRA data.  In \\secref{relating}, we show that, with a few well-motivated theoretical assumptions, the DAMA and Super-K event rates may be related.  Currently published Super-K results do not challenge the DAMA preferred region.  In \\secref{projection}, however, we show that there is significant potential for Super-K to extend its reach to dark matter masses from 5 to 20 GeV and provide sensitivity that is competitive with, or possibly much better than, direct detection experiments.  In \\secref{models}, we apply our analysis to three specific dark matter candidates that have been proposed to explain DAMA: neutralinos~\\cite{Bottino:2003iu}, WIMPless dark matter~\\cite{Feng:2008ya,Feng:2008dz,Feng:2008mu}, and mirror dark matter~\\cite{Foot:2008nw}.  These candidates illustrate and clarify the assumptions entering the analysis.  We present our conclusions in \\secref{summary}.  As this work was in preparation, a study appeared that also considered testing the dark matter interpretation of DAMA with data from Super-K~\\cite{Hooper:2008cf}. That work focused primarily on a model-independent approach and present Super-K data, considering neutralinos briefly as a case example, whereas this work considers neutralino, WIMPless, and mirror dark matter and emphasizes the much brighter prospects for future Super-K results. ",
        "conclusions": "\\label{sec:summary}  The DAMA/LIBRA signal is currently of great interest, and alternative methods for corroborating or excluding a dark matter interpretation are desired.  In this study, we have shown that the preferred DAMA region implies that Super-K, through its search for dark matter annihilation to neutrinos, has promising prospects for testing DAMA.  We have given a conservative estimate of the projected sensitivity of Super-K.  By using fully contained events, we expect that current super-K bounds may be extended to dark matter masses of 5 to 10 GeV. In the region of most interest for the DAMA result with $m_X \\sim 5-10~\\gev$, the neutralino models of Ref.~\\cite{Bottino:2003iu} and WIMPless models can potentially be tested, provided the sensitivities expected at this low mass range are actually realized.  For mirror matter, however, the mass range of interest is $10-30~\\gev$, and it is unlikely that Super-K can place limits on this model.  We thus have the intriguing prospect that the direct detection result of DAMA/LIBRA could be sharply tested by an indirect detection experiment in the very near future."
    },
    "0808/0808.2754.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  This is a scientific strategy for the detection and characterization of extrasolar planets; that is, planets orbiting other stars. As such, it maps out--over a 15-year horizon--the techniques and capabilities required to detect and measure the properties of planets as small as Earth around stars as large as our own Sun. It shows how the technology pieces and their development fit together to achieve the strategy's primary goal: if planets like Earth exist around stars within some tens of light years of our own Solar System, those planets will be found and their basic properties characterized. Essential to this strategy is not only the search for and examination of individual planets, but also a knowledge of the arrangement, or \\emph{architecture}, of planetary systems around as large a number of stars as possible; this is the second goal of the strategy. The final goal of the strategy is the study of disks around stars, important both to understand the implications of the variety of exoplanet systems for planet formation, and to determine how many nearby stars have environments around them clean enough of debris that planets may be sought and, if found, characterized.  It is important that the program target planets as small as the Earth, if we are to determine whether the conditions we find on our own world are a common outcome of planetary evolution. The only example of a habitable world we have is our own one-Earth-mass planet, and indeed our nearest neighbor, Venus, is nearly the same mass but uninhabitable by virtue of closer proximity to the Sun. Searching for planets, for example, five times the mass of Earth is easier but should they turn out to lack habitable atmospheres we would not know whether this is by chance or a systematic effect of the higher mass.   This goal by no means excludes detection and study of larger planets, which themselves will provide a wealth of discoveries and new knowledge, but the Task Force strongly believes that public interest ultimately lies in knowing whether planets like our own are a common outcome of cosmic evolution. For that reason, we have set the bar high.  The charge to the Exoplanet Task Force (``ExoPTF\"), which conducted the study for the Astronomy and Astrophysics Advisors Committee (``AAAC\"), is reproduced after the bibliography. The ExoPTF membership, listed above, was chosen from the astronomical commu\"nity to represent the range of techniques and expertise involved in exoplanet research today: in instrumentation, observation, and theory. The Task Force met five times from February to September 2007. Representatives of NASA and the NSF were present as \\emph{ex officio} members, and a number of presentations from outside experts were solicited to provide the Task Force with the most up to date information available.  Early in the process, ``white papers\" were solicited on the Task Force's public website, inviting interested individuals to submit ideas for techniques, missions, and theoretical investigations in the area of planet search and detection. Eighty-five white papers were received, read, and discussed by the Task Force. A number of the white papers provided important information that, when verified with outside experts, influenced the ultimate strategy; the sum total of the white papers provided an impressive indication of the breadth and depth of thinking in the astronomical community.  Periodic briefings on the Task Force's progress were provided to the AAAC committee, and general reports were provided to the community at selected scientific conferences and workshops. A draft report was reviewed by the AAAC in October 2007 and a final report delivered in March 2008.  The principal goal of this strategy--the detection of habitable planets--is perhaps among the most challenging and at the same time most rewarding goals in modern science. In a sense, habitable planets are in an ``anti-sweet spot\", wherein the position they occupy around nearby Sun-like stars does not correspond to the highest senstivity zone for any of the techniques discussed here. Further, Earth-sized planets will have extremely low signal-to-noise, making analysis of the detected signatures of these objects extremely difficult. The Task Force was encouraged, however,  by the extraordinary progress made in the last half-decade in detecting planets down to less than ten times the mass of the Earth, in characterizing extrasolar giant planets spectroscopically, in technological progress accomplished in key approaches that will be needed to detect and study Earth-sized planets, and in the innovative and entrepeneurial character of the community in proposing and in some cases executing novel approaches to detection and characterization. It is the vigor of the field and its practitioners that ensures ultimate success in this difficult endeavor.  The Task Force firmly believes that the strategy presented herein is reasonably paced, technically possible, and scientifically compelling. It is robust against surprises both in the areas of science ({\\it e.g.,} frequency of occurrence of planets) and technological development. But the strategy requires a level of stamina not typical of the Federal enterprise. Important and exciting results will come from the near-term portions of the strategy (\\emph{i.e.,} within 5 years), but the ultimate goal of finding and characterizing a planet like our own, around a star like our Sun, will require the full 15-year time horizon for its realization.  \\pagebreak  % ============================================================================ % Begin New PART - SCIENTIFIC AND PHILOSOPHICAL SIGNIFICANCE % ============================================================================  \\chapter{Scientific and Philosophical Significance of Detecting other Earths}  The idea that humankind's home in the cosmos is just one of countless worlds has been favored by many cultures throughout human history.  The thread of western thought on this matter goes back to antiquity, but the view of the cosmos promulgated in medieval Europe put the Earth in a special place and discouraged speculation on the possibility of other inhabited worlds. Nonetheless, the German philosopher and theologian Albertus Magnus wrote in the 13th century: ``Do there exist many worlds, or is there but a single world? This is one of the most noble and exalted questions in the study of Nature.\" Two centuries later, Nicolas of Cusa asserted on a philosophical basis that the heavens were teeming with inhabited worlds.  The notion of a plurality of worlds was put on a scientific footing in the mid-16th century by Copernicus' cosmology that displaced the Earth from the center of the Solar System, so that it became possible to imagine an abundance of unseen worlds in the form of other solar systems, as Giordano Bruno wrote: \\\\  \\begin{quote} \\emph{``There are countless suns and countless earths all rotating around their suns in exactly the same way as the seven planets of our system. We see only the suns because they are the largest bodies and are luminous, but their planets remain invisible to us because they are smaller and non-luminous. The countless worlds in the universe are no worse and no less inhabited than our Earth.\"}\\\\ \\end{quote}  These ideas were not popular in the chilly counter-Reformation environment of the times, and indeed Bruno was burned at the stake in 1600. Less than a decade later, however, Galileo Galilei saw through his telescope first that our Moon--rather than being a perfect sphere--was in fact a mountainous world in its own right, and then that planet-like bodies (moons) orbited Jupiter just as our Moon orbits the Earth. The progressive displacement of humanity's position from the cosmic center became irreversible. Today we know the Earth is one of countless bodies that orbit the Sun; the Sun orbits the center of the Milky Way Galaxy along with $\\simeq 200$ billion other stars; there are tens of billions of galaxies like the Milky Way in the cosmos.  And yet, one special aspect of the Earth remains: of all the planets in the solar system, only Earth demonstrably is teeming with life. If life exists on other planets or moons of the Solar System, it does not have the profound impact on the properties of those bodies that life has on Earth, and for all we know Earth is the only abode of life around the Sun. No world in the solar system is remotely suitable for life as we know it in the way that Earth is; we must look beyond our own Solar System to the stars if we are to determine whether other habitable worlds like the Earth exist. To complete the Copernican revolution, we must detect and characterize ``extra-solar\" planets the size of our home world.  Beginning in the early 1990s, planets around other stars began to be detected, first by their indirect effects on the motion of their parent stars, then directly. Today, over 260 planets are known to exist around other stars, about 10\\% of which can be studied directly. All of these, with just a couple of exceptions, are more than ten times more massive than the Earth, and we could not today detect our own home world in orbit around a Sun-like star a mere 4 light years away--the distance to the closest star Proxima Centauri. But we are close to detecting and characterizing Earth-sized planets around smaller stars, and the means required to do so around a star like the Sun are so well understood that we argue here that the required technologies could be fielded within fifteen years. A strategy that maximizes the chances for success while minimizing risk forms the core of this report.  Complementary to detecting and characterizing individual planets the size of the Earth is determining whether the architecture of our own solar system--rocky planets huddled close to the parent star and multiple giant planets in more distant orbits--is typical. We have enough information today to say that upwards of 15\\% of Sun-like stars harbor giant planets, and that planets down to the size of Neptune also seem to be common, but have no information on the occurrence of planets the mass or size of the Earth. Because the formation of rocky planets like Earth may have been quite different from that of giant planets, extrapolation is risky. The scientific importance of understanding both the statistical occurrence of Earth-sized planets and the architectures of planetary systems is high, because these are crucial to understanding the processes by which planetary systems form. They therefore form a part of our strategy in conjunction with direct detection and characterization, as well as the study of planet forming disks and disk debris. Together, these elements will ensure that, no matter what the ultimate commonality or rarity of Earth-sized planets might be, the end result of the strategy will be a deep understanding of how planets fit into the overall scheme of cosmic evolution.  Humankind stands today at the threshold of answering one of humankind's most ancient questions: is our home world the only suitable abode for life like us in the cosmos? Most humans would answer no, but only because the Copernican mindset is so firmly a part of our modern culture. What if we were rare enough that no such other Earth existed within the reach of modern astronomical telescopes, on the ground or in space? How would we react to the confirmed emptiness the universe presented to us? Would our specialness lead to a profound sense of faith, or a deep loneliness? Would it stiffen humankind's resolve to survive, to solve its political and environmental problems? Would our solitary state lead us to be more mindful of the precious uniqueness of our intellect and awareness?  Conversely, were we to find a planet like the Earth--rocky, similar in size to our world, possessed of an atmosphere-- around nearby stars, our perception of the cosmos around us would change in an equally profound fashion.  We would look up still with wonder, but a handful of stars in our night sky will forever after hold a special place in our imagination, tempting us with wild dreams of flight. Surely that too would make us refocus our energies to hasten the day when our descendants might dare to try to bridge the gulf between two inhabited worlds. \\\\  \\begin{quote} {\\emph{``Which one is it, Mommy?\" asked the older of her two children. They had walked away from the campfire, and gazing now at the familiar pattern of stars in the night sky, a question far different from any ever asked by thousands of generations of human beings drifted off in the cool night air. ``Look at the bright star over there,\" the woman responded to them ``now move your eyes a little to the right, and you'll see that slightly fainter star. The planet belongs to that one. It's almost exactly the size of the Earth, is just a little closer to its sun than we are to ours, and the space telescope that your mommy helped build found oxygen in its atmosphere. That world has air that creatures like us could breathe.\" ``Who lives there?\" asked the younger one. ``No one knows,\" the woman replied, ``but maybe they are looking at us, right now, wondering the same thing.\" }} \\end{quote}  \\pagebreak  % ============================================================================ % Begin New PART - THE KNOWN EXTRASOLAR PLANETS % ============================================================================  \\chapter{The Known Extrasolar Planets}  %\\setcounter{section}{0}  % - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - - ",
        "conclusions": ""
    },
    "0808/0808.3895_arXiv.txt": {
        "abstract": "We have studied the kinematics traced by the water masers located at the centre of the planetary nebula (PN) K3-35, using data from previous Very Large Array (VLA) observations. An analysis of the spatial distribution and line-of-sight velocities of the maser spots allows us to identify typical patterns of a rotating and expanding ring in the position-velocity diagrams, according to our kinematical model. We find that the distribution of the masers is compatible with tracing a circular ring with a $\\simeq$0\\farcs021 ($\\simeq$100~AU) radius, observed with an inclination angle with respect to the line of sight of 55$\\degr$. We derive expansion and rotation velocities of 1.4 and 3.1 km s$^{-1}$, respectively. The orientation of the ring projected on the plane of the sky, at PA$\\simeq$158$\\degr$, is almost orthogonal to the direction of the innermost region of the jet observed in K3-35, suggesting the presence of a disc or torus that may be related to the collimation of the outflow. ",
        "introduction": "High angular resolution observations of molecular gas have revealed the presence of dense equatorial discs and tori towards several late asymptotic giant branch (AGB) stars and young planetary nebulae (PNe): for instance, \\citet{bie88}, \\citet{for98}, \\citet{sah98b}, \\citet{buj03}. The interaction between the post-AGB wind and such equatorial structures has been proposed as one of the possible physical mechanisms in shaping the bipolar and multipolar morphologies seen in PNe and proto-PNe (PPNe) \\citep{bal87,mel95,sok00,ick03}.  The origin of molecular discs and tori around late AGB stars is not completely clear, but a possible explanation is the presence of a binary system \\citep{mor87,liv88,taa89,sok06}. In this case, when one of the stars enters the AGB phase, some of the ejected material can be retained, generating an extended torus. However, in the case of the high velocity jets observed in some late AGB stars, post-AGB stars, and young PNe \\citep{fei85,sah98a,ima02,rie03} it is believed that a more effective collimation mechanism(s) should be present in the innermost region of the object, such as an accretion disc \\citep{mor87,sok94,sok00} or a stellar magnetic field produced by a rotating star \\citep{pas85,che94,gar97}. Furthermore, recent observations reveal compelling evidence for magnetic fields in post-AGB stars and PNe, associated with equatorial discs and/or jets (e.g. \\citealt{sab07,vle06}).  The process for the formation of accretion discs in PPNe has been investigated in detail by \\citet{rey99}. They find that a disc forms when a close binary system (with a substellar companion) undergoes common envelope evolution. For the case in which a low-mass secondary is disrupted during a dynamically unstable mass transfer process, an accretion disc, with a radius of $\\sim$10~AU and a mass of $\\sim$2$\\times10^{-3}$~M$_{\\sun}$, forms within $\\sim$100~yr. Recently, \\citet{rij05} from their 3D simulations based on a two-wind model with a warped disc, suggested that, in order to explain the observed multipolar and point-symmetric shape of PNe, the required discs are quite small ($\\sim$10--100~AU). Furthermore, these disc-like structures should be dense ($10^7-10^8$~cm$^{-3}$) and in Keplerian rotation.  Interferometric CO observations show larger toroidal molecular structures with sizes in the range of $\\simeq$1000--6000 AU in PPNe and young PNe. These structures seem to be systematically in expansion with a mean velocity of $\\sim$7~km~s$^{-1}$, such as in M~1-92 \\citep{buj98}, M~2-9 \\citep{zwe97}, \\mbox{M~2-56} \\citep{cas02}, or KjPn~8 \\citep{for98}. These velocities are comparable to or below those found in expanding AGB envelopes \\citep{hug07}. Importantly, rotation has been observed in the Red Rectangle as well as slower expansion ($\\sim$0.8~km~s$^{-1}$), superimposed on rotation, according to \\citet{buj05}. Note that the sizes of the tori are about one or two orders of magnitude more than the sizes of the disc-like structures proposed by \\citet{rey99} and \\citet{rij05}.  Maps of water maser emission allow the identification of a disc-like structure with a radius of 0\\farcs12 in IRAS 17347-3139 \\citep{deg04}, which corresponds to $\\simeq$100--750~AU at the source distance \\citep{gom05a}. In this case, the gas kinematics suggests the presence of both rotation and expansion in the disc traced by the water masers. In this context, it is very important to study in detail the kinematics of much smaller disc-like structures that might be related to the collimation of the observed bipolar outflows in some PNe.  K3-35 is a young PN that shows a bipolar outflow in optical images \\citep{mir00}. At radio wavelengths, K3-35 exhibits a bright core and two bipolar lobes with an S-shape \\citep{mir01}. The distance to this object has been estimated to be $\\sim$5~kpc \\citep{zha95}, using an statistical method. However, we note the large uncertainty of this type of estimate, since the application of different statistical methods could give distances varying by factors of $\\sim$3 \\citep{phi04}. The characteristic S-shape morphology of the radio lobes can be successfully reproduced by a precessing jet evolving in a dense circumstellar medium \\citep{vel07}. Water maser emission has been found in three regions: two regions located at the tips of the bipolar radio jet about $\\simeq$1$''$ from the centre (regions N and S), and another region towards the core of the nebula within $\\simeq$0\\farcs02 (region C), suggesting the presence of a torus \\citep{mir01}. In addition, OH maser emission has been detected towards the centre of K3-35 (within $\\simeq$0\\farcs04), showing circular polarisation that suggests the presence of a magnetic field \\citep{mir01,gom05b}.  We decided to study the spatio-kinematical distribution of the water masers reported by \\citet{mir01} towards the centre of K3-35 in order to identify possible expansion and/or rotation motions. The paper is organised as follows. In Section 2, we present a simple kinematical model of a ring including both expansion and rotation, and we calculate the pattern delineated in the position-velocity diagrams. We also describe the least-squares fit procedure that we used. In Section 3, we present the current observational data of H$_2$O masers in the PN K3-35. We then apply our model to the H$_2$O masers located towards the centre of the PN K3-35, making a comparison between the results and the observations. Finally, in Section 4, we discuss the implications of our results.  \\section[]{Rotating and Expanding Ring Model} \\subsection{Model} We assume a narrow, uniform, rotating, and expanding ring of radius $R$, arbitrarily oriented with respect to the line of sight. Its projection on the plane of the sky is an ellipse with semimajor and semiminor axes $a$ and $b$, respectively. We define the two frames of reference shown in Fig.~1. Both coincide with the plane of the sky, one of them has its origin at the centre of the ellipse and is oriented such that the $x'$-axis is along the major axis of the projected ellipse, and the other one has the axes parallel to the RA and Dec axes. The semimajor and semiminor axes are related to the ring radius by $a=R$ and $b=R\\cos i$, where $i$ is the inclination angle between the line of sight and the normal to the ring plane as shown in Fig.~2.  The equation of the ellipse is given by \\begin{equation} \\frac{x'^2}{a^2}+\\frac{y'^2}{b^2}=1, \\end{equation} and the transformation equations between the coordinate systems are \\begin{equation} x'=(x-x_0)\\cos\\theta+(y-y_0)\\sin\\theta, \\end{equation} \\begin{equation} y'=-(x-x_0)\\sin\\theta+(y-y_0)\\cos\\theta, \\end{equation} where ($x_0,y_0$) is the position of the centre of the ellipse, and $\\theta$ is the angle between the $x$-axes of the two frames of reference, and is defined as positive clockwise. The angle $\\theta$ is related to the position angle (PA) of the major axis of the ellipse by $\\rmn{PA}=90\\degr-\\theta$.  \\begin{figure} \\includegraphics[width=72mm]{fig1.eps} \\caption{Reference systems. Both the $x'y'$- and the $xy$-coordinate systems are in the plane of the sky. The $x'y'$-system has its origin at the centre of the ellipse ($x_0,~y_0$), and $\\theta$ is the angle between the $x$-axis and the $x'$-axis. The $x$ and $y$ axes are parallel to the RA and Dec axes.} \\end{figure}  \\begin{figure*} \\includegraphics[width=150mm]{fig2.eps} \\caption{Position-velocity diagrams for a rotating and expanding ring model. The top panels correspond to a positive inclination angle and the bottom panels correspond to a negative inclination angle. The sense of rotation (clockwise or counterclockwise) as seen from the observer is indicated. The $x'$ and $y'$ axes are those defined in Fig.~1.} \\end{figure*}  Let $v_s$, $v_{\\rmn{rot}}$, and $v_{\\rmn{exp}}$ be the local standard of rest (LSR) systemic velocity, the rotation velocity, and the expansion velocity of the ring, respectively. Then the observed LSR velocity of a point in the ring can be expressed as \\begin{equation} V_{\\rmn{LSR}}=v_s+\\frac{x'}{a}v_{\\rmn{rot}}\\sin i+\\frac{y'}{a}v_{\\rmn{exp}}\\tan i. \\end{equation} Hence the observed $V_{\\rmn{LSR}}$ will be a linear function of either $x'$ or $y'$, if only one type of motion (rotation or expansion, respectively) is present in the ring \\citep{usc05}. Using equation (1), equation (4) can be written either in terms of the $x'$ or $y'$ coordinate as \\begin{equation} \\frac{[V_{\\rmn{LSR}}-v_s-(x'/a)v_{\\rmn{rot}}\\sin i]^2}{(v_{\\rmn{exp}}\\sin i)^2}+\\frac{x'^2}{a^2}=1, \\end{equation} \\begin{equation} \\frac{[V_{\\rmn{LSR}}-v_s-(y'/a)v_{\\rmn{exp}}\\tan i]^2}{(v_{\\rmn{rot}}\\sin i)^2}+\\frac{y'^2}{(a\\cos i)^2}=1. \\end{equation} Therefore, equations (5) and (6) indicate that the observed $V_{\\rmn{LSR}}$ has a quadratic form (ellipse) expressed in terms of $x'$ or $y'$, when both motions are present.  In this model, we do not solve the radiative transfer through the ring, but we assume that the emission at a given $V_{\\rmn{LSR}}$ comes from the point of the ring having this line-of-sight velocity component. We then use this information to construct position-velocity diagrams.  The orientation of the major axis of the ellipse in the position-velocity ($x'$-$V_{\\rmn{LSR}}$ or $y'$-$V_{\\rmn{LSR}}$) diagrams changes depending on the value of the inclination angle and on whether the sense of rotation is clockwise or counterclockwise as seen from the observer's point of view, as shown in Fig.~2. Accordingly, we are able to distinguish between a positive or negative value of the inclination angle and the sense of the rotation by doing a comparison between the position-velocity diagram delineated by the maser emission and the position-velocity diagrams expected for a rotating and expanding ring. It is important to note that similar position-velocity diagrams can be obtained considering contraction instead of expansion in the ring, but changing the sign of the inclination angle and the sense of rotation. This ambiguity can be solved by constraining the value of the inclination angle (positive or negative) with additional information (see Section 3.2).  \\subsection[]{Least-squares fit} We carried out a least-squares fit of an ellipse to the observed emission, to estimate its spatial distribution on the sky. In order to do this, we considered the curve on the $xy$-plane given by equations (1)--(3) for a given set of values of the parameters $(x_0,y_0)$ (the centre of the ellipse on the plane of the sky), $a$ and $b$ (the semimajor and semiminor axes, respectively) and $\\theta$ (the angle between the $x$-axes of the two frames of reference). We then compute the minimum distances $d_j$ between the position $(x_j,y_j)$ of each of the masers and the ellipse (these distances are measured along straight lines that pass through the maser and intersect the ellipse at right angles to the curve). With these distances, we define the $\\chi^2$ as \\begin{equation} \\chi^2={1\\over {N-5}}\\sum_{j=1}^N \\left({d_j\\over \\sigma_j}\\right)^2\\,, \\label{chi2} \\end{equation} where $N$ is the number of data points (i.e. the factor $N-5$ is the number of degrees of freedom), and $\\sigma_j$ is the error associated with the position of the maser spots. We then find the set of parameters $(x_0,y_0,a,b,\\theta)$ (see above) which give the minimum of $\\chi^2$.  As a first approximation, we began by setting the values of $\\sigma_j$ equal to the observational uncertainties $\\Delta_j$ for the measured positions of the masers, and then finding the ellipse that gives the minimum value of ${\\chi^2}$ (see equation \\ref{chi2}). The actual structure of the maser-emitting region could be a ring of finite width, for which a broad ellipse is a rough approximation of its projection on the plane of the sky. Therefore, we do not expect maser spots to exactly trace an ellipse. We can characterise the width of the ring, assuming that the fitted ellipse traces the mean projected angular distance of the maser emission to the central star, and the actual emission will be distributed around this ellipse, with a dispersion $\\Delta_e$. We treat $\\Delta_e$ as a source of error for the ellipse fit, additional to the measured error, so that ${\\sigma_j}^2={\\Delta_j}^2+{\\Delta_e}^2$. We then try different values of the width parameter ($\\Delta_e$), until we obtain a fit with a minimum ${\\chi^2}(\\Delta_e)=1$. We consider $2 \\Delta_e$ as the characteristic width of the maser ring.  The $(x_0,y_0,a,b,\\theta)$ parameters obtained from the minimization of ${\\chi^2}$ give the best elliptical fit to the observed positions of the masers. With these spatial parameters, we carried out a kinematical fit, using the LSR velocities of the maser components in order to define a ${\\chi^2}_v$ of the form \\[ {\\chi^2}_v={1\\over {N-3}}\\sum_{j=1}^N {1\\over {{\\sigma_v}^2}} \\] \\begin{equation} \\times\\left(V_{\\rmn{LSR}},j-v_s-{{x_j}'\\over a}v_{\\rmn{rot}}\\sin i- {{y_j}'\\over a}v_{\\rmn{exp}}\\tan i\\right)^2\\,, \\end{equation} where $N-3$ is the number of degrees of freedom, and $({x'}_j,{y'}_j)$ are given by equations (2)--(3). Here $\\sigma_v$ is the uncertainty in the observed LSR velocity that we adopt as the spectral resolution of the observations. The minimization of ${\\chi^2}_v$ yielded the best values for the systemic ($v_s$), rotation ($v_{\\rmn{rot}}$), and expansion ($v_{\\rmn{exp}}$) velocities.  \\begin{table*} \\centering \\begin{minipage}{112mm} \\caption{H$_2$O Masers in K3-35 (Region C)} \\label{symbols} \\begin{tabular}{@{}ccrcllccl} \\hline $V_{\\mathrm{LSR}}$ & \\multicolumn{3}{c}{Flux Density} & \\multicolumn{2}{c}{Position\\footnote{ Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds. Data from \\citet{mir01} and \\citet{gom03}.}}   & \\multicolumn{3}{c}{Position Uncertainty\\footnote{Relative position uncertainties (2$\\sigma$) between maser spots. The position of the 1.3~cm continuum emission peak is $\\alpha(\\rmn{J2000})=19^{\\rmn{h}}27^{\\rmn{m}}$44\\fs0233, $\\delta(\\rmn{J2000})=21\\degr30'$03\\farcs441. The relative position uncertainty between the continuum and the H$_2$O masers is 0\\farcs002. The accuracy of the absolute positions is 0\\farcs05.}} \\\\ (km~s$^{-1})$ & \\multicolumn{3}{c}{(mJy)}  &  $\\alpha$(J2000) & $\\delta$(J2000) & \\multicolumn{3}{c}{(arcsec)} \\\\ \\hline 24.6 &  &   23 &  &   19 27 44.0243   & 21 30 03.438   &  &  &  0.010 \\\\ 24.0 &  &   61 &  &   19 27 44.0242   & 21 30 03.428   &  &  &  0.004 \\\\ 23.3 &  &  218 &  &   19 27 44.02246  & 21 30 03.4460  &  &  &  0.0012 \\\\ 22.6 &  & 1010 &  &   19 27 44.022254 & 21 30 03.45146 &  &  &  0.00024 \\\\ 22.0 &  & 1572 &  &   19 27 44.022364 & 21 30 03.45335 &  &  &  0.00014 \\\\ 21.3 &  &  945 &  &   19 27 44.02247  & 21 30 03.4564  &  &  &  0.0003 \\\\ 20.7 &  &  201 &  &   19 27 44.02257  & 21 30 03.4625  &  &  &  0.0011 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "From our model, we conclude that a ring is a likely explanation for the distribution and kinematics shown by the water masers located towards the centre of the PN K3-35. This ring may be arising from the innermost region of a disc or torus probably formed at the end of the AGB phase. Since masers trace regions with very stringent physical conditions, it is not possible to know whether the ring is tracing part of a toroidal or a disc-like structure. The kinematics of the ring in K3-35 suggests the presence of both rotating and expanding motions, as was also proposed in the young PN IRAS 17347-3139 \\citep{deg04}. The estimated expansion and rotation velocity values for K3-35 are similar (a few km~s$^{-1}$) to those obtained from water maser observations of IRAS 17347-3139.  The calculated expansion velocity of the ring ($\\simeq$1.4~km~s$^{-1}$) in K3-35 is comparable to, but below, the expansion velocities of the tori inferred from interferometric CO observations in some PPNe and young PNe. For instance, the expansion velocities in M~1-92, M 2-9, M 2-56, and KjPn 8, have values in the range of $\\simeq$5--8~km~s$^{-1}$ \\citep{buj98,zwe97,cas02,for98}. It is noteworthy that the estimated expansion velocity in K3-35 obtained from water maser observations is quite similar to the slow expansion velocity ($\\simeq$0.8~km~s$^{-1}$) deduced from CO observations in the PPN Red Rectangle \\citep{buj05}. However, the estimated sizes of the tori detected in CO observations in those sources are about one order of magnitude more than the size of structure traced by the water masers in K3-35, suggesting that the structures observed in CO are probably related to the outermost equatorial region, while the water masers could be arising from the innermost one, such as a circumstellar disc, closer to the central star.  \\begin{table} \\caption{Results of the rotating and expanding ring model fit to the H$_2$O masers towards the centre of K3-35.} \\label{symbols} \\begin{tabular}{@{}ccrcl} \\hline Parameter & \\multicolumn{4}{c}{Value$^a$} \\\\ \\hline $a$             &  & 0\\farcs021   & $\\pm$ & 0\\farcs003    \\\\ $b$             &  & 0\\farcs012   & $\\pm$ & 0\\farcs002    \\\\ $x_0\\,^*$       &  & 0\\farcs001   & $\\pm$ & 0\\farcs001    \\\\ $y_0\\,^*$       &  & 0\\farcs004   & $\\pm$ & 0\\farcs004    \\\\ PA              &  & 158$\\degr$     & $\\pm$ & 10$\\degr$ \\\\ $i$             &  & 55$\\degr$      & $\\pm$ & 7$\\degr$ \\\\ $v_s$           &  & 22.8    & $\\pm$ & 0.5~km~s$^{-1}$ \\\\ $v_{\\rmn{exp}}$ &  & 1.4     & $\\pm$ & 0.9~km~s$^{-1}$ \\\\ $v_{\\rmn{rot}}$ &  & 3.1     & $\\pm$ & 0.8~km~s$^{-1}$ \\\\ \\hline \\end{tabular} \\medskip  $^a$ Note: Uncertainties are 2$\\sigma$.\\\\ $^*$ Coordinates of the centre of the ellipse relative to the position of the 1.3~cm continuum emission peak. \\end{table}  In the case of K3-35, the water masers may be delineating an ellipse on the plane of the sky with a semimajor axis of $\\simeq$100 AU and a PA$\\simeq$158$\\pm$10\\degr, suggesting the presence of a disc or torus projected on the sky. The major axis of the ellipse is almost perpendicular (within the uncertainties) to the direction of the bipolar emission traced by the innermost region of the jet, with a PA$\\simeq$65\\degr observed by \\citet{mir01}, suggesting that the disc or torus could be physically related to the collimation of the bipolar outflow.  Recently, \\citet{hug07} found that jets typically appear a few hundred years after the torus formation in PPNe and young PNe. This time sequence provides evidence that jets and tori are physically related. In the case of K3-35, \\citet{vel07} modelled its radio continuum emission as a precessing dense jet with an age of $\\sim$40~yrs. On the other hand, our estimate of the dynamical age of the ring traced by the maser emission is about $\\sim$350 years, using its radius ($\\sim$100~AU) and assuming a uniform expansion velocity of $\\sim$1.4~km~s$^{-1}$. These values would, in principle, suggest a lag in time similar to the value found by \\citet{hug07}. However, the dynamical age we derive is not reliable enough as an age estimate. Proper motion studies of the water maser emission (e.g. using e-MERLIN and VLBA) could provide better estimates.  The kinematics of the ring in K3-35 suggests the presence of both expansion and rotation. A rough estimate for the central stellar mass can be obtained by assuming that the total energy (kinetic plus gravitational) of the masing gas is close to zero. In this case, $M\\simeq (R/2G)(v_{\\rmn{exp}}^2+v_{\\rmn{rot}}^2)$, where $R\\simeq100~(D/5~\\rmn{kpc})$~AU, is the radius of the ring. Hence, $M\\simeq[0.7~(D/5~\\rmn{kpc})\\pm 0.3]~\\rmn{M}_{\\sun}$, where $D$ is the source distance. The error in the mass only includes the errors in the fitted parameters. This estimate of the central mass is in agreement with the core mass required for a PNe, according to evolutionary models \\citep{kwo03}."
    },
    "0808/0808.0896_arXiv.txt": {
        "abstract": "We present an observational study about the  effects of the interactions in the kinematics, stellar population and abundances of the components of the galaxy pair AM\\,2306-721. Rotation curves for the main and companion galaxies were obtained, showing a deprojected velocity amplitude of  175 km s$^{-1}$ and 185 km s$^{-1}$, respectively. The interaction between the main and companion galaxies was modeled using numerical N-body/hydrodynamical simulations, with the result indicating that the current stage of the merger would be about 250 Myr after perigalacticum. The spatial variation in the distribution of the stellar population components in both galaxies was analysed  by fitting combinations of stellar population models of different age groups. The central region of main galaxy is  dominated by an old (5-10\\,Gyr) population, while significant contributions from a young (200\\,Myr)  and intermediate (1\\,Gyr) components are found in the disk, being enhanced in the direction of the tidal features. The stellar population of  the companion galaxy is overall much younger, being dominated by components with 1\\,Gyr or less, quite widely spread over the whole disk. Spatial profiles of the oxygen abundance were obtained from the a grid of photoionization models using the $R_{23}$ line  ratio. The disk of the main galaxy shows a clear radial gradient, while the companion galaxy presents an oxygen abundance relatively homogeneous across the disk. The absence of an abundance gradient in the secondary galaxy is interpreted in terms of mixing by gas flows from the outer parts to the center of the galaxy due to the gravitational interaction with the more massive primary. ",
        "introduction": "It is widely accepted by some time that merging and interaction events play an important role in the formation and evolution of galaxies. Mergers change the mass function of galaxies, creating a progression from small galaxies to larger ones; the merging process can also change the morphology of the constituents, transforming gas-rich spirals in quiescent ellipticals. Interactions can also trigger a wide  set of physical and morphological phenomena, such as tidal tails, bridges and shells, kinematically decoupled cores and star formation enhancements (see review of \\citealt{struck99}).  Interacting galaxies also show enhanced star formation when compared with isolated galaxies. Such enhancement was initially proposed  by \\citet{larson78} to explain the wider range of optical colors found in galaxies in pairs. Since then, numerous studies have confirmed these results, especially in the central regions,  through measurements of optical emission lines  \\citep{kennicutt84, kennicutt87, donzelli97, barton03, woods07}, infrared emission \\citep{joseph85, sekiguchi92, geller06} and radio  emission \\citep{hummel81}. Recent studies have also shown that this enhancement is a function of the projected galaxy pair separation (e.g. \\citealt{barton00, lambas03, nikolic04}), being stronger in low-mass  than in high-mass galaxies (e.g. \\citealt{woods07,ellison08}). In particular, \\citet{ellison08} found that the star formation rate as  measured by the H$\\alpha$ equivalent width for galaxies in  pairs selected from the Sloan Digital Sky Survey, is some  70\\% higher when compared to a control sample of galaxies with equal stellar mass distribution.  There is a connection between the interaction strength and the morphological distortion in binary galaxies. According to \\citet{mihos96} models, the response of the gas to a close passage depends dramatically on the mass distribution of the galaxy, with the irregularities in the gas velocity field tracing the disturbances in the gravitational potential of the galaxy as observed, for example, in some galaxies in the Virgo cluster \\citep{rubin99}. Combined N-body/hydrodynamic simulations show that galaxy-galaxy mergers disturb the gas velocity field significantly, and hence lead to asymmetries and distortions in the rotation curves \\citep{kronberger06}. However, according to these authors, no severe distortions are observable about 1 Gyr after the first encounter.  The gas motions created by the interaction can also significantly alter the chemical state of the galaxies \\citep{koeppen90,dalcanton07}, and modify the usually smooth radial metallicity gradient often found in isolated disk galaxies \\citep{henry99}. Recently, \\citet{kewley06} found that O/H abundance in the central region of nearby galaxy pairs is systematically lower that that of isolated objects. These authors suggest that the lower metallicity is a consequence of gas infall caused by the interaction, but very few observational studies have been published analysing in detail the influence of different levels of  interactions in the metallicity distribution and enrichment properties of galaxies.  Ferreiro \\& Pastoriza 2004 (hereafter FP04), in a study of the integrated photometry and star formation activity in a sample of interacting systems from the Arp-Madore catalogue \\citep{arp87}, found that the galaxies involved have bluer colours than those of isolated galaxies of the same morphological type, indicating an enhancement of star forming activity.  This enhancement was also  previously suggested by \\citet{donzelli97}  to explain the slightly larger  values of H$\\alpha\\, +\\, $\\ion{N}{ii} equivalent widths found in these systems, when compared with normal isolated spiral galaxies. From their sample, we have selected several systems to start a more comprehensive study of the effects of the interactions in the kinematics, stellar population and abundances of the galaxies in the so-called ``minor merger'' systems, defined here as physical pairs with mass ratio in the range of $0.04 < M_{secondary}/M_{primary} < 0.2 $.  This paper presents the results for the system AM\\,2306-721. This pair is composed by a peculiar spiral with disturbed arms  (hereafter, AM\\,2306A) in interaction with an irregular galaxy (hereafter, AM\\,2306B). Both galaxies contain very luminous  \\ion{H}{ii} regions with H$\\alpha$ luminosity in the range of $8.30\\times\\,10^{39}\\,<\\,$L(H$\\alpha\\,)\\, <\\,1.27\\times\\,10^{42}$erg\\, s$^{-1}$ as estimated from H$\\alpha\\,$ images \\citep{ferreiro08}; and high star formation rate in the range of 0.07 to 10  $ M_{\\sun}$/yr \\citep{ferreiro08}. The present paper is organized as follows: in Section \\ref{datared},  we summarize the observations and data reduction. Section \\ref{vel} describes the gas kinematics of each galaxy and Section \\ref{numsim} present the numerical N-body/hydrodynamical simulations of the interaction. Section \\ref{sintese} presents the stellar population synthesis. The metallicity analysis is  presented in Section \\ref{emission}, and the conclusions are summarized in Section \\ref{final}. ",
        "conclusions": "\\label{final} An observational study about the  effects of the interactions in the kinematics, stellar population and abundances of the galaxy pair AM\\,2306-721 is performed. The data consist of long-slit spectra in the  wavelength range of 3\\,350 to 7\\,130\\AA\\, obtained with the  Gemini Multi-Object Spectrograph at Gemini South. The main findings are the following: \\begin{enumerate} \\item Rotation curves of the main and companion galaxies with an deprojected velocity amplitude of  175 km s$^{-1}$ and 185 km s$^{-1}$ respectively were obtained. An estimate of the dynamical mass was derived for each galaxy, using the deprojected velocity amplitude. For the main galaxy, its dynamical mass is $ 1.29 \\times 10^{11} M_{\\sun}$ within  a radius of 18 kpc; and  for the companion galaxy, the estimated dynamical mass  is $ M(R)= 8.56 \\times 10^{10} M_{\\sun}$ within a radius of  10.7 kpc. \\item In the main galaxy, radial velocity deviations from the disk rotation of about 100 km\\,s$^{-1}$ were detected, which are  probably due to the interaction with the companion galaxy. \\item In order to reconstruct the history of the AM\\,2306-721 system and to predict the evolution of the encounter, we modeled the interaction between AM\\,2306A and AM\\,2306B through numerical N-body/hydrodynamical simulations. The orbit that best reproduces the observational properties is found to be hyperbolic, with an eccentricity $e=1.15$ and perigalacticum of $q=5.25$ kpc; the current stage of the system would be about 250 Myr after perigalacticum. \\item The spatial variations  of the stellar population components of the galaxies were analysed  by fitting combinations of stellar population models of different ages (2.5 Myr, 200 Myr, 1 Gyr, 5 Gyr and 10 Gyr)  and solar metallicity. The central regions of the main galaxy  are dominated by the old (5-10\\,Gyr) population, with some significant contribution from a young 200\\,Myr  and intermediate 1\\,Gyr component along the disk of the galaxy. On the other hand, the stellar population in the companion galaxy is overall much younger, being dominated by the 2.5\\,Myr, 200\\,Myr and 1\\,Gyr components, which are quite widely spread over the whole disk. \\item Oxygen abundance spatial profiles were obtained using a grid of photoionization models and the $R_{23}$=([\\ion{O}{ii}]$\\lambda\\,3727+$[\\ion{O}{iii}] $\\lambda\\,4959+$[\\ion{O}{iii}]  $\\lambda\\,5007)/$H$\\beta$ line  ratio. The disk of the main galaxy shows a clear radial oxygen abundance gradient,  that can be fitted as a linear function $12+\\ohlog =8.75(\\pm0.06)-0.025(\\pm 0.007)R_{G}$; in the companion galaxy the oxygen abundance is relatively homogeneous across the galaxy disk, with the observed  mean value of $12+\\ohlog =8.39$ \\item The absence of an abundance gradient in the secondary galaxy is interpreted as it having been  destroyed  by gas flows from the outer parts to the center of the galaxy. \\end{enumerate}  \\thanks"
    },
    "0808/0808.2294_arXiv.txt": {
        "abstract": "We show that repeated sound waves in the intracluster medium (ICM) can be excited by a single inflation episode of an opposite bubble pair. To reproduce this behavior in numerical simulations the bubbles should be inflated by jets, rather than being injected artificially. The multiple sound waves are excited by the motion of the bubble-ICM boundary that is caused by vortices inside the inflated bubbles and the backflow (`cocoon') of the ICM around the bubble. These sound waves form a structure that can account for the ripples observed in the Perseus cooling flow cluster. We inflate the bubbles using slow massive jets, with either a wide opening angle or that are precessing. The jets are slow in the sense that they are highly sub-relativistic, $v_j \\sim 0.01c-0.1c$, and they are massive in the sense that the pair of bubbles carry back to the ICM a large fraction of the cooling mass, i.e., $\\sim 1-50 M_\\odot \\yr^{-1}$. We use a two-dimensional axisymmetric (referred to as 2.5D) hydrodynamical numerical code (VH-1). ",
        "introduction": "\\label{sec:intro}  Higher X-ray emissivity arcs, termed ripples, are observed in the intracluster medium (ICM) of the Perseus cluster (Fabian et al. 2003, 2006; Sanders \\& Fabian 2007), as well as in A2052 (Blanton et al. 2007). The ripples are thought to be sound waves that are excited by the inflation of bubbles near the cluster's center. Some of the ripples are observed on the outer edge of bubbles; they are compressed ICM that will turn into shock waves, and later to sound waves. It is possible that some of the other ripples are compressed by weak jets (Soker \\& Pizzolato 2005), rather than being sound waves detached from their parent bubble. Beside being interesting by their own merit, the sound waves can heat the ICM and partially offset the radiative cooling of the ICM (e.g., Br\\\"uggen et al. 2005; Fabian et al. 2005; Fujita \\& Suzuki 2005; Graham et al. 2008). \\par  To obtain repeated sound waves previous numerical simulations assumed repeating bubble inflation episodes (e.g., Ruszkowski et al. 2004a,b; Sijacki \\& Springel 2006a, b). This had to be done because these simulations used \\emph{artificial bubbles}, i.e., numerically injected spherical bubbles at off-center locations, as is commonly done (e.g., Br\\\"uggen 2003; Br\\\"uggen \\& Kaiser 2001; Gardini 2007; Jones \\& De Young 2005; Pavlovski et al. 2008; Reynolds et al. 2005; Robinson et al. 2004; Ruszkowski et al. 2004a, 2004b, 2007, 2008; Scannapieco \\& Br\\\"uggen 2008). This of course is not the way bubbles are formed in clusters. Bubble are formed by jets [a different approach of injecting magnetic energy instead of a jet (Nakamura et al. 2008; Xu et al. 2008) requires further study]. When large, almost spherical bubbles (termed \\emph{fat bubbles}) are inflated by jets, some phenomena appear, that are not revealed when artificial bubbles are used (Sternberg \\& Soker 2008b). Among these are: (1) more stable bubbles, and (2) bubble-ICM boundary that is stochastically vibrates due to the vortices inside the bubble and the back flowing ICM.  In previous works we found that in order for the jets to inflate fat bubbles they should have a large opening angle, or they should rapidly precess (Soker 2004, 2006; Sternberg et al. 2007; Sternberg \\& Soker 2008a). In addition, jets that are relatively slow, $v_j \\sim 10^4 \\km \\s^{-1} \\ll c$,  and with a mass loss rate of the order of $2\\dot M_j \\sim1-50 M_\\odot \\yr^{-1}$ (for both jets), are more likely to inflate fat bubbles. The same results can be seen in the simulations of Omma et al. (2004), Alouani Bibi et al. (2007) and Binney et al. (2007), who use similar parameters to ours, but instead of conical jets with wide opening angles they use cylindrical jets with large initial radius. This is also shown by the simulation of Heinz et al. (2006), who manages to inflate a bubble (although not a bubble attached to the center) by launching a jet with a velocity of $3\\times 10^4 \\km \\s^{-1}$ and a mass loss rate of $35 M_\\odot \\yr^{-1}$ in one jet. In this paper we continue our study of jet-inflated bubbles, and show that such bubbles can excite several consecutive sound waves without the need to invoke periodic jet launching episodes. We do not deal with the propagation of sound waves and their properties. These seem to require more sophisticated treatment (Graham et al. 2008). We limit ourself to show that a proper treatment of bubble inflation can better explain the excitation of sound waves in the ICM. \\par ",
        "conclusions": "\\label{sec:discussion}  We followed the response of the ICM to the inflation of bubbles by jets. The inflation of a bubble by a jet results in vortices inside the bubble and a backflow of the ICM around the bubble (Sternberg et al. 2007; Sternberg \\& Soker 2008a, b). Both processes cause the bubble-ICM to change shape with time and to become corrugated. In the case of a precessing jet (section \\ref{sec:precessing}), the changes in the jet axis cause a more prominent change in the bubble-ICM shape. This motion of the bubble-ICM boundary sends shocks and sound waves into the ICM. This effect cannot be reproduced by artificial bubbles, i.e., by numerically inserting spherical bubbles at off-center locations. \\par  Our main result is that one episode of a bubble pair inflation can excite several sound waves along each radial direction. These form high-density arcs, the ripples. The front ripple is actually a weak shock. There is no need to introduce periodic (or semi-periodic) episodes of bubble inflation. Nonetheless, multiple episodes will make more ripples, and make the ripple structure more complicated. This is the case in Perseus, where the main central bubble pair was formed by jets that changed their direction and intensity (Dunn et al. 2006).  We also observe a dense arc at the bubble's front that is apparent part of the time. The dense arc is the excitation of a new sound wave at the compression phase. Part of the time the region at the bubble's front is at the rarefaction phase, the wave trough, and no dense arc is observed there. \\par  Our results can be incorporated into a broader scope. In a previous paper (Sternberg \\& Soker 2008b) we found that to follow the evolution of bubbles in the ICM one must inflate them by jets, rather than introduce them artificially. The evolution of bubbles and their influence on the ICM, e.g., sound waves, is crucial for the deposition of heat from the central active galactic nuclei to the ICM. Our works show that in order to understand the heating of the ICM one must properly inflate bubbles. \\par"
    },
    "0808/0808.0772_arXiv.txt": {
        "abstract": "The effect of neutrino trapping on the longitudinal dielectric function at low densities has been investigated by using different relativistic mean field models. Parameter sets G2 of Furnstahl-Serot-Tang and Z271 of Horowitz-Piekarewicz, along with the adjusted parameter sets of both models, have been used in this study. The role of the isovector adjustment and the effect of the Coulomb interaction have been also studied. The effect of the isovector adjustment is found to be more significant in the Horowitz-Piekarewicz model, not only in the neutrinoless matter, but also in the matter with neutrino trapping. Although almost independent to the variation of the leptonic fraction, the instability region of matter with neutrino trapping is found to be larger. The presence of more protons and electrons compared to the neutrinoless case is the reason behind this finding. For parameter sets with soft equation of states at low density, the appearance of a large and negative $\\varepsilon_L (q,q_0=0)$ in some parts of the edge of the instability region in matter with neutrino trapping is understood as a consequence of the fact that the Coulomb interaction produced by electrons and protons interaction is larger than the repulsive isovector interaction created by the asymmetry between the proton and neutron numbers. ",
        "introduction": "\\label{sec_intro} At low densities both the relativistic and the non-relativistic mean field models predict a liquid-gas phase transition region for nuclear matter leading, for dense star matter, to a non-homogeneous phase commonly called pasta phase, which is formed by a competition between the long range Coulomb repulsion and the short range nuclear attraction~\\cite{Provi07}. This transition has substantial consequences on the properties of stellar matter and neutrino transport~\\cite{Duco07}. Considerable efforts to comprehend the uniform ground state stability of multi-component systems consisting of electrons, neutrinos, protons, and neutrons as a good approximation of this transition have been recently devoted, not only in the zero temperature approximation, but also for finite temperature~\\cite{Pethick,Douchin,Carr,Horowitz01,Provi1,Provi2,Provi3,Provi4,anto06,Muller:1995ji}. It is obvious that in order to understand the physics inside the non-homogeneous (unstable) regions like the mechanism of nuclear creation with slab-like or rod-like shape, we have to go beyond the mean field approximation. Attempts in this direction are discussed in Refs.~\\cite{Provi07,Napoli07,Watanabe04,Horowitz04,Sonoda}. Moreover, in the collapsing supernova core and at sub-nuclear densities, the transition of nuclear shape from sphere to other exotic shapes has significant effects to the neutrino mean free path. However, how these effects modify the neutrino mean free path is not fully understood yet~\\cite{Horowitz04,Sonoda}.  Another motivation to study this transition comes from the fact that a neutron star is expected to have a solid inner crust of nonuniform neutron-rich matter above its liquid mantle~\\cite{Carr} and the mass of its crust depends sensitively on the density of its inner edge and on its equation of state (EOS)~\\cite{Douchin}. On the other hand, the critical density ($\\rho_c$), a density at which the uniform liquid becomes unstable to a small density fluctuation, can be used as a good approximation of the edge density of the crust~\\cite{Carr}. By generalizing the dynamical stability analysis of Ref.~\\cite{Horo1} in order to accommodate the various nonlinear terms in the relativistic mean field (RMF) model of Horowitz-Piekarewicz~\\cite{Horowitz01}, Carriere {\\it et al.}~\\cite{Carr} found a strong correlation between $\\rho_c$ in the neutron star and the density dependence of nuclear matter symmetry energy ($a_{\\rm sym}$). This leads to a suggestion that a measurement of the neutron radius in $^{208}{\\rm Pb}$ will provide useful information on the $\\rho_c$ \\cite{Carr,Horowitz01}.  In our previous work~\\cite{anto06}, the critical densities of uniform matter with and without neutrino trapping have been calculated and analyzed by means of different RMF models. In this analysis it is shown that the interplay between the dominant contribution of the matter composition and the effective masses of mesons and nucleons leads to higher critical densities for matter with neutrino trapping. Furthermore, it was also found that the predicted critical density is insensitive to the number of trapped neutrinos as well as to the RMF model used. However, the discussion about the reason behind these findings was not quite robust. On the other hand, as we mentioned above, the neutrino transport is very crucial in the dynamics of the core-collapsing supernova due to the fact that the neutrinos carry most of the energy away and will lose also their energies by exciting collective nuclear and plasmon modes~\\cite{Provi2}. Similar situation can be also found in the neutron stars. Moreover, it was also shown in Ref.~\\cite{Provi2} that the behavior of electrons in matter depends strongly on the wave length or momentum of the external perturbation $q$. Note that this momentum is related to the energy transfer of the neutrinos that propagate in matter.  The present paper reports on the extension of our previous investigation~\\cite{anto06} by calculating the longitudinal dielectric function of ERMF models and analyzing the relation between the obtained results and the isovector sector adjustment, the presence of the long-range Coulomb interaction, as well as the presence of electrons. The purpose of this work is to explain the reason behind the appearance of each point along the onset of the instability. To this end, we should emphasize here that we need to calculate the dielectric function in the edge of non-homogeneous regions because information from the critical density alone is insufficient. Furthermore, it is also important to emphasize that our definition of the instability is not the non-homogeneous area, but rather it is connected with the points where these non-homogeneities start to appear. As a consequence, the assumption of the uniform matter in the calculation is still valid.  This paper is organized as follows. The RMF models and some constraints used in the present analysis are briefly discussed in Sec.~\\ref{sec_models}.  In Sec.~\\ref{sec_LDF}, a discussion of the longitudinal dielectric function is given. In Sec.~\\ref{sec_results} we present the graphical results of the onset of instability along with the corresponding discussions. Finally, we give the conclusion in Sec.~\\ref{sec_conclu}. ",
        "conclusions": "\\label{sec_conclu} We have studied how the instability region starts to appear in low-density matter described by the Horowitz-Piekarewicz and Furnstahl-Serot-Tang models. To this end we have utilized the longitudinal dielectric function at $q_0=0$. The importance of the electron and Coulomb terms in matter with neutrino trapping has been investigated. It is found that the adjustment of the isovector terms has a more significant effect in the Horowitz-Piekarewicz model, i.e., producing a stronger repulsive isovector contribution which leads to a stronger correlation between its low density instability region and the $a_{\\rm sym}$ compared to the model of Furnstahl-Serot-Tang for matter without neutrino trapping. In the case of matter with neutrino trapping, the parameter sets with stiff EOS at low density lead to a large and positive $\\varepsilon_L (q,q_0=0)$. This demonstrates that, although the onsets of the instability of parameter sets with stiff and soft EOS at low densities are similar, the driving mechanisms are different. This fact might have an effect to the neutrino transport in matter. In both models the effect of the variation of the leptonic fraction is negligible, but the effect of the neutrino trapping on the onset of the instability region is significant. The presence of more protons and electrons in matter with neutrino trapping is the reason behind this phenomenon. The Coulomb term is found to be decisive in enlarging the stability of matter in this density region. The presence of the large and negative $\\varepsilon_L (q,q_0=0)$ in some parts of the instability region of matter with neutrino trapping originates from the fact that the isovector term is insufficient to cancel the attractive Coulomb interaction contributions generated by the presence of electrons (and protons)."
    },
    "0808/0808.0919_arXiv.txt": {
        "abstract": "We performed cosmological, magneto-hydrodynamical simulations to follow the evolution of magnetic fields in galaxy clusters, exploring the possibility that the origin of the magnetic seed fields are galactic outflows during the star-burst phase of galactic evolution. To do this we coupled a semi-analytical model for magnetized galactic winds as suggested by \\citet{2006MNRAS.370..319B} to our cosmological simulation. We find that the strength and structure of magnetic fields observed in galaxy clusters are well reproduced for a wide range of model parameters for the magnetized, galactic winds and do only weakly depend on the exact magnetic structure within the assumed galactic outflows. Although the evolution of a primordial magnetic seed field shows no significant differences to that of galaxy clusters fields from previous studies, we find that the magnetic field pollution in the diffuse medium within filaments is below the level predicted by scenarios with pure primordial magnetic seed field. We therefore conclude that magnetized galactic outflows and their subsequent evolution within the intra-cluster medium can fully account for the observed magnetic fields in galaxy clusters. Our findings also suggest that measuring cosmological magnetic fields in low-density environments such as filaments is much more useful than observing cluster magnetic fields to infer their possible origin. ",
        "introduction": "Magnetic fields have been detected in galaxy clusters by radio observations, via the Faraday rotation signal of the magnetized cluster atmosphere towards polarized radio sources in or behind clusters \\citep{2002ARA&A..40..319C} and from diffuse synchrotron emission of the cluster atmosphere \\citep[see][for recent reviews]{2004IJMPD..13.1549G,2008SSRv..134...93F}. However, our understanding of their origin is still very limited.  At present, models for the origin of seed fields can be classified in three main groups. In the first, a magnetic field is created in shocks through the ''Biermann battery'' effect \\citep{1997ApJ...480..481K,Ryu..1998,2001ApJ...562..233M}. A subsequent turbulent dynamo boosts it to the field strength observed in galaxy clusters. A second class of models invokes processes that took place in the early universe. In general, they predict that magnetic seed fields fill the entire volume of the universe; however the coherence length of the field crucially depends on the details of the models \\citep[see][for a review]{Grasso..PhysRep.2000}. Finally, galactic winds \\citep[e.g.][]{Volk&Atoyan..ApJ.2000} or AGN ejecta \\citep[e.g.][ and references therein]{1997ApJ...477..560E,Furlanetto&Loeb..ApJ2001} can produce magnetic fields and pollute the proto-cluster region. In such models, the magnetic field can also originate from an early population of dwarf, starburst galaxies \\citep{1999ApJ...511...56K} at relatively high redshift ($z \\approx 4 - 6$).  In previous work, non radiative simulations of galaxy clusters within a cosmological environment which follow the evolution of a primordial magnetic seed field were performed using Smooth-Particle-Hydrodynamics (SPH) codes \\citep{1999A&A...348..351D,2002A&A...387..383D,2005JCAP...01..009D} as well as Adaptive Mesh Refinement (AMR) codes \\citep{2005ApJ...631L..21B,2008A&A...482L..13D,2008ApJS..174....1L}. Although these simulations are based on different numerical techniques they show good agreement in the predicted properties of the magnetic fields in galaxy clusters, when the evolution of an initial magnetic seed field is followed. This work has also  demonstrated, that the properties of the final magnetic field in galaxy clusters do not depend on the detailed structure of the assumed initial magnetic field. The spatial distribution and the structure of the predicted magnetic field in galaxy clusters is primarily determined by the dynamics of the velocity field imprinted by cluster formation \\citep{1999A&A...348..351D,2002A&A...387..383D} and compares well with measurements of Faraday rotation.  The creation of magnetic fields in shocks through the ''Biermann battery effect'' \\citep{1997ApJ...480..481K,Ryu..1998}, and subsequent turbulent dynamo action can be followed as well as a prediction can be made for magnetic field values from velocity fields inferred in cosmological simulations \\citep{2008Sci...320..909R}. Both methods predict magnetic field strengths in filaments with somewhat higher values \\citep[e.g. see ][]{2004PhRvD..70d3007S} than found in simulations that follow the evolution of a primordial magnetic seed field.  Faraday rotation can be observed in several radio galaxies located at different radial distances with respect to the cluster center. Motivated by numerical simulations \\citep{2001A&A...378..777D}, the observed magnetic field is often modelled with a radially-declining field strength and a power law spectral structure. >From such observations, once can constrain the power law spectral index \\citep{2004A&A...424..429M,2006A&A...460..425G} or directly reconstruct the power spectrum of the magnetic field \\citep{2003A&A...412..373V,2005A&A...434...67V}. Given the sparse observational data available at the moment, a degeneracy exists between the central value of the magnetic field and its rate of radial decline \\citep[see for example][]{bonafede08,2008A&A...483..699G}, for which detailed predictions from simulations can be useful in breaking the degeneracy. The simulations must therefore examine different possible magnetic field origins in galaxy clusters in order to test the robustness of the inferred magnetic field properties.  Recently the validity of models that produce a cluster magnetic field from galactic winds has been supported by a semi-analytic modelling of galactic winds \\citep{2006MNRAS.370..319B}. However, these models are unable to predict how the magnetic fields produced by the ejecta of the galaxies are compressed and amplified by the process of structure formation. Therefore, the structure of the final magnetic field in galaxy clusters cannot yet be predicted by these models. In this work we extend these studies by directly incorporating the galactic outflow model in magneto-hydrodynamical simulations of structure formation.  The paper is structured as follows: In section \\ref{sim setup} we present the details of the cosmological setup, concentrating especially on the coupling of the semi-analytic model to the cosmological simulations. Details of the wind model and  the seed magnetic field are presented in section \\ref{wind_model} and appendix \\ref{E2B}. The general results of our simulations are presented in section \\ref{results_all} and in section \\ref{results} we discuss our findings, particulary how galaxy clusters formed in our simulations. Finally, we present our conclusions in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In this work we performed cosmological, magnetohydrodynamic simulations following the evolution of magnetic fields on large scale structures and in galaxy clusters. Coupling a semi-analytic model for magnetized galactic winds as suggested by \\citet{2006MNRAS.370..319B} with our cosmological simulation we explored the possibility that the magnetic fields in galaxy clusters originate from galactic outflows during star-burst phases, further processed by structure formation. We compared our results with the ones obtained by following a primordial magnetic seed field \\citep{2005JCAP...01..009D}. Performing several simulations, we explored the effect of various parameters of the adapted semi-analytic model relevant for the strength of the magnetic seed field from the galactic outflows.  We also explored the effect of the magnetic field configuration assumed for the galactic outflows and of the seeding strategy. Our general findings are: \\begin{itemize} \\item The typical magnetic field strengths of several $\\mu$G in galaxy clusters as obtained from observations of Faraday rotation are well reproduced for a wide range of parameters of the galactic outflow model. \\item The general shape of the predicted Faraday rotation profile within clusters compares well with the sparse observational data available. Models that assume a field strength of $5\\mu$G within the galactic halo reproduce the observed Faraday rotation profiles better than models with a ten times stronger or a ten times weaker halo magnetic field. \\item The properties of the final magnetic field in galaxy clusters do not depend on the exact field configuration within the magnetic outflows.  This confirms previous studies that the structure of the magnetic field in galaxy clusters is primarily driven by the velocity field induced by the structure formation process. \\item In massive galaxy clusters, the magnetic field amplification saturates around values of several $\\mu$G. The mass (or temperature) scale on which this happens depends on the strength of the magnetic seed field, and probably also on the resolution of the simulation. \\item In systems where saturation effects start to play a significant role, we observe only a small scatter in the magnetic scaling relations, and in the shape of the radial, magnetic profiles. \\item In clusters where saturation effects are negligible the strength and configuration of the magnetic field strongly depends on the dynamic state. Therefore, we observe a large scatter in the magnetic scaling relations and in the shape of the radial, magnetic profiles. \\item Within galaxy clusters, the structures predicted from synthetic Faraday rotation maps do not depend significantly on the magnetic seed field, and agree well with observed ones. \\item In low density environments imprints of the magnetic seed fields are still present and can be observed, in principle. In these environments the galactic outflows forming at later epochs contribute significantly to the magnetic field configuration. \\end{itemize}  In summary, we confirm that galactic outflows and the subsequent action of structure formation can explain the properties of the magnetic fields observed in massive galaxy clusters. Their strength and shape is attributable to the velocity field induced by structure formation. We find that there are no measurable imprints of the magnetic seed fields left in the synthetic Faraday rotation maps of our simulated clusters. However, low density environments, like filaments, still contain significant information on the magnetic seed fields. Therefore, they may be used to discriminate proposed scenarios once information on magnetic fields within these regions becomes available with the next generation of radio instruments."
    },
    "0808/0808.2139_arXiv.txt": {
        "abstract": "In the microquasar V4641 Sgr the spin of the black hole is thought to be misaligned with the binary orbital axis. The accretion disc aligns with the black hole spin by the Lense-Thirring effect near to the black hole and further out becomes aligned with the binary orbital axis.  The inclination of the radio jets and the Fe\\ka\\ line profile have both been used to determine the inclination of the inner accretion disc but the measurements are inconsistent.  Using a steady state analytical warped disc model for V4641 Sgr we find that the inner disc region is flat and aligned with the black hole up to about $900\\, R_{\\rm g}$. Thus if both the radio jet and fluorescent emission originates in the same inner region then the measurements of the inner disc inclination should be the same. ",
        "introduction": "Microquasars are binary systems in which material is accreted from a normal star on to a compact object. They differ from typical X-ray binaries by the strong presence of a persistent or episodic radio jet \\citep{MR99}. The compact object is usually associated with a neutron star or a stellar mass black hole. To date there are over fifteen microquasars for which the compact object has been dynamically confirmed to be a stellar mass black hole \\citep{O03}. However, only four of these systems have well resolved relativistic radio jets \\citep[XTE{\\thinspace J1550--564}, GRO{\\thinspace J1655--40}, GRS{\\thinspace 1915+105} and V4641 Sgr,][]{G03}.  It is usually assumed that the inclination of the jet axis is perpendicular to that of the orbital plane of the binary. However, it has been shown by \\cite{Macc} and more recently by \\cite{MTP08} that the alignment time-scale in microquasars is usually a significant fraction of the lifetime of the system. If the black hole in such microquasars were formed with misaligned angular momentum, as expected from supernova-induced kicks, then it would be likely that the system would remain misaligned for most of its lifetime.  Precise measurements of both the orbital plane and jet inclination are known for two systems. GRO{\\thinspace \\jb} is a microquasar thought to contain a rapidly rotating black hole \\citep{Z97,R08a} with a mass constrained to be larger than $6.0\\msun$ \\citep{OB97}. \\cite{H95} measured a jet inclination of $85$\\deg$\\pm{\\thinspace 2}$\\deg\\ to the line-of-sight and it is thus misaligned by at least $10$\\deg\\ to the orbital plane, with an inclination of $70.2$\\deg$\\pm{\\thinspace 1.9}$\\deg\\ \\citep{G01}. \\cite{MTP08} presented a detailed investigation of the alignment timescale of this system and found that it is consistent with the lifetime of the secondary star.   V4641{\\thinspace Sgr} was discovered as an X-ray source independently with the Wide Field Cameras on \\bepo\\ on 1999 February 20 \\citep{intZ99} and with the Proportional Counter Array on the {\\it Rossi X-ray Timing Explorer} (\\rxte) on 1999 February 18 \\citep{MSM99}. Spectroscopic observations made between 1999 September 17 and 1999 October 16 led to a mass function $f(M) = 2.74\\pm0.12$\\msun\\ \\citep{O01}. From the lack of X-ray eclipses, combined with the large amplitude of the folded light curve, they deduced an orbital inclination angle in the range $60$\\deg$\\le i_{\\rm orbit} \\le 70$\\deg$.7$ and mass $8.73 \\le M_{\\rm BH} \\le 11.70$\\msun\\ for the compact object. Radio observations of \\vs\\ \\citep{H00} made during the 1999 September outburst found the jet expanding at apparent superluminal velocities with a proper motion ranging between $0.22\\arcsec-1\\arcsec$ per day. Based on the radio information, \\cite{O01} suggested that that the jet must be highly beamed and have an inclination along the line-of-sight of $i_{\\rm jet} \\le 10$\\deg.  This differs significantly from from the inclination of the binary orbital axis.  X-ray spectral analysis made on \\bepo\\ observations of \\vs\\ \\citep{intZ00} revealed the presence of a strong Fe\\ka\\ emission with an equivalent width between 0.3 and 1\\kev. They interpreted this as fluorescent emission using a photo-ionized medium. The presence of this Fe fluorescent emission was later confirmed by \\cite{R02} from a collection of \\rxte\\ data.  They found an emission line at about $6.6$\\kev\\ with an equivalent width of about $360$\\ev.  A more recent analysis of the source with \\rxte\\ data obtained during the outburst of 2003 August 5--17 \\citep{MB06} showed the presence of both a strong Fe\\ka\\ fluorescent emission line near 6.5\\kev\\ and a characteristic Compton hump at about $20$\\kev.  If these features are attributed to the reprocessing of a hard X-ray (powerlaw) continuum by cold matter in an accretion disc \\citep{RN03,RF07} then the degree of broadening observed implies that the emitting region is very close to the black hole. The shape of the line profile can then give an indication of both the the radius of the emitting material from the black hole as well as the inclination of the inner accretion disc \\citep{F89,L91}. In this way \\cite{M02} obtained an estimate for the inclination of the innermost part of the accretion disc of $43$\\deg$\\pm{\\thinspace 15}$\\deg ($90\\%$ confidence). ",
        "conclusions": "We find that, in the accretion disc in V4641 Sgr, the region thought to be both the origin of the jets and the emission site of the Fe\\ka\\ line is flat and aligned with the central black hole.  Thus we would expect the inclinations measured for the jets and with the Fe\\ka\\ line to be similar.  Because there is a significant difference between the two measurements we conclude that one or both of them must be inaccurate or our model incorrect.  It is important that this system be observed in more detail to resolve this in the near future."
    },
    "0808/0808.1315.txt": {
        "abstract": " ",
        "introduction": "\\label{sect:intro}  \\subsection{Scope of the report} \\label{sec:scope-document}  The AMBER Task Force described in document [RD \\ref{rd:atfplan}] the objectives, its methodology and its plan to investigate the various AMBER issues which prevent the instrument to fulfill its initial specifications and therefore its original science program.  The objectives of the February run was mainly to bring AMBER (see AMB-ATF-001 memo) into contractual specifications the accuracy of the absolute visibility, of the differential and of the closure phase through a fundamental analysis of the instrument status and limitations.  \\subsection{Outline of the report} \\label{sec:outline-report}  Before the run, a new implementation of the AMBER software by A.~Chelli and G.Duvert has been designed. This new and more accurate software using the same philosophy as the \\texttt{amdlib} v2.1\\footnote{\\texttt{amdlib} v2.1 is the last public release of the AMBER data reduction software.} is described in Sect.\\ \\ref{sec:dataproc}.  The first days of the run were dedicated to the alignment of AMBER and characterization of its behavior. Many issues were tackled and the results are reported in Sect.\\ \\ref{sec:hardware}. Then we focused our attention on the main objective of our run: the performances limitations due to phase beating and to the lack of absolute calibration reported in Sect.\\ \\ref{sec:performances}.  We end our report with the AMBER Task Force recommendations given in Sect.\\ \\ref{sec:recommendations}.  \\emph{Note: we tried to keep the main part of the report as concise as possible and we moved the details in annexes.}  \\subsection{Documents} \\subsubsection{Applicable documents} \\noindent \\begin{tabular}{p{6ex}p{0.45\\textwidth}p{0.37\\textwidth}} Code &Title &Number\\\\ \\hline \\ADlabel{ad:techspec} &AMBER Technical Specifications &VLT-SPE-ESO-15830-2074 issue 1.0 \\\\ &&dated 20.04.2000\\\\ \\hline \\end{tabular}  \\subsubsection{Reference documents} \\noindent \\begin{tabular}{p{6ex}p{0.45\\textwidth}p{0.37\\textwidth}} Code &Title &Number\\\\ \\hline \\RDlabel{rd:atfplan} &AMBER Task Force: Objectives, Methodology, Plan &VLT-PLA-AMB-15830-7003 issue 1.0\\\\ &&dated 13/12/2008\\\\ \\RDlabel{rd:amb-atf-001} &ATF detailed plan for Feb'08 run &AMB-ATF-001 issue 1.5 dated 28/01/2008\\\\ \\RDlabel{rd:amb-igr-018} &AMBER data processing in the Image space &AMB-IGR-018 issue 1.0 dated 26/10/2000\\\\ \\RDlabel{rd:amb-det-007} &Report of the AMBER detector intervention from September 10 to 19, 2007 &AMB-DET-007 issue 1.0 dated 28/09/2007\\\\ \\RDlabel{rd:OPM-LLS} &OPM Warm Optics Design Report &VLT-TRE-AMB-15830-1001 issue  2.2\\\\ &&dated 28/05/2001\\\\ \\RDlabel{rd:OPM-TR} &OPM Warm Optics Test Report &VLT-TRE-AMB-15830-1010 issue 1.0\\\\ &&dated 10/10/2003\\\\ \\hline \\end{tabular} \\bigskip   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": ""
    },
    "0808/0808.2233_arXiv.txt": {
        "abstract": "We have studied superionization and X-ray line formation in the spectra of \\zpup\\ using our new stellar atmosphere code (XCMFGEN) that can be used to simultaneously analyze optical, UV, and X-ray observations. Here, we present results on the formation of the \\ion{O}{6}\\ $\\lambda\\lambda 1032, 1038$ doublet. Our simulations, supported by simple theoretical calculations, show that clumped wind models that assume void in the interclump space cannot reproduce the observed \\ion{O}{6} profiles. However, enough \\ion{O}{6}\\ can be produced if the voids are filled by a low density gas. The recombination of \\ion{O}{6} is very efficient in the dense material but in the tenuous interclump region an observable amount of \\ion{O}{6} can be maintained. We also find that different UV resonance lines are sensitive to different density regimes in \\zpup\\ : \\ion{C}{4} is almost exclusively formed within the densest regions, while the majority of \\ion{O}{6} resides between clumps. \\ion{N}{5} is an intermediate case, with contributions from both the tenuous gas and clumps. ",
        "introduction": "One of the surprising discoveries of the {\\it Copernicus} satellite was the strong P-Cygnii profiles of superions, such as \\ion{O}{6} and \\ion{N}{5}, in the FUV spectra of many O and B stars \\citep{sno76}. The only viable explanation for the presence of \\ion{O}{6} is Auger ionization by X-rays from \\ion{O}{4} \\citep{cas79} which is the dominant form of oxygen in many O-type stars. The X-ray emission, necessary for Auger ionization, was later detected by the first X-ray telescopes \\citep[e.g.,][]{har79, sew79}.  The origin of the stellar X-ray emission was another enigma until the ``wind-shock'' mechanism \\citep{luc80}  became the accepted explanation. Massive stars posses strong line-driven winds in which the material is accelerated by numerous C, N, O, and Fe transitions \\citep[see e.g.,][]{pau90, CAK}. It was known from the conception of the line-driven wind theory that such flows are unstable and prone to the formation of dense clumps and shocks \\citep[see e.g.,][]{owo91,luc80}. The large-scale flow energy is converted to heat in the shock fronts producing high temperature plasma. Numerical simulations confirm \\citep{fel97, owo88} that at least the soft X-ray emission of early-type stars can be explained by this mechanism.  Evidence for density inhomogeneities (or clumped winds) is provided by variability studies of both WR stars \\cite[][and references therein]{lep99}, and O stars \\cite[e.g.,][]{eve98, lep08}. Further, density inhomogeneities allow the electron-scattering wings of emission lines to be reduced to the observed level while maintaining the strength of emission lines \\citep{hil91, ham98, hil99}. More recently, \\cite{cro02} and \\cite{hil03} found that they could not simultaneously fit the H$\\alpha$  and \\ion{P}{5} $\\lambda\\lambda$1120 profiles in normal O supergiants without assuming an inhomogeneous density distribution in the wind. Using a more statistical approach, \\cite{mas03} showed that the phosphorus ionization structure was consistent with expectations only if lower than conventional mass-loss rates were used in their analysis of \\ion{P}{5} $\\lambda\\lambda$1120. Additional observational evidence for wind clumping come from {\\it Chandra} and {\\it XMM-Newton} high-resolution X-ray spectra of O stars. These spectra revealed that X-ray lines suffer less absorption in the wind than predicted by ``smooth\" models \\citep[e.g.,][and references therein]{kra03b,coh05,wal07}.  Superionization has received only limited attention since the work of \\citeauthor{mac93} \\citeyearpar{mac93, mac94}. The effect was introduced into modern stellar atmosphere codes \\citep[e.g., {\\it WM-basic},][]{pau01}, but we are unaware of any work that has revisited the question in light of the high-resolution X-ray observations, improved X-ray emission calculations and the new results on clumping. With improvements to CMFGEN, we are developing tools and techniques to move towards this goal. As part of this effort we discovered that the interclump medium is crucial to explain the \\ion{O}{6} doublet profile in $\\zeta$~Pup. In this paper we demonstrate the effect and discuss its implications. In \\S\\ref{sec:obs} we briefly describe our code, the observations we used, and our models. We present and discuss our results in \\S\\ref{sec:result}--\\S\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  In this letter we present our first results on superionization in clumped winds, and showcase the potential of the interclump medium to produce observable features. Clumped wind models that use the classical ``volume filling factor\" approach (clumps with voids in between) cannot reproduce the observed \\ion{O}{6} profile in $\\zeta$~Pup. The recombination of \\ion{O}{6} is too efficient and the necessary fractional abundance cannot be sustained in the clumps. However, a tenuous interclump medium can contribute enough \\ion{O}{6} to produce an observable \\ion{O}{6} profile. Only a small amount of mass is necessary in the interclump medium, so its overall effect on the derived mass-loss rates is negligible.  Our result highlights the need for improved treatment of clumping in the winds of massive stars. It is impossible to achieve a simultaneous fit to all UV P-Cygnii profiles with a single component wind model for \\zpup. Our simulations suggest that in $\\zeta$~Pup, different UV resonance lines probe different density regimes. \\ion{C}{4} is formed almost exclusively in the dense material, while \\ion{O}{6} likely originates from the interclump medium. \\ion{N}{5} is an intermediate case with similar contributions from both components. In cooler O stars, when N$^{2+}$ become the dominant ionization stage, we might expect that \\ion{N}{5}\\ shows the same behavior as \\ion{O}{6}\\ in the hotter stars. Obviously other possible effects of the interclump medium, in both O and W-R stars, should be investigated."
    },
    "0808/0808.2375_arXiv.txt": {
        "abstract": "We investigate general aspects of molecular line formation under conditions which are typical of prestellar cores.  Focusing on simple linear molecules, we study formation of their rotational lines by radiative transfer simulations. We present a thermalization diagram to show the effects of collisions and radiation on the level excitation. We construct a detailed scheme (contribution chart) to illustrate the formation of emission line profiles. This chart can be used  as an efficient tool to identify which parts of the cloud contribute  to a specific line profile. We show how molecular line characteristics for uniform model clouds  depend on hydrogen density, molecular column density, and kinetic temperature.  The results are presented in a 2D plane to illustrate cooperative effects of the physical factors.  We also use a core model with a non-uniform density distribution and chemical stratification to study the effects of cloud contraction and rotation on spectral line maps. We discuss the main issues that should be taken into account when dealing with interpretation and simulation of observed molecular lines. ",
        "introduction": "Star formation is a fundamental process in the universe.  Dense, gravitationally bound, and genuinely starless cores, which we call here ``prestellar cores'', are the earliest visible precursors of forming stars.  The interplay between gravitation and thermal/magnetic pressure as well as the conservation of angular momentum are all driving forces behind prestellar core evolution. While the thermal structure is determined by dust properties and molecular composition, magnetic support is sensitive to the ionization degree. The interaction and relative importance of these processes as well as the role of external forces are still not well understood \\citep[see, e.g. reviews by][]{Klessen:2004, Bergin:2007}. Thus, it is important to study the physical and chemical evolution of these cores in detail to reveal the underlying physics.  The interiors of such  prestellar cores are well-shielded from interstellar or stellar radiation, leading to low internal temperatures. Under such conditions, the main component of these cores,  H$_2$, is not easily observable. Therefore, we have to rely on indirect methods to determine the physical structure of prestellar cores, e.g., on observations of thermal dust emission or emission lines from other molecules, like CO, CS and N$_2$H$^+$.  Observations of spectral lines have an important advantage  over continuum observations since they also carry information about kinematics of the gas. A disadvantage of molecular tracers can be that they are only present under certain conditions and can be frozen-out on dust surfaces. Numerous studies, including single-dish and interferometric observations along with their theoretical analysis, have been performed over the past years, significantly deepening our understanding of the physical and chemical processes in star-forming regions \\citep[see, e.g. reviews of][]{Myers:1999,Evans:1999,DiFrancesco:2007, Bergin:2007}.  However, deriving physical properties from molecular lines is a difficult (inverse) problem  because of the many factors which can affect the line formation.  Even if we assume that the studied object is spherically symmetric and uniform, we have to specify at least 5 independent parameters describing the formation of molecular lines. These are density $n({\\rm H}_2)$, kinetic temperature $T_{\\rm kin}$, molecular column density $N({\\rm mol})$, radial velocity $V_r$, and micro-turbulent velocity $V_{\\rm turb}$. In more realistic models, additional parameters should also be considered, e.g., electron concentration and external radiation field, with spatial distributions of all  the above parameters.  An exact analytical treatment of this multi-parameter system is, in most cases, impossible. Optically thin line formation can be described analytically to some extent \\citep[see, e.g.,][]{Wilson:2000}, but it is very difficult to do the same for optically thick lines. Also, it is difficult to use analytic methods to solve the inverse problem, i.e., to restore physical distributions of relevant parameters from observed spectra because of the non-local and non-linear nature of the radiative transfer problem. This is why numerical line radiative transfer (LRT) models are commonly used as an exploratory tool. Several reliable numerical methods and numerical codes have been developed by various groups for this purpose \\citep[see reviews by][]{Peraiah:2004,Zadelhoff:2002}.  There are also fast approximate numerical tools available for the molecular line radiative transfer analysis\\footnote{{\\it http://www.sron.rug.nl/\\~{}vdtak/radex/radex.php}} \\citep{vanderTak:2007}. However, input data for line modeling include not only physical conditions but also distributions of molecular abundances. All these data can be represented with some analytical prescription or extracted from chemical and dynamical models.  As recently outlined by \\citet{tsamis}, there are two alternative approaches to the diagnostics of protostellar objects. The first approach is a systematic study of the influence of various factors on emergent spectra in order to facilitate the analysis of forthcoming observations. Such studies (based mainly on approximate LRT methods) have already started more than 30 years ago  with the analysis of the excitation conditions for various molecules and effects of the density, temperature and velocity gradients on the emission line profiles \\citep[e.g.,][]{Lucas:1974,Goldreich:1974,Leung:1978,Stenholm:1980}. The alternative approach is a detailed study of an individual source. To infer its parameters, a number of models is constructed and the direct LRT problem is solved  for each model. Varying the model parameters, it is possible to find  the combination which provides the best agreement between observed and modeled spectra (maps) or their derived parameters, like, e.g., velocity centroids \\citep{walker1994} or line ratios \\citep{vanderTak:2007}. Both trial-and-error methods or any sophisticated numerical algorithm for isolating the ``best''\\  set of free parameters can be utilized \\citep{Keto:2004}. This second approach has already been successfully applied in a number of studies, aimed to extract detailed characteristics of pre-stellar and protostellar objects from observed spectra, in particular, their chemical and kinematical structure \\citep[e.g.,][]{Tafalla:2002,Tafalla:2004,Keto:2004, Evans:2005,Brinch:2007} as well as to test different star formation theories \\citep[e.g.][]{Pavlyuchenkov:2003, Offner:2007}.   Following the first approach, efforts of various authors  are mainly concentrated toward evolutionary stages later than the prestellar phase. Most cores investigated in detail are closer to the formation of a first hydrostatic core, when the collapse  is well-developed and can be described by Shu or Larson-Penston solutions.  This interest may be partly caused by the fact that first proofs of infall have been reported for Class\\,0 objects \\citep{walker1986,b335}, in which  the collapse has already resulted in the formation of a deeply embedded source. \\cite{zhou1992} considered spectral differences between Shu and Larson-Penston solutions, assuming constant molecular abundances, and identified four line properties, supposedly unique to collapsing cores. \\citet{rawlings1992} combined a dynamical  description, based on the Shu model, with a detailed chemical model and identified species  with broad line wings, indicative of infall motions. \\citet{ry2001} coupled a similar chemical and dynamical model to a more realistic radiation transfer model, implemented as an approximate $\\Lambda$-iteration code, and studied the sensitivity of emergent central optically thin and optically thick spectra to various model parameters. The assumption of  isothermality, used by \\citet{rawlings1992} and \\citet{ry2001}, was relaxed in \\citet{tsamis} in application to a ``B335-like'' Class~0 object. The sensitivity of line  profiles in this model to the intensity of the ambient radiation field was studied in \\citet{redman2004}. The influence of various parameters on Class~0 central spectra, with a particular emphasis on turbulence and rotation, was also studied by \\citet{wtb2001} under the  assumption of flat abundance distributions. The role of high angular resolution in studying such clouds was investigated by \\citet{choi}. More details about line modeling can be also found in review by \\citet{Evans:1999}.   Large amounts of data have also been accumulated for starless cores, most of which are presumably ``prestellar'' \\citep[e.g.,][]{infall}, together with the theoretical modeling of molecular lines. E.g., \\citet{Tafalla:2004} investigated the effects of depletion on the line intensity. \\citet{Lee:2004} studied the evolution of line profiles during core contraction based on the models where the chemical evolution is calculated along with dynamical evolution of the cloud. \\citet{Pavlyuchenkov:2007} investigated the combined effects of temperature and depletion, and \\citet{DeVries:2005} investigated the possibility to use approximate analytical models to extract velocity gradients from the line profiles. In principle, all the mentioned studies of line formation in protostellar objects are also relevant for prestellar cores. However, despite the obvious observational and theoretical progress, the interpretation of the data is still far from being straightforward, as infall and rotation velocities in prestellar cores are both comparable or even smaller than the sound speed. Difficulties in restoring the information about structural, thermal, and kinematic properties  are also caused by the lack of relevant methods for molecular line analysis.  In this paper, we systematically study the formation of molecular lines in prestellar cores and present different tools which can be used to analyze the formation of lines in detail. Using the linear molecules CO and HCO$^+$ as examples, we model their emission by means of non-LTE LRT simulations. The role of hydrogen density, molecular column density, kinetic temperature, infall, and rotation is examined and illustrated in the paper. Such an analysis may form the basis of a more sophisticated interpretation of observational data and can be useful for those who are going to use LRT simulations in their studies or want to get a global view of the factors influencing observed line profiles.   In Section~2 we describe the radiation transfer model used in the paper, introduce parameters to characterize line profiles, and provide some basic considerations on the line formation in prestellar cores. In Section~3 we show how molecular line characteristics depend on hydrogen density, molecular column density, and kinetic temperature for the parameter sample of uniform model clouds. In Section~4 we consider a non-uniform model cloud and study the effects of chemical stratification, contraction and rotation on spectral line maps. In Section~5 we discuss  additional problems related to the LRT analysis. Section~6 summarizes the main conclusions of this paper. ",
        "conclusions": " \\begin{itemize} \\item Densities in starless cores fall into the range where level populations are neither radiatively, nor collisionally dominated. \\item Large column densities do not necessarily lead to the appearance of self-absorption dips. \\item When the column density is fixed, a specific line can be optically thick only in a range of densities, being optically thin at densities both below and above this range. \\item The density which is ``traced'' by some transition depends on external factors (UV field and CR ionization), which shape the molecular distribution.  In particular, the ``traced'' density can be lower than the critical density. \\item Rotation and infall  may produce very similar spectra and in general can only be distinguished by spectral mapping. \\end{itemize}"
    },
    "0808/0808.2469_arXiv.txt": {
        "abstract": "We describe a new survey for unbound hypervelocity stars (HVSs), stars traveling with such extreme velocities that dynamical ejection from a massive black hole is their most likely origin. We investigate the possible contribution of unbound runaway stars, and show that the physical properties of binaries constrain low mass runaways to bound velocities. We measure radial velocities for HVS candidates with the colors of early A-type and late B-type stars. We report the discovery of 6 unbound HVSs with velocities and distances exceeding the conservative escape velocity estimate of Kenyon and collaborators. We additionally report 4 possibly unbound HVSs with velocities and distances exceeding the lower escape velocity estimate of Xue and collaborators.  These discoveries increase the number of unbound HVSs by 60\\% - 100\\%. Other survey objects include 19 newly identified $z\\sim2.4$ quasars. One of the HVSs may be a horizontal branch star, consistent with the number of evolved HVSs predicted by Galactic center ejection models.  Finding more evolved HVSs will one day allow a probe of the low-mass regime of HVSs and will constrain the mass function of stars in the Galactic center. ",
        "introduction": "Three-body interactions with a massive black hole (MBH) will inevitably unbind stars from the Galaxy \\citep{hills88, yu03}.  In 2005 we reported the discovery of the first HVS:  a 3 M$_{\\sun}$ main sequence star traveling with a Galactic rest frame velocity of at least $+700\\pm12$ km s$^{-1}$, more than twice the Milky Way's escape velocity at the star's distance of 110 kpc \\citep{brown05}. This star cannot be explained by normal stellar interactions: the maximum ejection velocity from binary disruption mechanisms \\citep{blaauw61, poveda67} is limited to $\\lesssim$300 km s$^{-1}$ for 3 M$_{\\sun}$ stars \\citep{leonard91, leonard93, tauris98, portegies00, davies02, gualandris05}.  Although runaways may reach unbound velocities for very massive stars, like the hyper-runaway HD 271791 \\citep{heber08, przybilla08c}, runaway ejection velocities are constrained by the properties of binary stars. A massive and compact object is needed to accelerate low mass stars to unbound velocities.  There is overwhelming evidence for a $\\sim4\\times10^6$ M$_{\\sun}$ MBH in the dense stellar environment of the Galactic center \\citep{schodel03, ghez08}. \\citet{hills88} coined the term HVS to describe a star ejected by the MBH.  The observational signature of a HVS is its unbound velocity.  Although not all unbound stars are necessarily HVSs -- fast-moving pulsars, for example, are explained by supernova kicks \\citep[e.g.][]{arzoumanian02} -- unbound low mass main sequence stars are most plausibly explained as HVSs.  Here we introduce a new HVS survey using the MMT to target HVS candidates with masses down to $\\sim$2 M$_{\\sun}$.  Discovering lower-mass HVSs should provide constraints on the stellar mass function of HVSs \\citep{brown06, kollmeier07}; the velocity distribution of low- versus high-mass HVSs may discriminate between a single MBH or binary MBH origin \\citep{sesana07b, kenyon08}.  Our survey strategy targets stars fainter and redder than the original HVS survey.  This strategy is successful:  we report the discovery of 6 unbound HVSs and 4 possibly unbound HVSs.  \\subsection{Recent Observations of HVSs}  Observers have identified a remarkable number of HVSs in the past 3 years. Following the discovery of the first HVS \\citep{brown05}, \\citet{hirsch05} reported a helium-rich subluminous O star leaving the Galaxy with a rest-frame velocity of at least $+717$ km s$^{-1}$.  \\citet{edelmann05} reported an 9 M$_{\\sun}$ main sequence B star with a Galactic rest frame velocity of at least $+548$ km s$^{-1}$, possibly ejected from the Large Magellanic Cloud.  \\citet{brown06, brown06b, brown07a, brown07b} reported 7 B-type HVSs discovered in a targeted HVS survey, along with evidence for an equal number of bound HVSs ejected by the same mechanism.  High-dispersion spectroscopy has shed new light on the nature of the HVSs. \\citet{przybilla08b} have recently shown that HVS7 is a chemically peculiar B main sequence star, with an abundance pattern unusual even for the class of peculiar B stars.  The star HVS3 (HE 0437-5439), the unbound HVS very near the LMC on the sky, has received the most attention.  HVS3 is a 9 M$_{\\odot}$ B star of half-solar abundance, a good match to the abundance of the LMC \\citep{bonanos08, przybilla08}. Stellar abundance may not be conclusive evidence of origin, however.  A- and B-type stars exhibit 0.5 - 1 dex scatter in elemental abundances within a single cluster, due to gravitational settling and radiative levitation in the atmospheres of the stars \\citep{varenne99, monier05, fossati07, gebran08a, gebran08b}.  An LMC origin requires that HVS3 was ejected from the galaxy at $\\sim$1000 km s$^{-1}$ \\citep{przybilla08}, a velocity that can possibly come from three-body interactions with an intermediate mass black hole in a massive star cluster \\citep{gualandris07, gvaramadze08}.  \\citet{perets08b} shows that the ejection rate of 9 M$_{\\sun}$ stars, however, is four orders of magnitude too small for this explanation to be plausible.  The alternative explanation is that HVS3 is a blue straggler, ejected by the Milky Way's MBH.  Theorists argue that a single MBH or a binary MBH can eject a compact binary star as a HVS \\citep{lu07, perets08b}; the subsequent evolution of such a compact binary can result in mass-transfer and/or a merger that can possibly explain HVS3 \\citep{perets08b}.  Proper motion measurements, underway now, will determine HVS3's origin.  Other recent HVS work highlights the link between stellar rotation and the origin of HVSs.  Main sequence B stars have fast mean $v\\sin{i}\\sim150$ km s$^{-1}$ \\citep[e.g.][]{abt02, huang06a}.  Hot blue horizontal branch (BHB) stars have slow mean $v\\sin{i}<10$ km s$^{-1}$ (because they have just evolved off the red giant branch \\citep[e.g.][]{behr03, behr03b}).  Interestingly, \\citet{hansen07} predicts that main sequence HVSs ejected by the Hills mechanism should be slow rotators, because stars in compact binaries have systematically lower $v\\sin{i}$ due to tidal synchronization.  \\citet{lockmann08}, on the other hand, predict that HVSs should be fast rotators, at least for single stars spun up and ejected by a binary black hole. To date HVS1, HVS3, HVS7, and HVS8 have observed $v\\sin{i}$ of $\\sim$190, $55\\pm2$, $55\\pm2$, and $260\\pm70$ km s$^{-1}$, respectively \\citep{heber08b, przybilla08, przybilla08b, lopezmorales08}.  As discussed by \\citet{perets07c}, we clearly require a larger sample of HVSs to measure the distribution of HVS rotations and discriminate HVS ejection models.  In \\S 2 we describe our new HVS survey strategy and summarize our observations.  In \\S 3 we discuss the Galactic escape velocity, our definition of a HVS, and the possible contribution of hyper-runaways to the population of unbound stars.  In \\S 4 we present the new unbound HVSs.  In \\S 5 we discuss a possible BHB star among the HVSs.  We conclude in \\S 6.  \\begin{figure}\t\t% \\plotone{f1.eps} \\caption{ \\label{fig:ugr} Color-color diagram illustrating the target selection for our new HVS survey ({\\it long dashed line}) and our old HVS survey ({\\it short dashed line}).  The six new HVSs ({\\it solid stars}) and four possible HVSs ({\\it solid dots}) scatter around the \\citet{girardi04} stellar evolution tracks for 2 - 4 M$_{\\sun}$ main sequence stars ({\\it solid lines}).  Average color uncertainties are indicated by the errorbar on the upper right.  Previous HVS discoveries ({\\it open stars}) and the \\citet{xue08} BHB sample ({\\it x's}) are also marked.} \\end{figure} ",
        "conclusions": "We describe a new targeted HVS survey, a spectroscopic survey of faint stars $19<g'_0<20.5$ with early A-type and late B-type colors. Recent observations confirm that 3 of our B-type HVSs are 3-4 M$_{\\sun}$ main sequence stars.  The observational signature of a HVS is its unbound velocity, which we determine by comparing observed radial velocities and distances to Galactic potential models.  We argue that the known properties of binaries and the rarity of massive stars make hyper-runaways like HD 271791 rare.  A MBH ejection remains the most plausible origin of unbound low-mass stars.  Our HVS survey is 59\\% complete and, combined with our original HVS survey, shows a remarkable velocity distribution:  26 stars with $v_{rf}>+275$ km s$^{-1}$ and only 2 stars with $v_{rf}<-275$ km s$^{-1}$. Here we report the discovery of 6 new unbound HVSs in excess of the conservative escape velocity model of \\citet{kenyon08}, and 4 additional unbound HVSs in excess of the escape velocity model of \\citet{xue08}.  One of the new HVSs may be an evolved BHB star.  The \\citet{kenyon08} ejection models predict BHB HVSs are $\\sim$10 times less abundant than the main sequence HVSs in our survey, consistent with the existence of $1\\pm1$ BHB stars in our HVS sample.  Of course, the exact number of BHB HVSs depends on the mass function of stars near the central MBH. BHB HVSs therefore have the potential to probe the low-mass regime of HVSs and constrain the mass function of stars in the Galactic center.  HVSs are fascinating because their properties are tied to the nature and environment of the MBH that ejects them \\citep{levin06, baumgardt06, merritt06, ginsburg06, ginsburg07, demarque07, gualandris07, sesana06, sesana07, sesana07b, sesana07c, lu07, kollmeier07, hansen07, perets07, perets07c, perets08b, perets08a, sherwin08, svensson08, oleary08, lockmann08}, and their trajectories probe the dark matter halo through which they move \\citep{gnedin05, yu07, wu08, kenyon08}. The angular distribution of HVSs on the sky reveals significant anisotropy that may also be related to the Galactic potential \\citep{brown08d}. Our ultimate goal is to find a statistical sample of $\\sim$100 HVSs to measure the distribution of HVS properties and discriminate HVS ejection models.  ~"
    },
    "0808/0808.1140_arXiv.txt": {
        "abstract": "We report on the variability of 443 flat spectrum, compact radio sources monitored using the VLA for 3 days in 4 epochs at $\\sim 4$ month intervals at 5 GHz as part of the Micro-Arcsecond Scintillation-Induced Variability (MASIV) survey.  Over half of these sources exhibited 2-10\\% rms variations on timescales over 2 days.  We analyzed the variations by two independent methods, and find that the rms variability amplitudes of the sources correlate with the emission measure in the ionized Interstellar Medium along their respective lines of sight.  We thus link the variations with interstellar scintillation of components of these sources, with some (unknown) fraction of the total flux density contained within a compact region of angular diameter in the range 10-50$\\mu$as. We also find that the variations decrease for high mean flux density sources and, most importantly, for high redshift sources. The decrease in variability is probably due either to an increase in the apparent diameter of the source, or a decrease in the flux density of the compact fraction beyond $z \\sim 2$.  Here we present a statistical analysis of these results, and a future paper will the discuss the cosmological implications in detail. ",
        "introduction": "The discovery of centimeter-wavelength Intra-Day Variability (IDV) or ``flickering'' in  Active Galactic Nuclei (AGN) by \\citet{hee84} initially raised concerns that some AGN possess brightness temperatures over six orders of magnitude above the $10^{12}\\,$K inverse Compton limit for incoherent synchrotron emission \\citep[e.g.][]{qui89}. However, considerable evidence has now accumulated to demonstrate that interstellar scintillation (ISS) in the turbulent, ionized interstellar medium (ISM) of our Galaxy is the principal mechanism responsible for the IDV observed in AGN, as was proposed by \\citet{HR87}.  Two more recent observational techniques provide compelling evidence for the prevalence of ISS.  Time delays of 1--8\\,min are observed in the arrival times of the flux density variations between telescopes on different continents for the three intra-hour variable sources B0405--385, B1257--326 and J1819$+$3845 \\citep{jau2000,big2006,dtdb2002}.  The delay arises due to the finite time required for the stochastic fluctuations associated with the ISM to drift across the Earth.  A second observational signature of ISS relates to the modulation of IDV variability timescales with a period of exactly one year.  This arises because the Earth's orbital motion about the Sun contributes to the effective velocity with which the interstellar scattering material moves relative to an Earth-bound observer; the variations are slow as the Earth moves parallel to the material and fast as it moves anti-parallel to it. Annual cycles in IDV variability timescales are reported in at least seven sources \\citep[e.g.][]{den2003,ric2001,jau2001,big2003,jau2003}, including several whose long variability timescales preclude detection of time delays in the scintillation pattern over intercontinental distances.  For many lines of sight through the ISM the slowest variations are expected in September if the motion of the turbulent material is comparable to the local standard of rest (LSR).  The recognition of ISS as the dominant cause of IDV has not entirely alleviated the brightness temperatures problems posed by these sources. A source must be small to scintillate; in the weak scintillation case most frequently observed at frequencies near 5 GHz \\citep{wal98}, the source angular size must be comparable to or smaller than the angular size of the first Fresnel zone, $\\theta_{\\rm F}=\\sqrt{c/2 \\nu \\pi L}$. Here $L$ is the distance to the scattering region, which we will refer to as the screen even though in some cases it may be better described as a slab extending from the Earth out to distance $\\sim L$. $\\theta_{\\rm F}$ is typically tens of microarcseconds for screen distances of tens to hundreds of parsecs, which implies source components with angular sizes two to three orders of magnitude finer than the scales probed by VLBI. The long time-scale over which IDV has been observed in some sources suggests that such scintillating components can be relatively long-lived despite their small physical sizes.   This paper reports on the results of a Micro-Arcsecond Scintillation-Induced Variability (MASIV) survey for IDV in AGN. This year-long survey conducted observations of between 500--700 AGN over each of four epochs of three or four days duration in 2002 and 2003 at 4.9\\,GHz with the VLA. The aim of the survey was to provide a catalogue of at least 100 AGN which vary on timescales of hours to days to provide the basis of detailed studies of the IDV AGN population drawn from a well-defined sample.  A description of the observations in epochs 2, 3 and 4 is presented in \\S\\ref{obs} as a supplement to descriptions of the first epoch observations and MASIV source selection in \\citet{lov2003} (hereafter Paper 1). In \\S\\ref{VarClass} we describe how the time series of flux density for each source in each epoch was classified as variable or non-variable.  In \\S\\ref{DsClass} we describe how we have quantified the amplitude and timescale for the variations from an analysis of the structure function combined from all 4 epochs and apply corrections for noise and other sources of flux density error.  The basic hypothesis of the paper that the variations are predominantly due to interstellar scintillation is presented and examined in \\S\\ref{sec:ISS}, including the influence of parameters of the interstellar medium (\\S\\ref{VarB} emission measure and galactic latitude) and of the sources themselves (\\S\\ref{VarAlpha} mean flux density, spectral index). We now have redshifts for more than half of the sources and we present the dependence of the variability on redshift in \\S\\ref{VarZ} - a result which shows that the MASIV survey provides a new cosmological probe.  In \\S\\ref{CompareVar} we consider whether the variability is intermittent over the four epochs. Our data are listed in Table \\ref{tab:BigTable} and our conclusions are presented in \\S\\ref{Conc}. ",
        "conclusions": "\\label{Conc}  We have reported results from the four epochs of the MASIV survey.  There were 710 sources with flat spectra ($\\alpha<-0.3$) near 5~GHz selected in weak and strong flux density groups surveyed for variability in four epochs over a year.  These flat spectrum sources are predominantly quasars with compact emission probably from a core and jet, many with effective diameters small enough to show interstellar scintillation (ISS).  In each epoch the flux density was measured using sub-arrays of the VLA every~2 hrs for about 12~hrs each day for 3-4~days. Sources were removed from the study if they showed evidence for changing correlated flux density due to confusion or resolution of their more extended structure, leaving 443 sources which were analyzed for variability in two ways.  The first was a binary classification based on the raw modulation index (visual method) in which 43\\% of the sources were classified as variable in 2, 3 or 4 of the epochs. The second was a fit to the epoch-averaged structure function parameterized by $D(2\\mbox{d})$ and a timescale $\\tau_{\\rm char}$. In view of the uncertainties in the latter we classified sources as slow ($>3$~days), medium (0.5-3~days) or fast ($<0.5$~days) if $D(2\\mbox{d})$ exceeded $4\\times10^{-4}$.  By this criterion 37\\% of the sources varied with more than 1.4\\% modulation index over 2~days, which is similar to the 43\\% variables by the visual classification.  We found that $D(2\\mbox{d})$ and timescale varied both with coordinates in the Galaxy and also with source-based quantities. This confirms that the variations are dominated by ISS, which depends on both the strength of scattering and the distance to the scattering region and also on the fraction of flux density in its most compact component and its effective angular diameter.  The following is a summary of our findings:  \\begin{itemize}   \\item The amplitude of 2-d  variability increases with increasing emission measure estimated from H$\\alpha$ intensity for each line of sight. Emission measure is the column density for the square of the electron density which is expected to be strongly correlated with inhomogeneity in the ionized medium that causes ISS. This result provides observational evidence that ISS is the dominant cause of the variations.  We find fast variations dominate for low emission measure, as expected since such regions will be seen out of the plane and closer to the Earth, and slow variations dominate for high emission measure which are typically seen at greater distances toward the Galactic plane especially for southern latitudes where the H$\\alpha$ intensity is high.  \\item The amplitude of 2-d ISS variability varies significantly with Galactic latitude but differs substantially between positive and negative latitudes.  The expected behaviour is complicated;  greater path lengths at low latitudes, where the scattering should be stronger, cause the scattering to be slower which should reduce the rms over 3 d. However, the observed timescales show that there are more sources with fast variations at high latitudes and more sources with slow variations at low latitudes in both hemispheres, in clear support of ISS as the dominant cause.  \\item The ISS modulation index tends to decrease with increasing mean flux density, as expected if the compact emission is limited by synchrotron self absorption or inverse Compton losses to have a maximum brightness temperature.  In that case the expected angular diameter $\\propto \\sqrt{\\bar{S}}$ which will quench the ISS of the stronger sources.  \\item There is little change in the ISS amplitude with spectral index for our sample with $\\alpha > -0.3$.  \\item There is evidence that the ISS can be intermittent on times of 3-6 months for some sources, but this is hard to quantify from the 3~day observing sequences, when the time scale of the variations is of the same order.  \\item  We model $D(2\\mbox{d})$ as a function of compact source component fractional flux density and angular diameter, from which we find compact diameter to lie in the range 0.005 -- 0.15 milli arcseconds and brightness temperatures in the range $10^{12} - 10^{14}$K.  \\item The most far-reaching result reported here is the discovery of a decrease in both the fraction of sources that scintillate and in their scintillation amplitude beyond redshifts around 2.  We conclude that there is an increase in the typical angular diameter of the most compact radio-emitting regions of the quasars beyond reshift 2. The possible interpretations of this exciting result will be presented in a companion paper \\citep{macq08}.  \\item A further surprise (at least to us) was the apparent absence of the very rapid variables (IHV). J1819+3845 fell in our sample, but it was the only source to show remarkable large amplitude variability. J0929+5013 showed rapid variability in the January 2002 epoch \\citep{lov2003} but, although monitored closely, revealed only slower, many-hour variability in the three later epochs. We had expected to find more of these rapid variables especially given that two of the three known, J1819+3845 and PKS1257-326 were discovered serendipitously.  This strongly suggests that the IHV sources lie behind discrete local interstellar clouds which cover a small fraction of the sky.  \\end{itemize}"
    },
    "0808/0808.2432_arXiv.txt": {
        "abstract": "Spherical models of collisionless but quasi-relaxed stellar systems have long been studied as a natural framework for the description of globular clusters. Here we consider the construction of self-consistent models under the same physical conditions, but including explicitly the ingredients that lead to departures from spherical symmetry. In particular, we focus on the effects of the tidal field associated with the hosting galaxy. We then take a stellar system on a circular orbit inside a galaxy represented as a ``frozen\" external field. The equilibrium distribution function is obtained from the one describing the spherical case by replacing the energy integral with the relevant Jacobi integral in the presence of the external tidal field. Then the construction of the model requires the investigation of a singular perturbation problem for an elliptic partial differential equation with a free boundary, for which we provide a method of solution to any desired order, with explicit solutions to two orders. We outline the relevant parameter space, thus opening the way to a systematic study of the properties of a two-parameter family of physically justified non-spherical models of quasi-relaxed stellar systems. The general method developed here can also be used to construct models for which the non-spherical shape is due to internal rotation. Eventually, the models will be a useful tool to investigate whether the shapes of globular clusters are primarily determined by internal rotation, by external tides, or by pressure anisotropy. ",
        "introduction": "Large stellar systems can be studied as collisionless systems, by means of a one-star distribution function obeying the combined set of the collisionless Boltzmann equation and the Poisson equation, under the action of the self-consistent mean potential. For elliptical galaxies the relevant two-star relaxation times do actually exceed their age; an imprint of partial relaxation may be left at the time of their formation \\citetext{if we refer to a picture of formation via incomplete violent relaxation; \\citealp{Lyn67}, \\citealp{Alb82}}, but otherwise they should be thought of as truly collisionless systems, generally characterized by an anisotropic pressure tensor. In turn, for globular clusters the relevant relaxation times are typically shorter than their age, so that we may argue that for many of them the two-star relaxation processes have had enough time to act and to bring them close to a thermodynamically relaxed state, with their distribution function close to a Maxwellian. This line of arguments has led to the development of well-known dynamical models for globular clusters \\citep{Kin65,Kin66}.  King models are based on a quasi-Maxwellian isotropic distribution function $f_K(E)$ in which a truncation prescription, continuous in phase space, is set heuristically to incorporate the presence of external tides; but otherwise, they are fully self-consistent (i.e., no external fields are actually considered) and perfectly spherical. Empirically, the simplification of spherical symmetry is encouraged by the fact that in general globular clusters have round appearance. Indeed, as a zero-th order description, these models have had remarkable success in applications to observed globular clusters \\citep[e.g., see][and references therein]{Spi87, DjoMey94,McLMar05}. In recent years, great progress has been made in the acquisition of detailed quantitative information about the structure of these stellar systems, especially in relation to the measurement of the proper motions of thousands of individual stars \\citep[see][]{Lee00,McL06}, with the possibility of getting a direct 5-dimensional view of their phase space. Such progress calls for renewed efforts on the side of modeling. More general models would allow us to address the issue of the origin of the observed departures from spherical symmetry. In fact, it remains to be established which physical ingredient among rotation, pressure anisotropy, and tides is the primary cause of the flattening of globular clusters \\citep[e.g., see][and references therein]{Kin61, FreFal82,Gey83,WhiSha87,DavPru90,HanRyd94,Ryd96,Goo97,Ber08}  As in the case of the study of elliptical galaxies \\citep[e.g., see] [and references therein]{BerSti93}, different approaches can be taken to the construction of models. Broadly speaking, two complementary paths can be followed. In the first, ``descriptive\" approach, under suitable geometrical (on the intrinsic shape) and dynamical (e.g., on the absence or presence of dark matter) hypotheses, the available data for an individual stellar system are imposed as constraints to derive the internal orbital structure (distribution function) most likely to correspond to the observations. This approach is often carried out in terms of codes that generalize a method introduced by \\citet{Sch79}; for an application to the globular cluster $\\omega$ Cen, see \\citet{Ven06}. In the second, ``predictive\" approach, one proposes a formation/evolution scenario in order to identify a physically justified distribution function for a wide class of objects, and then proceeds to investigate, by comparison with observations of several individual objects, whether the data support the general physical picture that has been proposed. Indeed, King models belong to this latter approach.  The purpose of this paper is to extend the description of quasi-relaxed stellar systems, so far basically limited to the spherical King models, to the non-spherical case. There are at least three different ways of extending spherical isotropic models of quasi-relaxed stellar systems (such as King models), by modifying the distribution function so as to include: (i) the explicit presence of a non-spherical tidal field; (ii) the presence of internal rotation; (iii) the presence of some pressure anisotropy. As noted earlier in this Introduction, these correspond to the physical ingredients that, separately, may be thought to be at the origin of the observed non-spherical shapes. We will thus focus on the construction of physically justified models, as an extension of the King models, in the presence of external tides and, briefly, on the extension of the models to the presence of rigid internal rotation. A first-order analysis of the triaxial tidal problem addressed in this paper was carried out by \\citet{HegRam95} and the effect of a ``frozen'' tidal field on (initially) spherical King models was studied by \\citet{Wei93} using N-body simulations. For the axisymmetric problem associated with the presence of internal rotation, a first-order analysis of a simply truncated Maxwellian distribution perturbed by a rigid rotation was given by \\citet{KorAna71}; different models where differential rotation is considered were proposed by \\citet{PreTom70} and by \\citet{Wil75}, also in view of extensions to the presence of pressure anisotropy, which goes beyond the scope of this paper. Models that represent a direct generalization of the King family to the case with differential rotation have also been examined, with particular attention to their thermodynamic properties \\citep{LagLon96,LonLag96}. In principle, the method of solution that we will present below can deal with the extension of other spherical isotropic models with finite size that are not of the King form \\citetext{e.g., see \\citealp{WooDic62,Dav77}; see also the interesting suggestion by \\citealp{Mad96}}.  The study of self-consistent collisionless equilibrium models has a long tradition not only in stellar dynamics, but also in plasma physics \\citep[e.g., see][]{Har62,AttPeg99}. We note that in both research areas a study in the presence of external fields, especially when the external field is bound to break the natural symmetry associated with the one-component problem, is only rarely considered.  The paper is organized as follows. Section 2 introduces the reference physical model, in which a globular cluster is imagined to move on a circular orbit inside a host galaxy treated as a frozen background field; the modified distribution function for such a cluster is then identified and the relevant parameter space defined. In Sect.~3 we set the mathematical problem associated with the construction of the related self-consistent models. For models generated by the spherical $f_K(E)$, Sect.~4 and Appendix A give the complete solution in terms of matched asymptotic expansions. Alternative methods of solutions are briefly discussed in Sect.~5. The concluding Section 6 gives a summary of the paper, with a short discussion of the results obtained. In Appendix B we show how the method developed in this paper can be applied to construct quasi-relaxed models flattened by rotation in the absence of external tides. In Appendix C we show how the method can be applied to other isotropic truncated models, different from King models.  Technically, the mathematical problem of a singular perturbation with a free boundary that is faced here is very similar to the problem noted in the theory of rotating stars, starting with Milne \\citep[see][]{Tas78,Mil23,Cha33,Kro42,ChaLeb62,MonRox65}. The problem was initially dealt with inadequate tools; a satisfactory solution of the singular perturbation problem was obtained only later, by \\citet{Smi75,Smi76}. ",
        "conclusions": "Spherical King models are physically justified models of quasi-relaxed stellar systems with a truncation radius argued to ``summarize\" the action of an external tidal field. Such simple models have had great success in representing the structure and dynamics of globular clusters, even though the presence of the tidal field is actually ignored. Motivated by these considerations and by the recent major progress in the observations of globular clusters, in this paper we have developed a systematic procedure to construct self-consistent non-spherical models of quasi-relaxed stellar systems, with special attention to models for which the non-spherical shape is due to the presence of external tides.  The procedure developed in this paper starts from a distribution function identified by replacing, in a reference spherical model, the single star energy with the relevant Jacobi integral, thus guaranteeing that the collisionless Boltzmann equation is satisfied. Then the models are constructed by solving the Poisson equation, an elliptic partial differential equation with free boundary. The procedure is very general and can lead to the construction of several families of non-spherical equilibrium models. In particular, we have obtained the following results:  \\begin{itemize} \\item We have constructed models of quasi-relaxed triaxial stellar systems in which the shape is due to the presence of external tides; these models reduce to the standard spherical King models when the tidal field is absent. \\item For these models we have outlined the general properties of the relevant parameter space; in a separate paper (Varri \\& Bertin, in preparation) we will provide a thorough description of this two-parameter family of models, also in terms of projected quantities, as appropriate for comparisons with the observations. \\item We have given a full, explicit solution to two orders in the tidal strength parameter, based on the method of matched asymptotic expansions; by comparison with studies of analogous problems in the theory of rotating polytropic stars, this method appears to be most satisfactory. \\item We have also discussed two alternative methods of solution, one of which is based on iteration seeded by the spherical solution; together with the use of dedicated $N$-body simulations, the ability to solve such a complex mathematical problem in different ways will allow us to test the quality of the solutions in great detail. \\item By suitable change of notation and physical re-interpretation, the procedure developed in this paper can be applied to the construction of non-spherical quasi-relaxed stellar systems flattened by rotation (see Appendix B). \\item The same procedure can also be applied to extend to the triaxial case other isotropic truncated models (such as low-$n$ polytropes), that is models that do not reduce to King models in the absence of external tides (see Appendix C). \\end{itemize}  We hope that this contribution, in addition to extending the class of self-consistent models of interest in stellar dynamics, will be the basis for the development of simple quantitative tools to investigate whether the observed shape of globular clusters is primarily determined by internal rotation, by external tides, or by pressure anisotropy."
    },
    "0808/0808.2574_arXiv.txt": {
        "abstract": "To reconstruct dark energy models the redshift $z_{eq}$, marking the end of radiation era and the beginning of matter-dominated era, can play a role as important as $z_{t}$, the redshift at which deceleration parameter experiences a signature flip. To implement the idea we propose a variable equation of state for matter that can bring a smooth transition from radiation to matter-dominated era in a single model. A popular $\\Lambda \\propto \\rho$ dark energy model is chosen for demonstration but found to be unacceptable. An  alternative $\\Lambda \\propto \\rho a^{3}$ model is proposed and found to be more close to observation. ",
        "introduction": "The current trend in modern cosmology is mainly focussed on issues of possible explanations of the observed late time acceleration in the expansion of our Universe\\citep{riess,perl}. The presence of a dark energy of unknown origin which dominates over matter in our Universe is now almost accepted. Unlike ordinary matter the dark energy has a repulsive effect and its dominance accounts for the observed acceleration. Several candidates of dark energy are found in literature. Among them the most elementary is the cosmological constant $\\Lambda$\\citep{padd,peebles}. Variable $\\Lambda$ models\\citep{car,wet,arbab,paddy,vis,shap,dutta} were introduced later in an effort to overcome the fine tuning problem. Among other candidates most popular are different kinds of scalar fields like quintessence\\citep{caldwell1}, K-essence\\citep{chiba}and phantom fields\\citep{caldwell2} to name only a few. The Chaplygin gas\\citep{kamen} has an equation of state that can produce negative pressure in later phase of evolution to account for the accelerated expansion of the Universe. Some geometrical origin of the late-time acceleration are also found in modified gravity models\\citep{nojiri}. All these models have their own merits and demerits. The exact nature of dark energy is therefore yet to be ascertained. For a detailed review on different dark energy candidates one may see the work of \\citet{sahni} or \\citet{copeland}.  Most of these models are usually studied only in the matter dominated phase assuming the presence of some kind of dark energy to explain the transition from deceleration to acceleration. This is due to the fact that radiation era was very short lived and dominated the Universe in its early phase of evolution whereas acceleration is a more recent phenomena which has occurred very late in the matter dominated epoch. The viability of such models are usually checked by comparing the variation of deceleration parameter $q$ with the red-shift $z$ obtained theoretically with that from observations. A good estimation of $z=z_t$, the redshift marking the transition from deceleration to acceleration, is available from observation. But observational data are limited over a small range of $z$ values around $z_t$. Hence many models that fit observational data are found to behave quite differently outside this range. To choose among these models we need at least another event that happens to be outside this range of $z$, preferably in the distant past. We point out that a very good estimate of $z=z_{eq}$, the redshift when matter density was equal to radiation density, is at our disposal. Thus instead of a single point we now have two well separated points through which the $q$ vs. $z$ curve for any viable cosmological model should pass. Such a model would be more reliable as its viability could be ensured over a large range of $z$ values. But then to estimate $z_{eq}$ we must have both radiation and matter in a single model. \\par In the next section we propose a variable equation of state for matter that mimics radiation in the past and pressureless dust in the later course of evolution. In section 3 we examine the $\\Lambda \\propto \\rho$ model with the proposed equation of state for matter. In section 4 an alternative variable $\\Lambda $ model is proposed. ",
        "conclusions": "We have shown in this present article that the viability of a dark energy model can be cheked in a better way by demanding that $z_{eq}$ should be $\\sim 10^{3}$. Introducing a variable equation of state for matter we have investigated the nature of expansion of the Universe in two variable dark energy models. The variable equation of state for matter ensures that the early epoch of the Universe is dominated by radiation. With the expansion of the Universe the pressure falls as $a^{-m}$ $(m>0)$ to produce an almost pressure-free late-Universe. The equation of state for dark energy is essentially chosen to be $\\alpha_{DE}=p_{DE}/\\rho_{DE}=-1$ while the $\\rho_{DE}$ is in general a function of the scale factor $a$. This in other way amounts to the same as the presence of a dynamical $\\Lambda$ in the field equations. It makes the conservation equation notably different. One has to allow interaction among the matter part and dark energy. The functional form of $\\Lambda$ is then chosen to produce late-time acceleration in the expansion of the Universe. A $\\Lambda \\propto \\rho$ model fails to meet observational constraints. A choice of $\\Lambda \\propto \\rho a^{3}$ is found to be a better alternative. Thus we conclude that meeting the additional $z_{eq}\\sim 10^{3}$ requirement may be an essential feature of reconstructing dark energy models."
    },
    "0808/0808.2742_arXiv.txt": {
        "abstract": "We report the discovery of the largest giant radio galaxy, J1420-0545: a FR type II radio source with an angular size of 17.4' identified with an optical galaxy at $z$=0.3067. Thus, the projected linear size of the radio structure is 4.69 Mpc (if we assume that $H_{0}$=71 km\\,s$^{-1}$Mpc$^{-1}$, $\\Omega_{\\rm m}$=0.27, and $\\Omega_{\\Lambda}$=0.73). This makes it larger than 3C236, which is the largest double radio source known to date. New radio observations with the 100 m Effelsberg telescope and the Giant Metrewave Radio Telescope, as well as optical identification with a host galaxy and its optical spectroscopy with the William Herschel Telescope are reported. The spectrum of J1420$-$0545 is typical of elliptical galaxies in which continuum emission with the characteristic 4000\\,\\AA\\hspace{1mm}discontinuity and the $H$ and $K$ absorption lines are dominated by evolved stars. The dynamical age of the source, its jets' power, the energy density, and the equipartition magnetic field are calculated and compared with the corresponding parameters of other giant and normal-sized radio galaxies from a comparison sample. The source is characterized by the exceptionally low density of the surrounding IGM and an unexpectedly high expansion speed of the source along the jet axis. All of these may suggest a large inhomogeneity of the IGM. ",
        "introduction": "The existence of bright hot spots at the edges of most of FR type II radio sources indicates that the jets ejected from a \"central engine\" in the AGN encounter resistance to their propagation. These supersonically advancing hotspots are likely to be confined by the ram pressure of the external environment, which on the largest scales is the intergalactic medium (IGM). The shocked jet material and shocked IGM form a cocoon surrounding the pair of jets. In a number of FR type II radio sources, such a cocoon is observed as a low-brightness bridge between the radio lobes \\citep{leahy}. While the lengthening of the cocoon is governed by the balance between the jet's thrust and the ram pressure of the IGM, the width of the cocoon is determined by its mean internal pressure, which is on the order of the kinetic energy delivered by the jets divided by the cocoon's volume. As both of these factors are functions of time, it is possible that the lateral expansion of the cocoon could be faster than the rate at which the jets delivered energy. Thus, the cocoon's pressure would decrease until it reached an equilibrium with the IGM pressure. However, in a number of papers (e.g. \\citealt{begel}; \\citealt{daly}; \\citealt{wdw}) there are arguments that the cocoons in many observed radio sources are still overpressured with respect to the IGM. Only the very tenuous material in the bridges of the oldest and largest sources is expected to have attained such an equilibrium state, in which the radiating particles and the magnetic field are balanced, and the pressure of the relativistic plasma equals the pressure of the gaseous environment (cf. \\citealt{nath}). Consequently, a determination of the physical conditions in these diffuse bridges of large-size double radio sources (as they are far away from the ram pressure-confined heads of the source) gives us an independent tool with which to probe the pressure of the IGM.  The above approach was applied by us \\citep{mkj} to a sample of giant radio galaxy (GRG) candidates located in the southern sky hemisphere with the aim of studying properties of the cosmological evolution of the IGM. The radio galaxy discussed in this paper is a member of that sample. The source J1420$-$0545 was discerned in the VLA 1.4 GHz surveys FIRST \\citep{becker} and NVSS \\citep{condon} as a 17.4' large FR type II radio structure consisting of two extended lobes with a total flux density of only 87 mJy and a central compact core. The core is unresolved with the 5\\arcsec$\\times$5\\arcsec\\hspace{1mm} beam in the FIRST survey and has a flux density of 2.7 mJy. The structure is highly collinear and symmetric; the ratio of separation between the core and the brightest regions in the lobes is 1.08, and the misalignment angle is 1.3$\\degr$. The FIRST map reveals also compact hot spots in both lobes, which implies strong shocks and {\\sl in situ} reacceleration of relativistic particles, finally inflating the lobes. In order to check whether the observed lobes are or are not possibly connected by a bridge undetected in the NVSS because of the lack of short baselines, we made supplementary single-dish 21 cm observations using the Effelsberg 100 m radiotelescope, \\footnote{The Effelsberg 100m telescope is operated by the Max-Planck-Institut f\\\"{u}r Radioastronomie (Bonn, Germany)} as well as 50 cm observations with the Giant Metrewave Radio Telescope (GMRT).\\footnote{The GMRT is a national facility operated by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research (Pune, India)} These observations, the final 1.4 GHz map of the investigated radio galaxy obtained from a combination of the NVSS and Effelsberg data, and the 619 MHz GMRT map are presented in \\S~2. The radio core position coincides perfectly with that of a galaxy with $R\\sim$19.65 mag and $B\\sim$21.82 mag at R.A.(J2000.0) = 14$^{\\rm h}$20$^{\\rm m}$23.80$^{\\rm s}$ and decl.(J2000.0) = $-$05\\degr45\\arcmin28.8\\arcsec, according to the magnitude calibration and astrometric position in the Digitized Sky Survey (DSS) data base. This DSS red magnitude suggests a rather distant galaxy whose photometric redshift estimate would be between 0.46 and 0.42, according to the Hubble diagrams for GRGs published by \\citet{schoen} and \\citet{lara}, respectively. However, the standard deviations in the apparent-magnitude distributions in those diagrams for a given redshift are about 1 mag. Therefore, supposing that the absolute magnitude $M_{\\rm R}$ of galaxies identified with very faint GRGs can be 1 $\\sigma$ lower than the mean value, we have estimated a value of $z\\approx 0.32$ for J1420$-$0545 (cf. \\citealt{mkj}). The independent photometry and optical spectroscopy of the host galaxy that confirm the above estimate are given in \\S~3, while some physical parameters and the dynamical age of the radio structure are estimated and discussed in \\S~4. ",
        "conclusions": ""
    },
    "0808/0808.3744_arXiv.txt": {
        "abstract": "\\noindent{\\it The high-lights of ground-based very-high-energy (VHE, $E>100$~GeV) gamma-ray astronomy are reviewed. The summary covers both Galactic and extra-galactic sources. Implications for our understanding of the non-thermal Universe are discussed. Identified VHE sources include various types of supernova remnants (shell-type, mixed morphology, composite) including pulsar wind nebulae, and X-ray binary systems. A diverse population of VHE-emitting Galactic sources include regions of active star formation (young stellar associations), and massive molecular clouds. Different types of active galactic nuclei have been found to emit VHE gamma-rays: besides predominantly Blazar-type objects, a radio-galaxy and a flat-spectrum radio-quasar have been discovered. Finally, many (presumably Galactic) sources have no convincing counterpart and remain at this point unidentified. A total of at least 70 sources are currently known. The next generation of ground based gamma-ray instruments aims to cover the entire accessible energy range from as low as $\\approx 10$~GeV up to $10^5$ GeV and to improve the sensitivity by an order of magnitude in comparison with current instruments. } ",
        "introduction": "The highest energy photons known are produced in astrophysical processes involving even more energetic particles presumably accelerated through stochastic acceleration mechanisms as suggested initially by \\citet{1949PhRv...75.1169F}. Therefore, observations of VHE photons provide a direct view of the astrophysical accelerators of charged particles and allow to identify the individual sources of cosmic rays: VHE photons open our view to the ``accelerator sky''. The origin of the Galactic population of charged cosmic rays has remained since its discovery by V. Hess in 1912 until today a long-standing question of astro- and particle physics. A widely favored model on the origin of cosmic rays assumes that diffusive shock acceleration takes place in the expanding blast waves of supernova remnants converting $10-20~\\%$ of the kinetic energy into cosmic rays. Under this assumption, Galactic supernova remnants provide sufficient power ($\\approx 10^{41}$~ergs~s$^{-1}$) in order to balance the escape losses of Galactic cosmic rays as well as produce a power-law type distribution of particle energy that closely resembles the cosmic ray spectrum measured locally \\citep[see e.g.][]{1964ocr..book.....G,2006astro.ph..7109H}.\\\\ Observations of VHE-gamma-rays, mainly by imaging air Cherenkov telescopes in the last decade, have surpassed the anticipated detection of a few supernova remnants and have established a rich and diverse collection of VHE sources. Specifically,  the currently active generation of imaging air Cherenkov telescopes (H.E.S.S.\\footnote{http://www.mpi-hd.mpg.de/hfm/HESS/HESS.html}, MAGIC\\footnote{http://wwwmagic.mppmu.mpg.de/}, and VERITAS\\footnote{http://veritas.sao.arizona.edu/}) have fulfilled and by far exceeded the expectations that were based upon the pioneering previous generation of experiments: the results obtained in the last years have shown that VHE emission is common to a variety of different source types - not only shell-type SNRs. \\\\ The current view of the source distribution in the VHE sky is shown in Figure~\\ref{fig:sky} where all known sources  are displayed (status of September 2008). Note, the sensitivity achieved  varies greatly across the sky. The best sensitivity is reached in the inner Galactic disk where a dedicated survey with the air Cherenkov telescopes of the H.E.S.S. experiment has been performed (see section~\\ref{subsect:imaging}).\\\\ The most remarkable feature of the gamma-ray sky is not evident in this picture: each source shown is an accelerator of particles up to and beyond TeV energies (``accelerator sky''). A similar conclusion can not be drawn when looking at source populations detected in other wavelength bands. This makes the VHE band a sensitive window to detect non-thermal particle accelerators. Future neutrino telescopes will almost certainly have the potential to detect some of the brighter VHE sources (and discover new sources that are optically thick for VHE gamma-rays). The detection of high energy neutrinos is experimentally a challenging task and requires tremendous efforts (see e.g. the Ice-CUBE neutrino telescope in the Antarctic ice). However, the observation of a neutrino source will ultimately demonstrate the presence of accelerated nuclei - largely model-independent. Different to the clear observational signature in the neutrino channel, VHE gamma-rays are sensitive to accelerated electrons (through inverse Compton scattering) as well as to accelerated nuclei (through neutral meson production and decay).  The scope of this article is limited to observational highlights obtained with ground based instruments and does not aim at presenting a complete review of the field of VHE astrophysics. The paper is structured in the following way: In  section~\\ref{section:obstech}, the experimental techniques of ground-based instruments (both imaging and non-imaging) are described before continuing with the census of today's VHE sky in section \\ref{section:census}. Section~\\ref{sec:secphy} provides a a short overview on VHE gamma-rays as probes of the interstellar medium including Lorenz-invariance violating effects related to structure of space-time at Planck-scales. Finally, in section~\\ref{section:future} this review is concluded with some comments on the future of the field.  \\begin{figure} \\begin{center} \\includegraphics[angle=90,width=1.00\\linewidth]{vhesky.png} \\end{center} \\caption{\\label{fig:sky} All VHE sources as known in September 2008 are shown as circles in a Hammer-Aitoff projection of a Galactic coordinate system. The Galactic  center is in the center of the picture. } \\end{figure} ",
        "conclusions": "The observational field of VHE gamma-ray astrophysics has been successfully driven in the last years by ground based imaging air Cherenkov telescopes.  The ``high energy frontier'' of Astrophysics will be expanded by the results expected from the Fermi satellite.  The all-sky sensitivity of the Fermi Large-Area-Telescope (LAT)  will most likely lead to many interesting new discoveries and surprises. The inter-play of the Fermi-LAT detections from space and follow-up observations from ground will provide mutual benefits for the communities as well as  certainly answer many  questions as well as stimulate new ones.  \\\\ While space-based observations of gamma-rays will be for a long time  limited to the results obtained with Fermi (until maybe pair conversion telescopes will be deployed on the surface of the moon), ground based observatories will continue to be developed and constructed during the next decade(s). To my knowledge, besides Fermi, no further space-based gamma-ray telescope, operating in the GeV energy range, is planned. Consequently, ground based gamma-ray detection techniques will be the only available experimental approach to observe high energy gamma-rays after the termination of the Fermi mission.\\\\ The extensions of H.E.S.S. and MAGIC into phase II until 2009 will lower the energy threshold and improve existing sensitivity moderately.  The next generation of ground based installations will become fully operational at the end of the next decade as envisaged in the AGIS\\footnote{Advanced gamma ray imaging system}\\citep{2007arXiv0709.0704K} and CTA\\footnote{Cherenkov telescope array} \\citep{2007AN....328..600H} projects. These installations aim at an improvement in sensitivity by a factor of 10 and a widened reach in energy. At that time, the new ground based instruments including larger ($>$km$^2$) installations like TenTen \\citep{2008NIMPA.588...48R} will explore new energy windows above $\\approx 10$~GeV as well as above 10~TeV.  \\\\ However, there still remains a ``blind spot'':  the transient sky above 10~GeV will neither be explored very well with Fermi (limited photon rate) nor with ground based instruments of the current generation (limited field of view and energy threshold). New installations like HAWC \\citep{2005AIPC..745..234S} will improve the situation in the energy range above a few 100~GeV. The lesson from the detection of fast transient events as seen from Cyg~X-1 and PKS~2155-304 however is that we are very likely missing the fast variability of XRBs and AGN. Proper coverage of these objects requires instruments with a large field of view and  an energy threshold well below 100~GeV.  \\label{section:future}   {\\small"
    },
    "0808/0808.3602_arXiv.txt": {
        "abstract": "The young system RX~J0529.3$+$1210 was initially identified as a single$-$lined spectroscopic binary. Using high$-$resolution infrared spectra, acquired with NIRSPEC on Keck II, we measured radial velocities for the secondary. The method of using the infrared regime to convert single$-$lined spectra into double$-$lined spectra, and derive the mass ratio for the binary system, has been successfully used for a number of young, low-mass binaries. For RX~J0529.3$+$1210, a long-period (462 days) and highly eccentric (0.88) binary system, we determine the mass ratio to be 0.78 $\\pm$ 0.05 using the infrared double-lined velocity data alone, and 0.73 $\\pm$ 0.23 combining visible light and infrared data in a full orbital solution. The large uncertainty in the latter is the result of the sparse sampling in the infrared and the high eccentricity: the stars do not have a large velocity separation during most of their $\\sim$1.3 year orbit.  A mass ratio close to unity, consistent with the high end of the 1$\\sigma$ uncertainty for this mass ratio value, is inconsistent with the lack of a visible light detection of the secondary component.  We outline several scenarios for a color difference in the two stars, such as one heavily spotted component, higher order multiplicity, or a unique evolutionary stage, favoring detection of only the primary star in visible light, even in a mass ratio $\\sim$1 system.  However, the evidence points to a lower ratio.  Although RX~J0529.3$+$1210 exhibits no excess at near-infrared wavelengths, a small 24~$\\mu$m excess is detected, consistent with circumbinary dust.  The properties of this binary and its membership in $\\lambda$ Ori versus a new nearby stellar moving group at $\\sim$90~pc are discussed. We speculate on the origin of this unusual system and on the impact of such high eccentricity, the largest observed in a pre$-$main sequence double$-$lined system to date, on the potential for planet formation. ",
        "introduction": "Understanding the process of star formation requires reliable observations of fundamental stellar properties so that theoretical models can be tested.  The copious population of young binary systems, however, complicates the problem.  The binary fraction can be high \\citep{2000prpl.conf..703M}, and may depend on the density of the star forming region \\citep{2003ApJ...583..358B} and spectral type of the primary \\citep{2006ApJ...640L..63L}, underscoring the importance of their study. Fortuitously, binary stars serve two purposes.  One, characterization of their frequency, separation distribution, and mass ratio distribution for a given star forming region (SFR) provides clues to the broad star forming properties (angular momentum, density, turbulence, etc.) of the parent molecular cloud.  Two, the special class of very small separation spectroscopic binaries with periods sufficiently short to enable the measurement of the individual stellar velocities, and thus the system's mass ratio, are potential targets for the dynamical determination of individual component stellar masses (e.g., Steffen et al. 2001; Prato et al. 2002a; Boden et al. 2005; Stassun et al. 2007). Knowledge of absolute masses and observable properties, such as effective temperature and luminosity, plays a key role in the improvement of pre-main sequence (PMS) evolutionary models \\citep{2001ApJ...553..299P}.  Once spectroscopic binaries have been identified, it is necessary to characterize their properties.  It is often the case with low mass-ratio systems that they are identified as single$-$lined spectroscopic binaries when observed in visible light (Mazeh et al. 2002), in which case the large difference in flux between the primary and secondary at short wavelengths prevents detection of the secondary component and thus the measurement of the mass ratio for the system.  In the Raleigh-Jeans regime, however, flux scales much less steeply as a function of mass. Thus, by observing single-lined spectroscopic binaries with infrared (IR) spectroscopy we are able to improve our chances of detecting the lower-mass secondary not seen in visible light. This technique was initially outlined in Prato et al. (2002b) and Mazeh et al. (2002, 2003). Our primary motivation for observing RX~J0529.3$+$1210 was to use this IR approach to determine the mass ratio of the system by converting it into a double-lined spectroscopic binary. Given that the secondary had not been detected in visible light, we anticipated a relatively low mass ratio for RX~J0529.3$+$1210.  With the advent of the {\\it R\\\"ontgensatellit (ROSAT)} all-sky survey, a number of researchers began a search for X-ray sources with optical counterparts of brightness consistent with membership in nearby SFRs, motivated in part by the goal of detecting the ``post-T Tauri'' population postulated by Herbig (1978).  In follow up observations of X-ray sources near Taurus, RX~J0529.3$+$1210 was identified variously as a PMS star (Neuhauser et al. 1997; Magazz\\`u et al. 1997) and a post-T Tauri star (Magazz\\`u et al. 1999); high-resolution visible light spectroscopy revealed its spectroscopic binary nature \\citep{neu97}.  \\citet{tor02} undertook a seven year campaign to characterize the orbits of all spectroscopic binaries identified in the X-ray sample of \\citet{neu95} and \\citet{neu97}, including RX~J0529.3$+$1210.  Table 1 summarizes the general properties of this system.  The $\\sim$3 dozen high-resolution visible light spectra taken of this binary suggested the presence of a secondary star; however, a conclusive identification was not possible.  A single-lined orbital solution was determined, although, probably owing to the extremely high eccentricity of the system the uncertainties are relatively large.  This paper describes the results of using high resolution IR spectroscopy to observe RX~J0529.3$+$1210 and determine the component radial velocities. RX~J0529.3$+$1210 is the most eccentric pre$-$main sequence spectroscopic binary known to date. This provides a unique context in which to speculate on the formation of the system and on the impact of the stellar dynamics on potential planet formation. In \\S2 we briefly describe our observations and data reduction. Our analysis and results appear in \\S3. Section 4 provides a discussion, and \\S5 summarizes our findings. ",
        "conclusions": "\\subsection{Where and How Old is RX~J0529.3$+$1210?}  The projected location of RX~J0529.3$+$1210 is associated with a high-density area in the CO maps of \\citet{dol01} for the $\\lambda$ Ori region.  However, \\citet{dol01} find a mean radial velocity for the strong lithium sources identified in $\\lambda$ Ori of 24.5~km~s$^{-1}$, with a dispersion of only 2.3~km~s$^{-1}$. Our center-of-mass velocity for RX~J0529.3$+$1210 is 18.38~km~s$^{-1}\\pm0.30$, indistinguishable from that found by \\citet{tor02}.  Thus, on the basis of radial velocities alone, it seems unlikely that the system is associated with $\\lambda$ Ori since RX~J0529.3$+$1210's center-of-mass velocity is inconsistent with that of $\\lambda$ Ori at the 3$\\sigma$ level.  Using the 2MASS JHK magnitudes of RX~J0529.3$+$1210, an effective temperature (T$_{eff}$) for a K7/M0 of 3900~K \\citep{2003ApJ...593.1093L,bro06}, and the nominal distance to $\\lambda$ Ori (400~pc; Barrado y Navascu\\'es 2005) we have estimated the luminosity and placed the system on an H-R diagram.  We find L $=$ 3.44 $\\pm$ 0.01 L$_{\\odot}$, yielding an age of $\\sim$0.1 Myr using the PMS tracks of \\citet{pal99}.  Accounting for a companion star with luminosity equal to that of the primary results in an age of $\\sim$0.5~Myr.  Using the tracks of \\citet{bar98}, or absolute K magnitude and the tracks of \\citet{2000A&A...358..593S}, we obtain similar age estimates.  \\citet{bar05} estimates the age of $\\lambda$ Ori to be between 3 and 10 Myr, \\citet{bar07} find an age of 5~Myr, and \\citet{dol01} describe a spread in ages from 1 to 10 Myr. A $<<$1~Myr object is expected to be at least somewhat embedded in its natal cloud and associated with circumstellar material, yet near-IR colors from 2MASS magnitudes indicate zero extinction, no veiling is detected in the spectra, and the H$\\alpha$ emission line equivalent width (Table 1) is only 2\\AA.  We consider two possible mechanisms for disk dissipation in this system, photoevaporation from nearby hot stars and the orbital dynamics of RX~J0529.3$+$1210 itself.  The projected separation of the B8 star HD 36104 \\citep{dol01} from RX~J0529.3$+$1210 is $\\sim$1~pc if both are assumed to be at a distance of 400~pc. \\citet{dol01} discuss the surprising lack of evidence for accretion disks around the PMS stars in the central $\\lambda$ Ori cluster and suggest that photoevaporation and/or possibly a supernova event 1$-$2 Myr ago played a role in the dispersion of disks in the local young low-mass stellar population.  RX~J0529.3$+$1210, located just north of the central cluster in the Barnard 30 dark cloud, could have experienced photoevaporation from the nearby B star at a young age, obliterating circumstellar material, although this is probably unlikely given the long survival time of proplyds in the Trapezium \\citep{1998AJ....115..263O}.  Given the extremely high eccentricity of this binary, however, the action of the companion star may have precluded formation of, or dissipated any, circumstellar material.  With an assumed primary star mass of 0.75 M$_{\\odot}$ and the mass function from Torres et al. (2002), we find a minimum secondary star mass of 0.40 M$_{\\odot}$, and thus a minimum total mass of 1.15 M$_{\\odot}$.  In conjunction with the 461.89 day period, this yields a periastron separation of 0.15~AU and an apastron separation of 2.30~AU. It is unknown if orbital evolution may have occurred, or even whether it is possible that this unusual system formed by capture, but if RX~J0529.3$+$1210 is located in $\\lambda$ Ori with an age of $<$1~Myr, then it seems likely that the system formed in a similar configuration to its present one.  Independent of the potentially harsh $\\lambda$ Ori environment and the orbital dynamics of RX~J0529.3$+$1210, the system manifests spectra of similar surface gravity to dwarf stars, as well as a relatively small lithium equivalent width, in comparison to classical T Tauri stars \\citep{1989AJ.....98.1444S,2005fost.book.....S}.  The lithium and H$\\alpha$ equivalent widths of \\citet{dol01} for $\\lambda$ Ori average greater than twice what is found for RX~J0529.3$+$1210 and are consistent with the classical and weak-lined T-Tauri limits defined by \\citet{1998AJ....115..351M}.   According to these limits, RX~J0529.3$+$1210 would be a post-T Tauri star and have one of the lowest lithium and H$\\alpha$ equivalent widths in the $\\lambda$ Ori region.  These characteristics are inconsistent with an age of $<$1~Myr.  Alternatively, RX~J0529.3$+$1210 could be an older, closer object. \\citet{mam07} describes a new candidate moving group, 32 Ori, consisting of a small cluster of X-ray bright, late type stars.  The $\\sim$10 young stars identified in the group are located around 5$^h$ 20$^m$ to 5$^h$ 30$^m$ and $+6\\deg$, at the proposed distance of $\\sim$90~pc.  RX~J0529.3$+$1210 is about 9~pc from the central clump of objects in 32 Ori and is one of the members used by Mamajek to define the common proper motion group (E. Mamajek 2008, private communication).  Given the proper motion of RX~J0529.3$+$1210, pmRA $=$  4.1 $\\pm$ 5.8 mas/yr and pmDec $=$ $-$30.7 $\\pm$ 5.8 mas/yr, compared to the proper motions of the stars $\\lambda$ Ori (pmRA $=$ 0.8 $\\pm$ 1.5, pmDec $=$ $-$2.3 $\\pm$ 1.5) and 32 Ori (pmRA $=$ 6.57 $\\pm$ 1.15, pmDec $=$ $-$32.45 $\\pm$ 0.48), evidence in support of membership in the 32~Ori group is compelling \\citep{1997A&A...323L..49P, 2004AJ....127.3043Z}.  Again combining the 2MASS JHK magnitudes of RX~J0529.3$+$1210, an effective temperature of 3900~K, and a distance now of 90~pc, we find L $=$ 0.17 $\\pm$ 0.02 L$_{\\odot}$, giving an age for RX~J0529.3$+$1210 on the tracks of \\citet{pal99} of $\\sim$15$\\pm$5~Myr, consistent with the 25$\\pm$10~Myr age derived by \\citet{mam07} for the candidate group members.  \\citet{mam07} lists a group radial velocity of 18~km~s$^{-1}$, in excellent agreement with our measured center-of-mass velocity (Table 2).  Morales Calderon (2008, in prep) has observed the $\\lambda$ Ori region with the {\\it Spitzer} space telescope and finds a 3~$\\sigma$ detection of RX~J0529.3$+$1210 at 24~$\\mu$m with a 20~\\% excess above a 3900~K photosphere.  For a star with L$=$0.17~L$_{\\odot}$, the equilibrium temperature of a black-body grain with peak emission at 24~$\\mu$m corresponds to a distance of 4.36~AU. A more realistic treatment of the dust grain distribution would necessarily take into account the additional flux from the secondary star, yielding a larger distance from the center-of-mass of the system to the putative dust. However, given the estimated apastron separation of 2.30~AU, a lower limit of 4.36~AU for the dust radius from the system center illustrates the plausibility of a circumbinary debris disk.  Additional {\\it Spitzer} observations at shorter wavelengths have been taken and should reveal more information regarding the extent and location of the dust (Mamajek 2008, in preparation).  The lack of J$-$H, H$-$K, and K$-$L excesses is hardly surprising.  For black-body grains with a peak wavelength in the L band, the corresponding disk temperature occurs at a distance of $\\sim$0.1~AU, nearly coincident with the binary periastron.  Stable dust in this system is most likely to be located in a circumbinary distribution.  Furthermore, if the closer distance and therefore older age for this system implied by membership in 32 Ori are correct, then the presence of an evolved debris disk would not be unusual \\citep[e.g.,][]{tri08}.  Finally, the fact that the RX~J0529.3$+$1210 secondary was possibly detected in our AO images strongly supports the $\\sim$90~pc distance.  For the projected separation of 0.018\", determined using PSF fitting, and a distance of 90~pc, the corresponding distance in AU is 1.6$\\pm$0.54, not too different from our estimated apastron of 2.30~AU.  For a distance of 400~pc the separation would be 7.2$\\pm$2.4 AU. Based on the currently available data, it therefore appears that RX~J0529.3$+$1210 is associated with the new 32 Ori group, not with $\\lambda$ Ori, and has an age of $\\sim$15~Myr.   \\subsection{The Mass and Flux Ratios of RX~J0529.3$+$1210}  \\citet{tor02} note that the appearance of the velocity correlation function for RX~J0529.3$+$1210 suggests the presence of at least one other star in the system. For a primary star mass of 0.75~M$_{\\odot}$, the mass function implies a minimum secondary star mass of 0.40~M$_{\\odot}$ and a minimum mass ratio of 0.53, consistent to within 1$\\sigma$ of the mass ratio found from the full orbital solution for the system, 0.73$\\pm$0.23.  The approximate H-band flux ratio measured from the cross-correlation is 0.6$\\pm$0.1. The K-band flux ratio found in the best PSF fit to the January, 2008 AO images is 0.66$\\pm$0.18. For a M2.5 $+$ K7 pair, the spectral types which provided the best correlation (\\S 3), models of \\citet{pal99} and \\citet{bar98} imply an H-band flux ratio of 0.4$\\pm$0.1 based on components with a primary T$_{eff}=$3900~K and a secondary T$_{eff}=$3400~K \\citep{2003ApJ...593.1093L,joh66}.  Taken together, these properties suggest a relatively red color for the secondary star, hindering detection in visible light.  The IR primary and secondary velocities, presented in the Wilson plot (Figure 2), imply a range of mass ratios, from 0.73$-$0.83, well within the range we obtain by combining the IR and visible light data in a full orbital solution, $\\sim$0.50$-$0.93. Because the full solution includes phase information, it is generally more reliable, although for RX~J0529.3$+$1210 the sparse phase coverage during the epochs of largest velocity separation yields large uncertainties in the orbital solution.  Clearly more data are merited. For the case of a mass ratio relatively close to unity in this system (i.e. 0.8$-$1.0), it is necessary to explain the lack of a visible light detection of the secondary star. In the following paragraphs we outline several scenarios that would account for this.  A relatively red secondary and yet a mass ratio close to unity is possible if a third component is present, as first suggested by \\citet{tor02}. If the secondary star is actually a binary pair, with a very small separation and similar component masses, then the spectral types could be consistent with the best-fitting TODCOR templates. This inner binary pair would necessarily be in an orbit in a plane relatively perpendicular to our line of sight, such that the large velocity changes would not be obvious. This scenario would imply a near-unity mass ratio yet a smaller near-IR flux ratio, and a red enough secondary system to evade visible light detection.  The drawback is that it requires a very specific geometry.   Such a companion might also account for the 24~$\\mu$m excess.  A heavily spotted secondary star could mimic a lower temperature source, resulting in a best fitting secondary M2.5 template, a mass ratio close to 1.0, and a flux ratio $<$1 in the near-IR.  Again, the reddening effect of the cooler temperature in the spot covered areas would increase the difficulty of visible light detection.  For two T$_{eff}=$3900~K stars, if one has 50\\% of its surface covered with a 2900~K spot \\citep{bou89}, the secondary/primary bolometric flux ratio would be $\\sim$0.7.  For objects with ages between 10 and 30 Myr, some models \\citep{bar98,pal99} of PMS evolution indicate that targets in the mass range of the RX~J0529.3$+$1210 primary star, 0.6$-$0.8~M$_{\\odot}$, are located near the transition between the convective Hyashi track and the radiative Henyey track.  As a result, models show a sharp turn in the mass tracks towards higher temperatures. A similar but slightly lower mass secondary star that has not yet turned this corner off of the Hyashi track, and indeed, objects with masses lower than $\\sim$0.5~M$_{\\odot}$ never will, may have a similar mass as the primary star but a much lower temperature.  These models for the system behavior are highly speculative.  The absence of an unambiguous secondary star detection in visible light suggests that the mass ratio is {\\it not} close to unity. A value in the 0.7 to 0.8 range is most likely based on data to date.  \\subsection{Binary Formation Scenarios and the Potential for Planets}  The formation of such a highly eccentric multiple is challenging to understand as a primordial event given current theories of star formation \\citep{2005fost.book.....S}.  Figure 4 shows that pre$-$main sequence spectroscopic binaries are observed over a wide range of eccentricities.  Although RX~J0529.3$+$1210 represents the maximum of that range, it does not stand out particularly from the eccentricity distribution.  Thus, some form of dynamical evolution may have taken place in this system, but, if so, it is not distinguished by an outlying eccentricity. Possibilities for dynamical evolution range from an improbable capture event to disk excitation of stellar eccentricity \\citep{mat92}. If indeed this system is $>$10~Myr old, only fossil evidence, such as the detected 24~$\\mu$m excess, might remain of the primordial, presumably massive disk(s) responsible.  In a variation of the interactions proposed by \\citet{rei00}, a three body dynamical encounter between an object from outside of the binary and a companion short period binary could have stimulated the eccentricity of the system and tightened the orbit of a close (secondary) pair (\\S 4.2).  The eventual formation of a circumbinary (or circumtriple) debris disk could have been stimulated by dynamical disruption of the disks in the system in any one of these scenarios.  The formation of planets in binary systems has been observationally supported in recent years and is of primary significance since a majority of stars have companions \\citep{egg07}. \\citet{cun07} have determined the strict criteria for classifying stable and unstable planetary orbits in binary systems, however, their analysis is restricted to circular orbits. Many of the extra-solar planets observed in binaries are in wide separation systems, where the mass ratio for the stellar components is near unity and the eccentricity of the system is low.  \\citet{qui07} have modeled the formation of terrestrial planets around individual stars in binaries and find that, for periastron distances of $<$5~AU, such planet formation is restricted.  Considering that the periastron distance in the RX~J0529.3$+$1210 system is only 0.15~AU, we conclude that the probability of even a low-mass circumstellar planet in RX~J0529.3$+$1210 is remote.  The trend determined by \\citet{qui06} towards the formation of fewer circumbinary terrestrial planets in systems with apastron distances of $>$0.2~AU and non-zero eccentricities also bodes for poor circumbinary planetary stability around the stars in RX~J0529.3$+$1210.  \\subsection{Improving our Understanding of the RX~J0529.3$+$1210 System: Future Work}  Clearly it is imperative to improve the orbital solution for RX~J0529.3$+$1210, particularly the measurement of the mass ratio.  Visible light and especially infrared data are critical to obtain, preferably at high precisions. RX~J0529.3$+$1210 passes through its next maximum velocity separation in late December, 2009.  Densely sampled high signal-to-noise spectra taken during these epochs will yield a greatly improved precision for the orbital solution.  If RX~J0529.3$+$1210 is located at a distance of only 90~pc, an apastron separation of 2.30~AU implies a maximum projected separation on the sky of $\\sim$0.03$''$.  Much of the orbit of this system lies within reach of the diffraction limit of the 85~m baseline of the Keck Interferometer, 0.005$''$. Although the K$-$band magnitude of RX~J0529.3$+$1210, 9.2, is relatively faint for observations with this facility, planned improvements to the system may eventually enable resolved observations of this unusual binary, permitting the determination of individual component masses.  The 90~pc distance may also be confirmed by measuring the parallax, which is slightly greater than 10 mas at this distance, and is at the achievable limit of current instrumentation."
    },
    "0808/0808.2815_arXiv.txt": {
        "abstract": "We present the first detection of CS in the Antennae galaxies towards the NGC~4038 nucleus, as well as the first detections of two high-J (5-4 and 7-6) CS lines in the center of M~82. The CS(7-6) line in M~82 shows a profile that is surprisingly different to those of other low-J CS transitions we observed. This implies the presence of a separate, denser and warmer molecular gas component. The derived physical properties and the likely location of the CS(7-6) emission suggests an association with the supershell in the centre of M~82. ",
        "introduction": "\\label{sec:intro}   The molecule CS is a good tracer of dense gas (n(H$_2$)$\\geq$ 10$^{5}$-10$^{7}$ cm$^{-3}$) in massive star-forming regions in our own Galaxy (\\citealt{Plum92, Plum97, Andr07}) and in nearby galaxies such as the Magellanic Clouds \\citep{Niko07}, M~51 and Maffei~2 \\citep{Pagl95}, NGC~253 \\citep{Maue89a, Maue89b, Mart05}, IC~342 and M~82 \\citep{Henk85, Maue89a, Maue89b}. Multi-line studies of CS are crucial for the determination of the average gas densities in galaxies since CS transitions have excitation thresholds ranging from 10$^{4}$-10$^{5}$ cm$^{-3}$ for the CS(2-1) line \\citep{Bron96} up to $\\sim 2\\times 10^{7}$ cm$^{-3}$ for the CS(7-6) transition \\citep{Plum92}. In this letter, we report in the center of M~82 the first detections of the CS(5-4) and CS(7-6) transitions and the detection of the CS(5-4) line towards the Antennae galaxies (NGC~4038). We have also re-observed lower-J CS lines in M~82. These two well-known sources were chosen because interferometric submillimeter/millimeter maps previously obtained have shown high concentrations of molecular gas. High resolution $^{12}$CO(1-0) maps indicate that the nearby (D = 13.8Mpc, see \\citealt{Savi04}) Antennae interacting galaxies are likely sites of rapid high mass star formation \\citep{Wils00}. M~82 represents an excellent example of a nearby (D=3.25 Mpc, see \\citealt{Dumk01}) starburst galaxy. Low-J molecular line studies (e.g. \\citealt{Fuen06, Seaq06, Mart06}) have determined its average gas physical parameters. However, the molecular emission from the M~82 nucleus appears to come from multiple gas components. In particular the detection of abundant CH$_{3}$OH \\citep{Mart06} and HCN \\citep{Brou93} indicate the presence of high density gas in the nucleus. Our detection of the CS(7-6) line not only clearly confirms the presence of very high density gas but its analysis (see Sect.~\\ref{sec:discu}) suggests that it is located in the expanding superbubble likely to be associated with supernoave remnants \\citep{Weis99, Will99, Yao06}. ",
        "conclusions": "\\label{sec:discu}   The angular resolution of the CS(7-6) (beam size of 14\") does not allow us to resolve spatially the emitting region. However we can estimate its location from the radial velocity and the velocity width of this line by using the kinematic information provided by interferometric maps of other molecules. \\citet{Brou93} presented a 2\" resolution HCN(1-0) map of the south west part of M~82. From the channel map of the HCN(1-0) emission centered around 213.3 kms$^{-1}$($\\pm 40$kms$^{-1}$), it appears that the CS(7-6) emission may arise from the HCN(1-0) clump located at $\\alpha_{\\rm J2000}=09^h55^m52.2^s$ and $\\delta_{\\rm J2000}=69^\\circ40'46.1''$, within our beam. \\citet{Garc01, Garc02} presented high resolution interferometric HCO and H$^{13}$CO$^{+}$(1-0) maps of M~82. We re-analyzed the data and found that some of the clumps show narrow ($\\leq$ 60 kms$^{-1}$) emission, similar to that of our CS(7-6) line, confirming the likely location of the CS(7-6) emission at $\\approx$ (-5$''$, -2$''$) from the center. This is the location of the expanding molecular supershell \\citep{Weis99, Will99, Yao06}. The high density ($\\geq 10^{6}$cm$^{-3}$) and the high temperature ($\\approx$ 60-80K) may arise from the interaction between the expanding supershell and the ambient gas in M~82.  \\begin{table*} \\caption{Observational parameters.}\\label{tab:obs} \\renewcommand{\\footnoterule}{} \\hspace*{-2cm} \\begin{tabular}{l c c c c c c c c c } \\hline Source & Line & $\\nu$ & Tsys & beam & $\\int$(T$_{mb}$ dv)& V$_{LSR}$ & $\\Delta$V$_{LSR}$ & T$_{peak}$ & rms\\\\ & & (GHz) & (K) & size (\") &  (Kkms$^{-1}$) & (Kkms$^{-1}$) & (Kkms$^{-1}$) & (mK) & (mK)\\\\ \\hline M~82 & CS(2-1) & 97.980 & 281 & 25 & 13.3$\\pm$0.3 & 221.9$\\pm$2.7 & 225.7$\\pm$5.4 & 55.5 & 5.8\\\\ & CS(3-2)& 146.969 & 303 & 17 & 11.2$\\pm$0.3 & 219.7$\\pm$2.8 & 211.3$\\pm$6.4 & 50.1 & 4.6\\\\ & CS(4-3)& 195.954 & 1655 & 13 & 11.1$\\pm$1.3 & 161.0$\\pm$9.7 & 165.3$\\pm$23.2 & 63.4 & 25.8\\\\ & CS(5-4)& 244.936 & 437 & 20 & 4.2$\\pm$0.8 & 140.0$\\pm$38.0 & 211.2$\\pm$44.2 & 15.6 & 10.8\\\\ & CS(7-6)& 342.883 & 334 & 14 & 2.2$\\pm$0.2 & 213.8$\\pm$1.5 & 40.1$\\pm$2.8 & 52.5 & 14.1\\\\ \\hline NGC & CS(5-4)& 244.936 & 246 & 20 & 1.7$\\pm$0.1 & 1655.0$\\pm$4.4 & 98.7$\\pm$8.5 & 16.3 & 6.7\\\\ 4038&&&&&&&&&\\\\ \\hline \\end{tabular} \\end{table*}  \\begin{figure} \\includegraphics[height=12cm]{fig1c.jpg} \\caption{Spectra of the CS(2-1), CS(3-2), CS(4-3), CS(5-4) and CS(7-6) lines from top to bottom, respectively, measured towards the center of M~82. The observed position corresponds to the center of M~82 ($\\alpha_{\\rm J2000}=09^h55^m51.9^s$ and $\\delta_{\\rm J2000}=+69^\\circ40'47''$). The CS(2-1) and the CS(4-3) spectra have been smoothed to a common velocity resolution of 6.1kms$^{-1}$ while the CS(3-2), CS(5-4) and CS(7-6) spectra have a velocity resolution of 8.2kms$^{-1}$, $\\approx$ 10kms$^{-1}$ and 6.8kms$^{-1}$, respectively. The velocity scale (x axis) is expressed in kms$^{-1}$ units while the temperature scale (y axis) is expressed in main beam temperature units (T$_{MB}$), converted from the antennae temperature via the main beam efficiencies listed in Sect. ~\\ref{sec:obs}. The solid black lines superimposed onto each spectrum represents the Gaussian fit while the vertical black line marks the systemic M~82 V$_{LSR}$. }\\label{fig:1} \\end{figure} \\begin{figure} \\includegraphics[height=4cm]{fig2.jpg} \\caption{Spectrum of the CS(5-4) line towards NGC~4038. The observed position corresponds to the NGC~4038 nucleus ($\\alpha_{\\rm J2000}=12^h01^m52.8^s$ and  $\\delta_{\\rm J2000}=-18^\\circ52'05''$). The spectrum has been smoothed to a velocity resolution of 2.3 kms$^{-1}$. The solid black line superimposed on the spectrum represents the Gaussian fit while the vertical black line marks the systemic NGC~4038 V$_{LSR}$.}\\label{fig:2} \\end{figure} \\begin{figure} \\includegraphics[height=4cm]{fig3.jpg} \\caption{Rotation diagram derived from the CS lines towards M~82. The black lines represent the linear regression fits for two gas components while the dashed line corresponds to one single gas component. The black solid squares show the data with error bars. These error bars usually represent the main errors in rotational diagrams which actually correspond to those of the integrated intensities (order of 10-20\\%). This population diagram has been corrected for beam dilution effects assuming a source size of 10$''$ for the center of M~82. This value is estimated from the interferometic data presented in \\citet{Brou93} and \\citet{Garc01}.}\\label{fig:3} \\end{figure}   \\vspace*{1cm}   In summary, we have presented the first detection of a high density tracer (CS) in the Antennae galaxies (NGC~4038 nucleus) as well as the first detections of two high-J transitions of CS in the center of M~82. We find that multiple molecular gas components in M~82 are necessary to explain the observed line intensities. In particular, while low-J CS lines seem to arise from relatively low density gas ($\\sim 10^{5}$cm$^{-3}$), the molecular gas traced by the CS(7-6) line must be dense ($\\sim 10^{7}$cm$^{-3}$) and warm ($\\sim$ 70K) and appears to be associated with the expanding supershell in M~82. The high density and temperature may be due to the interaction between the expanding supershell and the ambient gas. Similar multi-line studies for the Antennae galaxies (NGC~4038/39) are necessary in order to determine the origin of its high density molecular gas component traced by the CS(5-4) line. In both sources, the CS column densities are consistent with model predictions of gas undergoing high-mass star formation \\citep{Baye08}.   Similar extended multi-transition multi-molecule studies performed on ultra-luminous infra-red galaxies by \\citet{Grev06b} and \\citet{Baan08} need to be carried out on nearby sources. In fact, \\citet{Mueh07} presented a detailed study of the para-H$_{2}$CO in M~82; however this molecule does not trace the dense gas component. Thus, in order to compare the physical and chemical properties of the very dense (extragalactic) gas between ULIRGs and nearby sources, CS observations (from 2-1 to 7-6) of ULIRGs should be performed."
    },
    "0808/0808.0398_arXiv.txt": {
        "abstract": "{} {The main purpose of this work is to improve the existing knowledge about the most powerful engines in the Universe -- quasars. Although a lot is already known, we still have only a vague idea how these engines work exactly, why they behave as they do, and what the relation is between their evolution and the evolution of their harboring galaxy.} {Methods we used are based on optical spectroscopy of visually bright quasars, many of which have recently been discovered as X-ray sources, but eventually missed in color-selected surveys. The spectra typically cover the 4200--7000 \\AA\\AA\\ region, allowing measurements of the characteristics of the hydrogen lines, the FeII contribution, and other lines of interest.} {We present accurate redshift estimates and Seyfert type classification of the objects. We also show that the contribution of the host galaxy to the optical continuum is non-negligible in many cases, as is the intrinsic AGN absorption. Consequences of not correcting for those factors when estimating different quasar parameters are discussed. We also find some evidence of a non-unity slope in the relation between the internal extinction based on the Balmer decrement and the one on the optical continuum slope, implying, if further confirmed, the intriguing possibility that some absorbing material might actually be located $between$ the continuum source and the broad-line region.} {} ",
        "introduction": "The generally accepted paradigm about the nature of quasars (or active galactic nuclei; AGN) invokes a supermassive black hole, surrounded by an accretion disk, where most of the emitted energy is generated. In the vicinities of this disk, there must be a region, producing the broad emission lines (BLR), traditionally presented as a system of clouds, moving in the gravitational field of the black hole. Farther out from the center should be a region producing narrow lines (NLR), which may be the only visible optical lines in the spectra, if the central region is hidden from the observer behind a thick torus, as the unification models explain the AGN type I/II differences (Antonucci, 1993, for a review). Studying the narrow lines is important, since they are fairly sensitive to the shape of the ionizing continuum, and especially to the unobserved EUV range, which, if restored properly, may help to test the general paradigm or to model the type of accretion disk (e.g. standard thin disk, ADAF, etc.). One of the least clear points in this picture so far appears to be the exact geometry (e.g. flat, spherical, conical), structure (system of clouds, irradiated accretion disk, etc.), and kinematics (inflow, outflow, Keplerian rotation) of the BLR matter. The BLR should play an important role in feeding the quasar; therefore, knowledge of its nature may reveal why some galactic centers are active and others are not, in case a supermassive black hole is likely to be present in both types (Magorrian et al.  1998).  The best way to probe the broad-line region might be through its optical emission-line spectrum. Although a single spectrum is certainly not sufficient to answer all the questions, studying the line profiles, ratios, and other properties in many objects seems to be one way of sheding some light on the problem. A promising new approach to systematizing quasar properties and differentiating between quasars is based on principle component analysis (PCA), where the quasar diversity can be expressed as a linear combination of a few eigenvectors of the correlation matrix, constructed from different measured quantities of a sample of quasars. Many of these measures are based on the optical spectra (Boroson \\& Green 1992; Sulentic et al. 2000), such as the width of the broad lines, equivalent width ratios. It is proposed that the main eigenvector (EV1), i.e. the main driver of the quasar diversity, is the accretion rate ratio (Marziani et al. 2001), one of the defining characteristics of the AGN. On the other hand, another important characteristic of the central engine, the black hole mass, is often calibrated through empirical relations, like the ``width of the broad lines, FWHM -- the monochromatic continuum luminosity, $\\lambda L_{\\lambda5100}\\AA$'', (Kaspi et al.  2005), ``the black hole mass -- bulge luminosity or bulge velocity dispersion'' (Magorrian et al.  1998; Ferrarese \\& Merritt, 2000), etc., since the only direct method known, the reverberation mapping (see Peterson \\& Horne, 2004, for a review), is observationally very expensive. It is clear that a large sample of quasars with relatively good spectra is needed for better understanding the importance of EV1 correlations for the quasar physics and to better calibrate different quasar empirical relations.  Nowadays we see that many bright broad-line quasars have been missed in color-selected samples (e.g. Palomar--Green, Schmidt \\& Green, 1983), mainly because of their reddened colors, due to significant host galaxy contribution and/or intrinsic absorption. Being mostly radio-quiet, many of those objects have only recently been discovered and identified as quasars through X-ray surveys. Studying such objects in detail is especially important for preventing introduction of different biases into our knowledge about the overall AGN population, its properties, and evolution. Furthermore, these partially obscured objects (reddened S1, S1.8, S1.9 types, etc.) may provide an important link between the ``pure'' type I AGN and the obscured population (type II / Seyfert 2), especially taking into account the building evidence that the latter may comprise a significant part of the entire AGN population (Zakamska et al.  2003). This paper should be considered as a part of a larger continuing effort to obtain spectra of relatively good quality for most of the nearby, visually bright quasars, including many discovered in recent years. In a subsequent paper, we will concentrate on the connection between the quasar host galaxies and different AGN characteristics.  The paper is organized as follows. In the next two chapters we describe the observations and present the main results, including individual notes on the objects. Discussion is presented in Sect. 4 and the summary, in Sect. 5.  \\begin{table*} \\caption{Log of observations.} \\label{table:1} \\centering \\begin{tabular}{lcccccccccc} \\noalign{\\smallskip} \\hline\\hline Object & RA J2000 & Dec J2000 & z & $\\sigma$(z) &  A$_{\\rm V}$ & Date & Instr. & Exp. & S/N\\\\ \\hline  2MASXJ005050+3536& 00 50 51 & +35 36 43 & 0.0585 & 0.0004 & 0.14 & 08.09.2004 & Rz 2.0 &   2700 &  36\\\\ MCG+08.17.060    & 09 13 46 & +47 42 00 & 0.0524 & 0.0005 & 0.05 & 19.03.2004 & Rz 2.0 & 2x2400 &  10\\\\ RXS J11401+4115  & 11 40 03 & +41 15 05 & 0.0717 & 0.0002 & 0.06 & 22.06.2005 & Sk 1.3 & 2x1200 &  12\\\\ PG 1211+143      & 12 14 18 & +14 03 13 & 0.0815 & 0.0003 & 0.11 & 23.06.2005 & Sk 1.3 & 2x1200 &  27\\\\ RXS J12308+0115  & 12 30 50 & +01 15 21 & 0.1183 & 0.0005 & 0.06 & 24.06.2005 & Sk 1.3 &   1200 &  25\\\\ RXS J16312+0955  & 16 31 16 & +09 55 58 & 0.0917 & 0.0005 & 0.18 & 17.08.2004 & Sk 1.3 & 2x2700 &  35\\\\ RXS J17233+3630  & 17 23 25 & +36 30 25 & 0.0400 & 0.0003 & 0.16 & 29.06.2005 & Rz 2.0 &   1800 &  25\\\\ IRAS 18423+2201  & 18 44 30 & +22 04 28 & 0.0464 & 0.0004 & 0.69 & 11.08.2005 & Sk 1.3 & 2x2400 &  15\\\\ NPM1G+27.0587    & 18 53 04 & +27 50 28 & 0.0620 & 0.0004 & 0.49 & 10.08.2005 & Sk 1.3 & 2x1800 &  31\\\\ &          &           &        &        &      & 04.05.2005 & Rz 2.0 &   1800 &  29\\\\ RXS J20440+2833  & 20 44 04 & +28 33 09 & 0.0492 & 0.0005 & 1.03 & 28.08.2005 & Sk 1.3 &   2400 &  38\\\\ &          &           &        &        &      & 30.06.2005 & Rz 2.0 &   1800 &  33\\\\ NPM1G$-$05.0589  & 21 03 38 &$-$04 55 40& 0.0575 & 0.0002 & 0.27 & 18.08.2004 & Sk 1.3 & 2x2700 &  31\\\\ RXS J21240$-$0021& 21 24 02 &$-$00 21 58& 0.0617 & 0.0003 & 0.15 & 17.08.2004 & Sk 1.3 & 2x2700 &  36\\\\ NPM1G+24.0470    & 21 39 41 & +24 24 18 & 0.0390 & 0.0005 & 0.23 & 10.08.2005 & Sk 1.3 &   2400 &  26\\\\ RXS J21592+0952  & 21 59 12 & +09 52 42 & 0.1003 & 0.0002 & 0.21 & 19.08.2004 & Sk 1.3 & 2x2700 &  39\\\\ RXS J22027$-$1304& 22 02 45 &$-$13 04 53& 0.0391 & 0.0003 & 0.14 & 28.08.2005 & Sk 1.3 & 2x2400 &  22\\\\ RXS J22160+1107  & 22 16 04 & +11 07 26 & 0.0514 & 0.0005 & 0.21 & 10.08.2005 & Sk 1.3 &   2400 &  16\\\\ RXS J22287+3335  & 22 28 46 & +33 35 08 & 0.0906 & 0.0004 & 0.29 & 11.08.2005 & Sk 1.3 & 2x2400 &  35\\\\ NPM1G$-$04.0637  & 22 53 11 &$-$04 08 49& 0.0257 & 0.0002 & 0.14 & 11.08.2005 & Sk 1.3 & 2x2400 &  24\\\\  \\hline\\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\subsection{Line profiles} While fitting the line profiles, we did not follow a multi-Gaussian/Lorentzian approach, because two or three, say Gaussians, fitting the broad profile can hardly have any physical meaning (in the sense that they can hardly represent different physical regions). Furthermore, each complex profile can be fitted to the desired level of accuracy with a large enough number of functions; however, the fit is not necessarily meaningful, not even unique in terms of degeneracy. We have to mention, though, that multi-Gaussian fits are often used by other authors, mainly for the sake of a better description of the profile (Warner et al.  2003; Popovic et al.  2004; see also Bachev et al.  2004, and the discussions there). The impression that the narrower profiles are better described by a Lorentzian, not Gaussian, (Sulentic et al.  2000; V\\'{e}ron-Cetty et al.  2001), however, may have no physical meaning, since it is not clear to what extent FWHM measures physically meaningful velocity-related quantity in a complex profile. A more adequate measure would be the second moment of the profile (the velocity dispersion), which is infinity for a Lorentzian. Even if this were not the case (say for a modified Lorentzian), a Lorentzian-type profile typically has higher dispersion than a Gaussian of the same FWHM. Therefore, it would be biased to conclude that from two profiles of the same velocity dispersion, the narrower one (in terms of FWHM!) resembles a Lorentzian more than a Gaussian.   \\subsection{The obscured AGN population} Finding more and more rather bright (i.e. 14--15 magnitude) AGN shows clearly that color-based quasar surveys, like PG for instance, are not complete. This is obviously so, because the quasars are very different, not only in terms of emission-line, radio, X-ray properties, but also in terms of the optical continuum slopes, which ultimately define the colors. The intrinsic quasar continuum obviously may not only depend on the exact accretion solution type and accretion parameters, such as black hole mass, accretion rate, accretion disk inner and outer radii, etc. The continuum shape and level can be severely altered by the host galaxy contribution and the internal absorption, and we found evidence of both in the sample we studied. The reddened continuum slopes can eventually explain why such bright object are not discovered earlier as AGN. On the other hand, a correct assessment of the intrinsic $\\lambda L_{5100}$ is tightly linked to the estimates of the bolometric luminosity and the central mass through empirical relations (the often used $L_{\\rm bol}\\simeq 9 \\lambda L_{5100}$, see also Sect. 3.2). It also affects other characteristics, like optical-to-X-ray slopes, radio-loudness (e.g. Ho, 2002), line equivalent widths, etc. Therefore, one needs to know the intrinsic quasar continuum for accurately computing the accretion parameters and exploring correlations like EV1. For instance, large quasar samples with well-studied spectral properties (e.g. Boroson \\& Green, 1992; Marziani et al.  2003) rarely make an attempt to assess the real continuum level through correcting for the starlight and the intrinsic absorption, presuming that their contribution is small. It may indeed be so for some bright quasars, but as we see here (see also Dietrich et al, 2005), this is not always the case. Furthermore, dangerous biases may be introduced -- the host galaxy will contribute to the continuum mostly for a lower-luminosity AGN, and the internal absorption may be connected with the orientation if associated with a thick torus, and the broad-line widths if the torus is coplanar to a flattened BLR. It should be easy to correct for the host galaxy in a spectrum of relatively good quality. Unfortunately, the host galaxy contribution depends largely on the slit size, seeing, etc. and therefore may differ significantly from observation to observation and from instrument to instrument. This also can have consequences for the reverberation-mapping studies, in the sense that a variable starlight contribution can at least add some extra noise in the cross-correlation function between the continuum and the lines.  \\subsection{Where is dust located?}  In this paper we did not make any attempt to correct the continuum for the intrinsic absorption; nevertheless, the latter is rather apparent for objects like NPM1G-04.063, where quasar continuum around H$\\beta$ appears to be totally absorbed.  One may make an attempt to take this absorption into account either by measuring the continuum slope or by comparing the Balmer decrement with the theoretical expectations (i.e. case ``B'' recombination). Both approaches may lead to incorrect results, however. The slope can, of course, vary for a number of reasons and the dust reddening is only one of them. On the other hand, the Balmer decrement can also be misleading since we do not know exactly where the absorbing material is. If it is uniformly spread outside the NLR, then such a correction can in principle be performed, and the absorption will presumably equally affect the NLR, the BLR, and the central continuum source. Different Balmer decrements for the broad and the narrow lines, as seen in some of the objects in our sample, suggest this may not always be the case. If the absorbing material is at least outside the BLR, then the broad line decrement could be used to assess the unabsorbed continuum level and shape. Unfortunately, we cannot rule out the presence of absorption inside the BLR inner radius (or mixed with the BLR matter), in which case the continuum will be affected, but indications for this might not be seen from the emission line ratios.  An important clue to solving the problem of the dust location can be provided by the relation between the optical continuum slope and the Balmer decrement. We found a correlation between the slope and the decrement absorption indices (Fig. 5). One is to note, however, that the wide range of both -- Balmer decrements and continuum slopes observed in our rather small sample is not typical. For instance, Dong et al. (2008) report a much narrower range of Balmer decrements for a large sample of SDSS quasars -- typically between $-0.5$ and 1, and slopes -- between 1.2 and 2.2 (both expressed in terms of $A_{\\rm V}$). Their sample is, however, blue-selected and is obviously biased against the obscured AGN population. This may explain why Dong et al. (2008) find no correlation between the decrements and the slopes in their sample.   If further confirmed, the relation  $A_{\\rm V}^{\\rm cont}\\simeq 2 A_{\\rm V}^{\\rm BLR}$ (Sect. 3.4) suggests that the reddening on average affects the center more than the BLR. Even adopting a different value for the intrinsic continuum slope (like the often assumed $a_{\\rm \\nu}^{\\rm intr}\\simeq 1$) only offsets the relation between the Balmer decrement and the continuum slope, but does not change the slope of the relation. Adopting a different value for the intrinsic Balmer decrement (e.g. 2.8 instead of 3.1) has a similar effect. One way to explain this result is to assume that the dust is associated with the BLR matter and the reddening affects the continuum coming from inside more than the BLR emission itself. Another explanation would be that the dust is much farther outside the BLR (probably associated with a torus), and we see the active nucleus just above the torus edge. As a result, the central continuum can appear heavily absorbed, but the BLR light is not that affected, taking their different spatial dimensions into account. If the latter explanation turns out to be the correct one, studying the line profiles of such partially obscured objects can be important for probing the exact BLR geometry and kinematics, since knowing the torus opening angle, the inclination can be estimated for these objects."
    },
    "0808/0808.1161_arXiv.txt": {
        "abstract": "In order to quantitatively test the ability of averaged inhomogeneous cosmologies to correctly describe observations of the large scale properties of the Universe, we introduce a smoothed template metric corresponding to a constant spatial curvature model at any time, but with an evolving curvature parameter. This metric is used to compute quantities along an approximate effective lightcone of the averaged model of the Universe. As opposed to the standard Friedmann model, we parameterize this template metric by exact scaling properties of an averaged inhomogeneous cosmology, and we also motivate this form of the metric by results on a geometrical smoothing of inhomogeneous cosmological hypersurfaces. The purpose of the paper is not to demonstrate that the backreaction effect is actually responsible for the Dark Energy phenomenon by explicitly calculating the effect from a local model of the geometry and the distribution of matter, but rather to propose a way to deal with observations in the backreaction context, and to understand what kind of generic properties have to hold in order for a backreaction model to explain the observed features of the Universe on large scales. We test our hypothesis for the template metric against supernova data and the position of the CMB peaks, and infer the goodness--of--fit and parameter uncertainties. We find that averaged inhomogeneous models can reproduce the observations without requiring an additional Dark Energy component (though a volume acceleration is still needed), and that current data do not disfavour our main assumption on the effective lightcone structure. We also show that the experimental uncertainties on the angular diameter distance and the Hubble parameter from Baryon Acoustic Oscillations measurements -- forseen in future surveys like the proposed EUCLID satellite project -- are sufficiently small to distinguish between a FLRW template geometry and the template geometry with consistently evolving curvature. \\\\ \\\\ {\\bf Keywords:} dark energy theory; CMBR experiments; supernova type Ia ",
        "introduction": "\\label{sec:intro}  On the very large scales the Universe appears to be close to a homogeneous and isotropic state. This is usually modeled by a locally isotropic and hence homogeneous solution of Einstein's equations, namely the standard Friedmann--Lema\\^\\i tre--Robertson--Walker (FLRW) metric. The observations of tiny temperature fluctuations of the Cosmic Microwave Background (CMB) radiation suggest that in the Early Universe deviations from this `background cosmology' were very small, thus motivating the use of linear perturbation theory about the FLRW solution. The experimental confirmation of the predicted baryon acoustic oscillations in the CMB power spectrum \\cite{debernardis,spergeletal}, as well as the distribution of galaxies and clusters on the large scales \\cite{lahav} provide a certain body of evidence in favor of the `concordance model' which results from this perturbative approach (see, however, \\cite{marraPN}, \\cite{huntsarkar}). This standard scenario relies on the assumption that the FLRW cosmology correctly describes the `background cosmology', i.e. the averaged inhomogeneous Universe,  at all times. Even though this hypothesis may be valid in the Early Universe, does it continue to hold even when the Universe becomes more and more structured at late times?  Answering such a question has become even more important in the light of the still unexplained Dark Energy phenomenon in the context of the FLRW paradigm. The luminosity distance measurements to type Ia supernova (SN-Ia) standard candles, when analyzed within the framework of the FLRW Universe, provide strong evidence for a missing component characterized by a negative pressure which, by inducing an accelerated phase of expansion, would be responsible for the observed dimming of far distant SN Ia (\\cite{Perlmutter,Riess}). The simplest scenario to account for these observations is a positive cosmological constant in Einstein's equations. This is often assumed to describe the energy contribution of quantum vacuum fluctuations. Nevertheless, because of the huge discrepancy between the particle physics expected value and the observed one, several alternative scenarios have been investigated. For instance phenomenological models such as a late time slow rolling scalar field (see reviews \\cite{copeland,pilar}) or the Chaplygin gas \\cite{chaplygin} have been proposed to describe this Dark Energy component. Alternatively, there have been several proposals to account for these effects through modification of the laws of gravitation (e.g. braneworlds \\cite{maartens:brane}, scalar--tensor gravity \\cite{esposito}, higher--order gravitational theories \\cite{riccilagrangians,capo}, AWE \\cite{AWE1,AWE2}).  Recently, a third alternative has been considered \\cite{rasanen:darkenergy,kolbetal} that aims at explaining Dark Energy as an effect caused by inhomogeneities. However, most of the approaches which include the effect of inhomogeneities still rely on the postulate that the FLRW solution reliably describes the effective (average) evolution of an inhomogeneous cosmology. For instance this is the case of inhomogeneous universe models in which distances are computed using perturbation theory about a FLRW background \\cite{ruth:luminosity,vanderveldetal}. Other efforts abandon the FLRW model and instead restrict inhomogeneities by strong symmetries, employing exact solutions to Einstein's field equations like the Lema\\^\\i tre--Tolman--Bondi (LTB) metric \\cite{celerier1, rasanen:LTB, nambu, LTBgron, LTBluminosity5, LTBluminosity1, LTBluminosity2, LTBluminosity4, alnes2, alnes3, singh1,Bolej,sussman2}.  In this work we shall test a different approach and exploit the key--insight that the (large--scale) kinematics of a homogeneous--isotropic {\\em state} does not necessarily follow the kinematics of a homogeneous--isotropic {\\em solution}, especially at late epochs characterized by the presence of large matter inhomogeneities. Indeed, the analysis of backreaction effects due to inhomogeneities suggests that there is a wider class of  (large--scale) homogeneous--(almost--)isotropic cosmological models, while smaller scales feature strong inhomogeneities and anisotropies that both are known to exist. In such a case it is natural to ask whether the emergence of Dark Energy can result from the breakdown of the underlying assumption associated with the FLRW cosmology.  If no assumption on the nature of the inhomogeneities is made, i.e. if we do not restrict them to be small deviations from a FLRW background or obeying strong symmetry restrictions, we can still look at effective (average) properties of Einstein's equations. In the simplest case such a programme can be realized by foliating spacetime into flow--orthogonal hypersurfaces, restricting the matter model to `dust', and spatially averaging the scalar parts of Einstein's equations with respect to a collection of free--falling observers (generalized fundamental observers). Whereas such a formalism and the dynamical equations that govern the behavior of the averaged inhomogeneous universe model are well--established \\cite{buchert:review}, the explicit geometry, which lies at the basis of how we measure distances, is left unspecified.  Here we suggest, as a next step, to complement the general kinematical properties of an averaged universe model with an explicit form of a {\\em template metric} that retains the main properties of the standard model of cosmology, such as its isotropy and homogeneity on large scales, but allows for structuring on small scales. The shift in emphasis is from postulating a {\\em strong cosmological principle} that assumes local isotropy about every point and hence homogeneity on all scales, to a {\\em weak cosmological principle} that only assumes (quasi--) isotropy and homogeneity on the largest observable scales. In this context, we retain by assumption the usual description of the Universe at early times, up to decoupling. However, at late times we will have to modify this description.  To summarize let us list our main goals and assumptions. The purpose of this paper is to study the influence of a geometrical effect induced by the coupling between backreaction and averaged spatial curvature, on top of the well--studied kinematical effect of backreaction on the evolution of the effective volume scale factor \\cite{buchert:review}. To realize this study, we propose an ansatz for the effective metric of the large-scale homogeneous model that is motivated by previous results on the smoothing of Riemannian metrics by the Ricci flow. This is used to define an effective background on which the photons propagate; such a background can be considered as a first refinement of the usual FLRW background geometry. In order to get an insight into the effects associated with our prescription, we consider a specific example of backreaction, namely a power law of the effective scale factor. Here we want to stress the fact that our aim is not to show that backreaction effects can be fully responsible for the Dark Energy phenomenon. Instead, we address the converse problem: what is necessary and what kind of generic effects are expected for a backreaction model to be consistent with the cosmological observations? And more specifically, we are interested in understanding what specificities can allow to distinguish between FLRW and averaged models.  The central concept underlying our investigation is that, although the 3--Ricci curvature distribution of an inhomogeneous cosmological slice can be smoothed {\\em at any time} into a constant curvature, the dynamical evolution of the averaged curvature can differ from the evolution of a constant--curvature (homogeneous) model. This deviation has recently been quantified in the framework of perturbation theory \\cite{lischwarz,lischwarz1,lischwarz2} (see also \\cite{Ian} for an estimation in the conformal Newtonian gauge), and since we have arguments why a non--perturbative treatment is necessary for the effects of interest, we shall consider the dominant perturbative mode within a general class of scaling solutions to a backreaction--driven cosmology. Of course, the scaling solutions cannot be expected to fully represent the realistic backreaction effect throughout the whole history of the Universe, but it is considered here for reasons of clarity and simplicity to illustrate the kind of effects expected from the non-trivial geometry, in analogy to studies using parameterizations of the equation of state for Dark Energy.   The paper is organized as follows. We introduce the backreaction context, the key--equations and free parameters of an averaged cosmological model in Section~2. In Section~3 we develop the ansatz for the metric, which is inspired by the study of the Ricci flow deformation of three--dimensional Riemannian initial data sets, and we use this effective metric to compute quantities along an approximate past lightcone, in particular the luminosity distances of cosmological objects. This template metric can be considered as an improvement over the FLRW ansatz that considers that photons follow the null geodesics of a locally homogeneous and isotropic metric. In Section~4 we discuss the constraints on the model parameters as inferred from SN Ia data and the multipoles of the CMB peaks and dips. In particular we calculate the cosmic distances in the template metric and determine the `best-fit' models with the help of exact scaling solutions to the backreaction problem, which include the leading perturbative mode. We conclude this section with a discussion on possible tests of our main assumptions. Finally, we summarize the results of the paper and present an outlook in Section~5. ",
        "conclusions": "In this paper we have addressed the problem of comparing averaged inhomogeneous cosmologies with observational data by proposing to fit the observations with the help of an improved  template metric, whose form is compatible with the kinematical integral properties of a general averaged model. This template metric has been motivated by the fact that the averaged curvature of space--like hypersurfaces of the space--time foliation is not expected to be constant in time. Indeed, the cosmological principle only requires the {\\it spacelike} quantities to be averaged, but does not impose anything on the evolution of these quantities. In other words, the FLRW universe models, which obey a strict cosmological principle, are a very particular subclass of models respecting a weaker cosmological principle presented in this paper. We consider the modified template metric as a first approximation tool for interpreting observations in a Universe that appears homogeneous {\\it on large scales}, but in which the backreaction effect cannot be neglected. That means that the proper effective lightcone along which the cosmological observations are made cannot be simply approximated by a FLRW lightcone. It is indeed important to notice that this template metric has only been introduced to compute quantities on the lightcone.\\\\ Thanks to this prescription for the lightcone, we then deduced constraints on the particular class of scaling backreaction, using the supernov{\\ae} luminosity/redshift distribution, and the positions of the peaks in the CMB spectrum. We found that the non--trivial geometry of the lightcone induces a slight change in the constraints, with respect to the same models in a FLRW geometry, leading to models compatible with the data for higher values of $\\Omega_{m}^{\\now\\CD}$, that is to say with a smaller amount of backreaction. This is particularly true for the leading perturbative mode ($n=-1$).  One should note that the model presented in this paper still needs an acceleration of the effective volume scale factor to reproduce the data. Recent results based on the the Lema\\^\\i tre--Tolman--Bondi (LTB) solution (see \\cite{Bolej} and references therein) show that, on the contrary, one can fit the data without an acceleration of the volume scale factor. The two results do not necessarily disagree. Indeed, in the LTB model, one fits the data with an inhomogeneous metric having a functional degree of freedom ($t_{b}(r)$ in \\cite{Bolej}), and then, one averages the best--fit model to find that there is no acceleration of the effective volume scale factor. In this paper, we have first introduced an effective homogeneous model that we fitted to the data, and instead of having a functional dependence in the effective metric, we have made the assumption that backreaction features a scaling behavior with the scale factor. The comparison between the two approaches could be done once more realistic inhomogeneous metrics are at hand, and a more realistic behavior of the backreaction effect could be implemented. Volume acceleration might be a consequence of over--evolving the backreaction at early times by a strict scaling ansatz.  One important feature of our results is that the cosmic history and the distances are strongly affected by the introduction of a non--FLRW geometry for the past lightcone, even if the constraints are only slightly different in the $(\\Omega_{m}^{\\CD},n)$ plane.  Finally, we discussed two ways of testing the assumption (\\ref{eq:defkappa}) that links the geometry of the lightcone with the kinematical properties of the model. On the one hand, we found that the supernov{\\ae} data are fully compatible with the assumption that the averaged curvature felt by photons along the past null--cone is linked with the averaged Ricci scalar of space--like hypersurfaces according to (\\ref{eq:defkappa}). Unfortunately, the current available data are not sufficiently precise to unambiguously show a preferred selection of this hypothesis; this analysis should be done again with future data providing more statistics. On the other hand, we have calculated the explicit form of a function $C_\\CD (z_{D})$, previously introduced in \\cite{chris} to measure possible violations of the Copernican principle. This form can be considered as a {\\it prediction} of the particular models studied in this paper, and it will be a crucial test demonstrating a quantitatively significant difference to the standard FLRW paradigm. We have shown that the future EUCLID satellite project might be able to distinguish between a FLRW template geometry and the template geometry with evolving curvature in an averaged model  by using joint measurements of a geometrical quantity ($d_{A}(z)$) and a kinematical property ($H(z)$).  The effective metric that we employed has been motivated by physical and mathematical arguments, but it cannot be considered as a fully satisfactory description of the lightcone. In particular, since weak lensing involves a series of local effects of the metric on lightrays rather than just an integrated effect, implementing constraints from weak lensing surveys require a more refined study of the lightcone structure and of its link with spatial averaging. This link can be achieved by implementing the averaging formalism directly on the lightcone; this is the subject of work in progress. Another improvement will come from a closure condition that is better than a simple scaling solution, and that will encode more precisely the time evolution of backreaction. Such a closure condition can be looked for in numerical simulations, analytical approximations (see \\cite{marraetal1, marraetal2, rasanenpeaks} for particular approximations, and \\cite{wiltshire,wiltshire05,wiltshire07} for an interesting perspective), or observations (see \\cite{buchertcarfora:Q} and references therein for remarks on the difficulties of this last approach). Moreover, the complete study of other observables like the full CMB spectrum is still unavailable and will be crucial for the construction and the test of a `concordance model' for averaged inhomogeneous cosmologies.  \\subsection*{\\sl Note added during revision} During the revision of this paper, \\cite{Rosenthal} published a preprint in which they pointed out that, in the effective model with a time varying curvature, the redshift can be calculated from first principles, as it is now done in Section 3.3 of this paper. This refinement of the redshift calculation indeed introduces quantitative changes, actually enhancing the differences between the model presented here and a FLRW model (as compared to the former version of or paper), but does not affect the conclusions of the paper. A few comments are in order to clarify the differences between our analysis and the one proposed in \\cite{Rosenthal}. One must note that instead of performing a full MCMC analysis of the model, the authors of \\cite{Rosenthal} fitted their model to a $\\Lambda$CDM fiducial model. Although it leads to correct models that actually reproduce the data (because $\\Lambda$CDM is a very good fitting model), it is by no way guaranteed that such a procedure provides all the acceptable values for the parameters, nor that it provides the most probable ones, as clearly shown by our analysis. \\subsection*{\\sl Acknowledgements} {\\small JL is supported by a Claude Leon Foundation Postdoctoral Fellowship. JL thanks C. Clarkson and B. Bassett for many valuable discussions and for having pointed out the use of $C$ as a test of the Copernican principle. TB acknowledges support and hospitality by Observatoire de Paris and Universit\\'e Paris 7, as well as by Universit\\'e de Gen\\`eve during working visits. MK acknowledges funding by the Swiss SNF for part of this work and thanks Syksy R\\\"as\\\"anen for valuable discussions.}"
    },
    "0808/0808.3164_arXiv.txt": {
        "abstract": "The detection and quantification of narrow emission lines in X-ray spectra is a challenging statistical task. The Poisson nature of the photon counts leads to local random fluctuations in the observed spectrum that often results in excess emission in a narrow band of energy resembling a weak narrow line. From a formal statistical perspective, this leads to a (sometimes highly) multimodal likelihood. Many standard statistical procedures are based on (asymptotic) Gaussian approximations to the likelihood and simply cannot be used in such settings. Bayesian methods offer a more direct paradigm for accounting for such complicated likelihood functions but even here multimodal likelihoods pose significant computational challenges. The new Markov chain Monte Carlo (MCMC) methods developed in \\citet{vand:park:08} and \\citet{park:vand:08}, however, are able to fully explore the complex posterior distribution of the location of a narrow line, and thus provide valid statistical inference. Even with these computational tools, standard statistical quantities such as means and standard deviations cannot adequately summarize inference and standard testing procedures cannot be used to test for emission lines.  In this paper, we use new efficient MCMC algorithms to fit the location of narrow emission lines, we develop new statistical strategies for summarizing highly multimodal distributions and quantifying valid statistical inference, and we extend the method of posterior predictive p-values proposed by \\citet{prot:etal:02} to test for the presence of narrow emission lines in X-ray spectra. We illustrate and validate our methods using simulation studies and apply them to the {\\it Chandra} observations of the high redshift quasar PG1634+706. ",
        "introduction": "\\label{park:sec:intro}  \\subsection{Scientific Background} \\label{park:sec:sci}  Modern X-ray observations show complex structures in both the spatial and spectral domains of various astrophysical sources. Nonetheless, active galalactic nuclei (AGN) including quasars' nuclei remain spatially unresolved even with the highest-resolution X-ray telescopes. Most of their energy is released within the unresolved core, and only spectral and timing information is available to study the nature of the X-ray emission. Generally speaking, emission and absorption lines constitute an important part of the X-ray spectrum in that they can provide information as to the state of plasma. One of the goals of X-ray data analysis is to understand the components present in the spectrum, and to obtain information about the emission and absorption features, as well as their locations and relation to the primary quasar emission. The detection of weak lines in noisy spectra is the main statistical problem in such analyses: Is a bump observed in the spectrum related to a real emission line or is it simply an artifact of the Poissonian noise?  Although quasars' X-ray spectra are usually featureless as expected based on the Comptonization process \\citep[see for example][]{mark:nowa:wilm:05,sobo:siem:zyck:04,siko:etal:97}, an important X-ray emission feature identified in AGN and quasars spectra is the iron K emission line \\citep[see recent review by][]{miller:07}. Determining the origin and the nature of this line is one of main issues in AGN and quasar research.  This line is thought to come directly from illuminated accretion flow as a fluorescent process \\citep{fabi:06}. The location of the line in the spectrum indicates the ionization state of iron in the emitting plasma, while the width of the line tells us the velocity of the plasma \\citep{fabi:06}. The iron line provides a direct probe of the innermost regions of accretion flow and matter in close vicinity of a black hole.  Absorption features associated with the outflowing matter (warm wind, partial covering absorber) have also been observed in recent X-ray observations \\citep{gall:etal:02,char:etal:02,poun:reev:07}. Although the location and width of absorption lines provide information as to the velocity of the absorber and its distance from the quasar, this article focuses on statistical issues in fitting the spectral location of narrow emission lines, i.e., identifying the ionization state.  There are two parts to the Fe-K-alpha emission line observed in AGN \\citep{yaqoob:etal:01}: one is a broad component thought to be a signature of a relativistic motion in the innermost regions of an accretion flow; the other is a narrow component that is a result of a reflection off the material at larger distances from the central black hole.  A detection of the broad component is challenging as it requires a spectral coverage over a large energy range, so the continuum emission is well determined and the broad line can be separated \\citep{reeves:etal:06}. The relativistic line profile is broad and skewed, and two strong peaks of the emission line that originates in a relativistic disk can be prominent and narrow. While the full profile of the broad line may not be easily separable from the continuum, these two peaks may provide a signature for this line in the X-ray spectrum.  The broad Fe-line gives an important diagnostic of the gas motion and can be used to determine the spin of a black hole \\citep{miller:06}; see also an alternative model for the ``red wing\" component by \\citet{miller:etal:08}. The narrow component of the line gives diagnostics of the matter outside the accretion disk and conditions at larger distances from the black hole; see Fe-line Baldwin effect discussion in \\citet{jiang:etal:06}. Both line components are variable and the line may ``disappear'' from the spectrum \\citep{yaq:etal:01}.  The spectral resolution of X-ray CCD detectors (for example 100-200~eV in ACIS on {\\it Chandra} or EPIC on XMM-{\\it Newton}) is relatively low with respect to the predicted width of narrow ($< 5000-30,000$~km~s$^{-1}$) emission or absorption lines in AGN and quasars. Observations with grating instruments (RGS or HEG) can provide high resolution X-ray spectra, but the effective area of the present X-ray telescopes is too low for efficient AGN detections, and only a handful of bright low redshift sources have been observed with gratings to date \\citep{yaqoob:07}. Therefore mainly the X-ray CCD spectra of lower resolution are used to study large samples of AGN and quasars \\citep[see for example][]{guai:06,page:etal:05,jimen:etal:05}.  Using these relatively low resolution X-ray detectors, the Fe-K-alpha emission line can be narrow enough to be contained entirely in a single detector bin.  In some cases (for example in {\\it Chandra}) the line may occupy a few bins.  In this article we focus on the statistical problem of fitting the spectral location of an emission line or a set of emission lines that are narrow. This is a common objective in high-energy analyses, but as we shall discuss fitting these relatively narrow features poses significant statistical challenges. In particular we find evidence that using line profiles that are narrower than we actually expect the emission line to be can improve the statistical properties of the fitted emission line location.    \\subsection{X-Ray Spectral Analysis} \\label{park:sec:hea}  X-ray spectra, such as those available with the {\\it Chandra X-ray Observatory} carry much information as to the quasar's physics. Taking advantage of the spectral capacity of such instruments, however, requires careful statistical analysis. For example, the resolution of such instruments corresponds to a fine discretization of the energy spectrum. As a result, we expect a low number of counts in each bin of the X-ray spectrum.  Such low-count data make the Gaussian assumptions that are inherent in traditional minimum $\\chi^2$ fitting inappropriate. A better strategy, which we employ, explicitly models photon arrivals as an inhomogeneous Poisson process \\citep{vand:etal:01}.  In addition, data are subject to a number of processes that significantly degrade the source counts, e.g., the absorption, non-constant effective area, blurring of photons' energy, background contamination, and photon pile-up. Thus, we employ statistical models that directly account for these aspects of data collection. In particular, we design a highly structured multilevel spectral model with components for both the data collection processes and the complex spectral structures of the sources themselves. In this highly structured spectral model, a Bayesian perspective renders straightforward methods that can handle the complexity of {\\it Chandra} data \\citep{vand:etal:01,vand:kang:04,vand:etal:06}. As we shall illustrate, these methods allow us to use low-count data, to search for the location of a narrow spectral line, to investigate its location's uncertainty, and to construct statistical tests that measure the evidence in the data for including the spectral line in the source model.    \\subsection{A Statistical Model for the Spectrum} \\label{park:sec:model}  The energy spectrum can be separated into two basic parts: a set of continuum terms and a set of several emission lines\\footnote{The model can be extended to account for absorption lines, but in this paper we focus on additive features such as emission lines.}. We begin with a standard spectral model that accounts for a single continuum term along with several spectral lines. Throughout this paper, we use $\\theta$ as a general representation of model parameters in the spectral model. The components of $\\theta=(\\theta^C,\\theta^L,\\theta^A,\\theta^B)$ represent the collection of parameters for the Continuum, (emission) Lines, Absorption, and Background contamination, respectively. (Notice that the roman letters in the superscripts serve as a mnemonic for these four processes.) Because the X-ray emission is measured by counting the arriving photons, we model the expected Poisson counts in energy bin $j\\in{\\cal J}$, where ${\\cal J}$ is the set of energy bins, as \\begin{eqnarray} \\Lambda_j(\\theta) \\ =\\ \\Delta_j f\\big(\\theta^C,E_j\\big)+ \\sum\\limits_{k=1}^K\\lambda_k\\pi_j\\big(\\mu_k,\\nu_k\\big), \\label{park:eq:ideal} \\end{eqnarray} where $\\Delta_j$ and $E_j$ are the width and mean energy of bin $j$, $f(\\theta^C,E_j)$ is the expected counts per unit energy due to the continuum term at energy $E_j$, $\\theta^C$ is the set of free parameters in the continuum model, $K$ is the number of emission lines, $\\lambda_k$ is the expected counts due to the emission line $k$, and $\\pi_j(\\mu_k,\\nu_k)$ is the proportion of an emission line centered at energy $\\mu_k$ and with width $\\nu_k$ that falls into bin $j$. There are a number of smooth parametric forms to describe the continuum in some bounded energy range; in this article we parameterize the continuum term $f$ as a power law, i.e., $f(\\theta^C,E_j)=\\alpha^CE_j^{-\\beta^C}$ where $\\alpha^C$ and $\\beta^C$ represent the normalization and photon index, respectively. The emission lines can be modeled via the proportions $\\pi_j(\\mu_k,\\nu_k)$ using narrow Gaussian distributions, Lorentzian distributions, or delta functions; the counts due to the emission line are distributed among the bins according to these proportions. While the Gaussian or Lorentzian function parameterizes an emission line in terms of center and width, the center is the only free parameter with a delta function; the width of the delta function is effectively the width of the energy bin in which it resides.  While the model in Equation~\\ref{park:eq:ideal} is of primary scientific interest, a more complex statistical model is needed to address the data collection processes mentioned in \\S\\ref{park:sec:hea}. We use the term {\\it statistical model} to refer to the model that combines the {\\it source} or {\\it astrophysical model} with a model for the stochastic processes involved in data collection and recording. Thus, in addition to the source model, the statistical model describes such processes as instrument response and background contamination.  Specifically, to account for the data collection processes, Equation~\\ref{park:eq:ideal} is modified via \\begin{eqnarray} \\Xi_l(\\theta) \\ =\\ \\sum\\limits_{j\\in{\\cal J}} M_{lj}\\Lambda_j(\\theta)d_j u(\\theta^A,E_j)+\\theta_l^B \\label{park:eq:obs} \\end{eqnarray} where $\\Xi_l(\\theta)$ is the expected observed Poisson counts in detector channel $l\\in{\\cal L}$, ${\\cal L}$ is the set of detector channels, $M_{lj}$ is the probability that a photon that arrives with energy corresponding to bin $j$ is recorded in detector channel $l$ (i.e., $\\mathbf{M}=\\{M_{lj}\\}$ is the so-called redistribution matrix or RMF commonly used in X-ray analysis), $d_j$ is the effective area (i.e., ARF, a calibration file associated with the X-ray observation) of bin $j$, $u(\\theta^A,E_j)$ is the probability that a photon with energy $E_j$ is {\\it not} absorbed, $\\theta^A$ is the collection of parameters for absorption, and $\\theta_l^B$ is a Poisson intensity of the background counts in channel $l$. While the scatter probability $M_{lj}$ and the effective area $d_j$ are presumed known from calibration, the absorption probability is parameterized using a smooth function; see \\citet{vand:hans:02} for details. To quantify background contamination, a second data set is collected that is assumed to consist only of background counts; the background photon arrivals are also modeled as an inhomogeneous Poisson process.   \\subsection{Difficulty with Identifying Narrow Emission Lines} \\label{park:sec:narrow}  Unfortunately, the statistical methods and algorithms developed in \\citet{vand:etal:01} cannot be directly applied to fitting {\\it narrow} emission lines. There are three obstacles that must be overcome in order to extend Bayesian highly structured models to spectra containing narrow lines. In particular, we must develop (1) new computational algorithms, (2) statistical summaries and methods for inference under highly multimodal posterior distributions, and (3) statistical tests that allow us to quantify the statistical support in the data for including an emission line or lines in the model. Our main objective in this paper is to extend the methods of \\citet{vand:etal:01} in these three directions, and to evaluate and illustrate our proposals. Here we discuss each of these challenges in detail.  \\paragraph{Challenge 1: Statistical Computation.} Fitting the location of narrow lines requires new and more sophisticated computational techniques than those developed by \\citet{vand:etal:01}. Indeed, the algorithms that we develop require a new theoretical framework for statistical computation: they are not examples of any existing algorithm with known properties. Although the details of this generalization are well beyond the scope of this article, we can offer a heuristic description; a more detailed description is given in Appendix~\\ref{ap:alg}. Readers who are interested in the necessary theoretical development of the statistical computation techniques are directed to \\citet{vand:park:04,vand:park:08} and \\citet{park:vand:08}.  The algorithms used by \\citet{vand:etal:01} to fit the structured Bayesian model described in \\S\\ref{park:sec:model} are based on the probabilistic properties of the statistical models. For example, the parameters of a Gaussian line profile can be fit by iteratively attributing a subset of the observed photons to the line profile and using the mean and variance of these photon energies to update the center and width of the line profile.  The updated parameters of the line profile are used to again attribute a subset of the photons to the line, i.e., to stochastically select a subset of the photons that are likely to have arisen out of the physical processes at the source corresponding to the emission line.  These algorithms are typically very stable. For example, they only return statistically meaningful parameters because the algorithms themselves mimic the probabilistic characteristics of the statistical model. The family of Expectation/Maximization (EM) algorithms \\citep{demp:lair:rubi:77} and Markov chain Monte Carlo (MCMC) methods such as the Gibbs sampler \\citep{gema:gema:84} are the examples of statistical algorithms of this sort.  A drawback of these algorithms is that in some situations they can be slow to converge. When fitting the location of a Gaussian emission line, for example, the location is updated more slowly if the line profile is narrower. This is because only photons with energies very close to the current value of the line location can be attributed to the line. Updating the line location with the mean of the energies of these photons cannot result in a large change in the emission line location. The situation becomes chronic when a delta function is used to model the line profile: The line location parameter sticks at its starting value throughout the iteration.  It is to circumvent this difficulty that we develop both new EM-type algorithms \\citep{vand:park:04} and new MCMC samplers specially tailored for fitting narrow lines.  Our new samplers are motivated by the Gibbs sampler, but constitute a non-trivial generalization of Gibbs sampling known as partially collapsed Gibbs sampling \\citep{vand:park:08,park:vand:08}; see Appendix~\\ref{ap:alg}. Our updated versions of both classes of algorithms are able to fit narrow lines by avoiding the attribution of photons to the emission line during the iteration. Such algorithms tend to require fewer iterations to converge regardless of the width of the emission line. Because they involve additional evaluation of quantities evolving the large dimensional redistribution matrix, $M$, however, each iteration of these algorithms can be significantly more costly in terms of computing time. A full investigation of the relative merit of the algorithms and a description of how the computational trade-offs can be played to derive optimal algorithms are beyond the scope of this paper. Except in Appendix~\\ref{ap:alg}, we do not discuss the details of the algorithms further in this article; interested readers are directed to \\citet{vand:park:04,vand:park:08} and \\citet{park:vand:08}.   \\paragraph{Challenge 2: Multimodal Likelihoods.} The likelihood function for the emission line location(s) is highly multimodal.  Each mode corresponds to a different relatively likely location for an emission line or a set of emission lines. Standard statistical techniques such as computing the estimates of the line locations with their associated error bars or confidence intervals implicitly assume that the likelihood function is unimodal and bell shaped. Because this assumption is clearly and dramatically violated, these standard summary statistics are unreliable and inadequate.  Unfortunately, there are no readily available and generally applicable simple statistical summaries to handle highly multimodal likelihoods. Instead we must develop summaries that are tailored to the specific scientific goals in a given analysis. Because general strategies for dealing with multimodal likelihood functions are little known to astronomers and specific strategies for dealing with multimodal likelihood functions for the location of narrow spectral lines do not exist, one of the primary goals of this article is to develop and illustrate these methods.  A fully Bayesian analysis of our spectral model with narrow emission lines is  computationally demanding, even with our new algorithms. Thus, we develop techniques that are much quicker and give similar results for the location of emission lines. These methods based on the so-called {\\it profile posterior distribution} do not stand on as firm of a theoretical footing as a fully Bayesian analysis, but are much quicker and thus better suited for {\\it exploratory data analysis}. The profile posterior distribution along with our exploratory methods are fully described and compared with the more sophisticated Bayesian analysis.  \\paragraph{Challenge 3: Testing for the Presence of Narrow Lines.}  In addition to fitting the location of one or more emission lines, we often would like to perform a formal test for the inclusion of the emission lines in the statistical model. That is, we would like to quantify the evidence in a potentially sparse data set for a particular emission line in the source.  Testing for a spectral line is an example of a notoriously difficult statistical problem in which the standard theory does not apply. There are two basic technical problems. First, the simpler model that does not include a particular emission line is on the boundary of the larger model that does include the line. That is, the intensity parameter of an emission line is zero under the simpler model and cannot be negative under the larger model. An even more fundamental problem occurs if either the line location or width is fit, because these parameters have {\\it no value} under the simpler model. The behavior (i.e., sampling distribution) of the likelihood ratio test statistic under the simpler model is not well understood and cannot be assumed to follow the standard $\\chi^2$ distribution, even asymptotically. \\citet{prot:etal:02} propose a Monte-Carlo-based solution to this problem based on the method of posterior predictive p-values \\citep{rubi:84,meng:94a}.  In this article we extend the application of Protassov~{\\it et~al}.'s solution to the case when we fit the location of a narrow emission line, a situation that was avoided in \\citet{prot:etal:02}.  \\subsection{Outline of the Article} \\label{park:sec:outline}  The remainder of the article is organized into four sections. \\S\\ref{park:sec:model-based} reviews Bayesian inference and Monte Carlo methods with an emphasis on  multimodal distributions, outlines our computation methods, proposes new summaries of multimodal distributions, and describes exploratory statistical methods in this setting. We introduce illustrative examples in \\S\\ref{park:sec:model-based}, but detailed spectral analysis is postponed in order to allow us to focus on our proposed methods. In \\S\\ref{park:sec:simul}, a simulation study is performed to investigate the statistical properties of our proposed methods, with some emphasis placed on the potential benefits of model misspecification. \\S\\ref{park:sec:quasar} presents the analysis of the high redshift quasar PG1634+706, and how to test for the inclusion of the line in the spectral model. Concluding remarks appear in \\S\\ref{park:sec:conclusion}. An appendix outlines the computational methods we developed specifically for fitting the location of narrow emission lines. ",
        "conclusions": "\\label{park:sec:conclusion}  This article presents methods to detect, identify, and locate narrow emission lines in X-ray spectra via a highly structured multilevel spectral model that includes a delta function line profile. Modeling narrow emission lines with a delta function causes the EM algorithms and MCMC samplers developed in \\citet{vand:etal:01} to break down and thus requires more sophisticated statistical methods and algorithms. The marginal posterior distribution of the delta function emission line location tends to be highly multimodal when the emission line is weak or multiple emission lines are present in the spectrum. Because basic summary statistics are not appropriate to summarize such a multimodal distribution, we instead develop and use HPD graphs along with a list of posterior modes. Testing for an emission line in the spectrum is a notoriously challenging problem because the value of the line intensity parameter is on the boundary of the parameter space, i.e., zero, under a model that does not include an emission line. Thus, we extend the posterior predictive methods proposed by \\citet{prot:etal:02} to test for the evidence of a delta function emission line with unknown location in the spectrum.  Using the simulation study in \\S\\ref{park:sec:simul}, we demonstrate the potential advantage of model misspecification using a delta function line profile in place of a Gaussian line profile. We show that the delta function line profile may provide more precise and meaningful summaries for line locations if the true emission line is narrow. When multiple lines are present in the spectrum, the marginal posterior distribution of a single delta function line location may indicate multiple lines in the spectrum.  Our methods are applied to the six different {\\it Chandra} observations of PG1634+706 in order to identify a narrow emission line in the X-ray spectrum. Given all the six observations, the most probable delta function line is identified at $2.845_{-0.055}^{+0.045}$~keV in the observed frame. The corresponding rest frame energy for the line is $6.64_{-0.13}^{+0.11}$~keV, which may suggest the high ionization of iron in the emitting plasma. There is some recent evidence that high ionization iron line can be variable on short timescales \\citep[see for example Mkn766 in][]{miller:etal:06}. Such variability would explain no detection of the emission line in one of the six {\\it Chandra} observations."
    },
    "0808/0808.3022_arXiv.txt": {
        "abstract": "We observed an area of 10 deg$^2$ of the Large Magellanic Cloud using the Infrared Camera on board AKARI. The observations were carried out using five imaging filters (3, 7, 11, 15, and 24 micron) and a dispersion prism (2 -- 5 micron, $\\lambda / \\Delta\\lambda$ $\\sim$ 20) equipped in the IRC. This paper describes the outline of our survey project and presents some initial results using the imaging data that detected over 5.9$\\times$10$^5$ near-infrared and 6.4$\\times$10$^4$ mid-infrared point sources. The 10 $\\sigma$ detection limits of our survey are about 16.5, 14.0, 12.3, 10.8, and 9.2 in Vega-magnitude at 3, 7, 11, 15, and 24 micron, respectively. The 11 and 15 micron data, which are unique to AKARI IRC, allow us to construct color-magnitude diagrams that are useful to identify stars with circumstellar dust. We found a new sequence in the color-magnitude diagram, which is attributed to red giants with luminosity fainter than that of the tip of the first red giant branch. We suggest that this sequence is likely to be related to the broad emission feature of aluminium oxide at 11.5 micron. The 11 and 15 micron data also indicate that the ([11] $-$ [15]) micron color of both oxygen-rich and carbon-rich red giants once becomes blue and then turns red again in the course of their evolution, probably due to the change in the flux ratio of the silicate or silicon carbide emission feature at 10 or 11.3 micron to the 15 micron flux. ",
        "introduction": "Owing to its proximity ($\\sim$ 50 kpc; e.g., \\cite{feast1987}) and face-on geometry, the Large Magellanic Cloud (LMC) is an ideal natural laboratory for the study of various astrophysical fields. It is located at a high Galactic latitude of $\\sim$ $-$36$^\\circ$, and we expect less contamination of foreground stars and less interstellar extinction. It is far enough to neglect its depth so that we can reasonably assume that objects in the LMC are all at the same distance from us. The apparent size of the LMC is also in a size, for which the entire galaxy can be surveyed in a reasonable amount of time to study the material circulation processes and star-formation history in a galactic scale. In the meanwhile, it is close enough to resolve and study individual objects even with relatively small telescopes. Moreover, the mean metallicity of the LMC is known to be small ($\\sim$ 1/4) compared to the solar, intriguing us to study the influence of low metallicity on various astrophysical phenomena.   \\begin{figure}[htbp] \\begin{center} \\FigureFile(86mm,86mm){figure1bw.ps}  % \\end{center} \\caption{Observed area of the AKARI IRC survey (thick solid outline), overlaid on the photographic image kindly provided by Mr. Motonori Kamiya. The thick dashed outline indicates the coverage of the \\textit{Spitzer}/SAGE survey (7$^\\circ$$\\times$7$^\\circ$, \\cite{meixner2006}), the thick dash-doted outline shows the coverage of the IRSF/SIRIUS near-infrared survey (6.3$^\\circ$$\\times$6.3$^\\circ$, \\cite{kato2007}), and the thin solid outline represents the coverage of the Magellanic clouds optical photometric survey (8.5$^\\circ$$\\times$7.5$^\\circ$, \\cite{zaritsky2004}).} \\label{fig:surveyregion} \\end{figure}   Because of these characteristics there have been a number of survey projects of the LMC in various wavelengths (Hard X-ray: \\cite{gotz2006}; Soft X-ray: e.g., \\cite{long1981} with Einstein, \\cite{sasaki2000}, \\cite{haberl1999} with ROSAT, and \\cite{haberl2003} with XMM; UV: \\cite{smith1987}; H$\\alpha$: \\cite{kennicut1986}, \\cite{gaustad2001}; Optical: \\cite{zaritsky2004}; NIR: e.g., \\cite{cioni2000a} with DENIS, \\cite{nikolaev2000} with 2MASS, and \\cite{kato2007} with IRSF/SIRIUS; MIR\\&FIR: e.g., \\cite{israel1986} with IRAS, \\cite{egan2003} with MSX, and \\cite{meixner2006} with \\textit{Spitzer}/SAGE; [CII]: \\cite{mochizuki}; CO: e.g., \\cite{mizuno2001} with NANTEN; HI: \\cite{luks1992}, \\cite{kim1998}).  A ground-based near-infrared (NIR) survey detected about fifteen million sources in the $\\sim$ 40 deg$^2$ area of the LMC (\\cite{kato2007}). On the other hand, previous mid-infrared (MIR) surveys detected only a few thousands of sources in $\\sim$ 100 deg$^2$ area of the LMC (\\cite{egan2003}), which is too shallow to compare with the survey observations in other wavelengths. Moreover, the angular resolutions of previous MIR surveys were not high enough to make secure cross-identifications with other survey catalogs. This situation is significantly improved with the advent of the \\textit{Spitzer Space Telescope} (SST; \\cite{werner2004}) and AKARI, which have instruments capable of deep mid-infrared observations with the high spatial resolution. The \\textit{Spitzer} SAGE project (\\cite{meixner2006}) carried out an uniform and unbiased imaging survey of about 49 deg$^2$ area of the LMC, providing photometry for about 2$\\times$10$^5$ sources with all IRAC \\citep{fazio2004a} bands at two epochs separated by 3 months.  We carried out near- to mid-infrared imaging and near-infrared spectroscopic survey toward the LMC with AKARI. The entire LMC was also observed as part of the All-Sky survey at 6 bands in the mid- to far-infrared with AKARI (\\cite{ishihara2006}; \\cite{kawada2007}). AKARI observations provide multi-band (11 bands) data of the LMC from the near- to far-infrared (FIR), which will give us a new insight into the various phenomena occurring in the LMC. In this paper, we provide an overview of our LMC survey project and present some initial results using the preliminary photometric catalog of bright point sources.  Then we will make general analyses using the catalog with an emphasis on the unique characteristics of the IRC observations compared to {\\it Spitzer} observations. The first AKARI/IRC LMC point source catalog is planned to be released to the public in 2009. ",
        "conclusions": "We carried out an imaging and spectroscopic survey of a 10 deg$^2$ area of the LMC using IRC onboard AKARI. In this paper we discuss the imaging data obtained in the survey and describe a preliminary photometric catalog of bright point sources. The first point source catalog is planned to be released to the public in 2009. We cross-identified our catalog with the existing optical, near-infrared, and mid-infrared photometry catalogs, and used the combined data to build color-magnitude diagrams using several combinations of interesting wavebands and also to calculate absolute bolometric magnitudes. By virtue of the S11 and L15 bands that are unique to IRC, we found new interesting features in the color-magnitude diagrams. The present IRC data and the upcoming results from AKARI All-Sky survey together with existing radio and near-infrared ground-based observational results will provide a significant database to study the star-formation history and material circulation within a galaxy."
    },
    "0808/0808.3813_arXiv.txt": {
        "abstract": "We present the first results of a multi-object spectroscopic (MOS) campaign to follow up cluster candidates located via weak lensing. Our main goals are to search for spatial concentrations of galaxies that are plausible optical counterparts of the weak lensing signals, and to determine the cluster redshifts from those of member galaxies. Around each of 36 targeted cluster candidates, we obtain $15-32$ galaxy redshifts. For 28 of these targets, we confirm a secure cluster identification, with more than five spectroscopic galaxies within a velocity of $\\pm 3000$km/s. This includes three cases where two clusters at different redshifts are projected along the same line-of-sight. In 6 of the 8 unconfirmed targets, we find multiple small galaxy concentrations at different redshifts, each containing at least three spectroscopic galaxies. The weak lensing signal around those systems is thus probably created by the projection of groups or small clusters along the same line-of-sight. In both the remaining two targets, a single small galaxy concentration is found.  We evaluate the weak lensing mass of confirmed clusters via two methods: aperture densitometry and by fitting to an NFW model. In most cases, these two mass estimates agree well. In some candidate super-cluster systems, we find additional evidence of filaments connecting the main  density peak to additional nearby structure. For a subsample of our most cleanly measured clusters, we investigate the statistical relation between their weak lensing mass ($M_{\\rm NFW}$, $\\sigma_{\\rm sis}$) and the velocity dispersion of their member galaxies ($\\sigma_{\\rm v}$), comparing our sample with optically and X-ray selected samples from the literature. Our lensing-selected clusters are consistent with $\\sigma_v = \\sigma_{\\rm sis}$, with a similar scatter to the optically and X-ray selected clusters. We thus find no evidence of selection bias compared to these other techniques. We also derive an empirical relation between the cluster mass and the galaxy velocity dispersion, $M_{200}=9.6\\times 10^{14}\\times (\\sigma_v/1000$km/s$)^{2.7}/E(z) h^{-1}M_\\odot$, which is in reasonable agreement with the prediction of $N$-body simulations in the $\\Lambda$CDM cosmology. ",
        "introduction": "\\label{sec:intro}  The development of weak lensing techniques, coupled with deep panoramic imaging surveys, has enabled us to locate clusters of galaxies via the gravitational distortion of background galaxies' shapes. Since the first, spectroscopically confirmed discovery of a shear-selected cluster by Wittman et al. (2001), there has been rapid progress toward a large, weak-lensing selected cluster catalogue. Miyazaki et al. (2003) first reported the detection of several significant shear-selected cluster candidates in an untargeted 2.1 deg$^2$ field. Hetterscheidt et al. (2005) found 5 cluster candidates in 50 randomly selected VLT FORS1 fields (0.64 deg$^2$ in total), all of which are associated with an overdensity of galaxies. Wittman et al. (2006) reported 8 candidates in the 8.6 deg$^2$ Deep Lens Survey. Gavazzi \\& Soucail (2007) found 14 cluster candidates in the 4 deg$^2$ CFHT Legacy Survey (Deep), of which nine have optical or X-ray counterparts and are thus secure clusters.  The first sizable sample of weak lensing shear-selected cluster candidates was presented by Miyazaki et al. (2007; hereafter P1). Their sample was obtained solely via peak finding in weak lensing density maps, and includes 100 significant peaks in a 16.7 deg$^2$ survey area.  Before such a sample is used for statistical cosmological or astronomical analyses, two additional follow-up observations are required. Firstly, each cluster candidate should be confirmed by independent observations, since a fraction of lensing peaks could be false positives from  e.g.\\ the chance tangential alignment of galaxies' intrinsic ellipticities (White, van Waerbeke \\&  Mackey 2002; Hamana, Takada \\& Yoshida, 2004; Hennawi \\& Spergel 2005). Secondly, the redshifts of confirmed clusters need to be determined in order to derive their physical quantities, including mass.  We have conducted a multi-object spectroscopic (MOS) campaign that accomplishes both goals. We have measured the redshifts of a few tens of galaxies within an expected cluster scale radius (or core radius, typically a few arcmins), and searched for spatial concentrations that are plausible optical counterparts of the weak lensing signals. Once a galaxy overdensity is found, it is easy to determine the cluster redshift from the redshifts of member galaxies. It is important to note that cluster confirmation based on prominent galaxy concentrations would not be very effective for very high mass-to-light ratio ($M/L$), galaxy-poor systems. Although our methodology can {\\it confirm} normal or galaxy-rich clusters, the absence of a galaxy concentration in our fairly sparsely-sampled data therefore does not necessarily prove that a weak lensing signal is false.  In addition to our primary goals, multi-object spectroscopic observations provide several useful by-products. If redshifts can be obtained for sufficient galaxies in a cluster, their line-of-sight velocity dispersion provides an estimate of the cluster's dynamical mass. MOS observations can also detect multiple structures along the same line of sight. Because of the relatively broad redshift window function of gravitational lensing, physically unrelated structures in the same line of sight may contribute to a single peak in a weak lensing density map, resulting in an overestimation of the cluster mass (White et al. 2002). It will therefore be important to quantify and properly account for such projections when computing statistics of cluster masses from weak lensing observations.  In this paper, we present results of cluster confirmations and cluster redshifts. We discuss the detailed weak lensing properties of each system and, for a clean subset of our clusters, examine statistical relations between the weak lensing masses and dynamical masses. A statistical analysis of the {\\it entire} sample, taking into account additional selection effects, will be presented in Green et al. (in preparation).  This paper is organized as follows. In \\S2, we discuss our selection of cluster candidates. In \\S3, we describe our new observations, data reduction and measurements of galaxy redshifts. In \\S4, we identify optical counterparts to cluster candidates, and measure their velocity dispersions and dynamical masses. In \\S5, we analyze the weak lensing signal of confirmed clusters. In \\S6, we investigate cluster scaling relations within our sample. In \\S7, we summarize our results. In Appendix~\\ref{appendix:nfw}, we calculate the gravitational lensing shear profile of a truncated NFW model. Detailed discussions of each system, including comparisons of the dynamic and lensing masses, follow in Appendices~\\ref{appendix:prop} (for clusters we have observed) and \\ref{appendix:xmmlss} (for observations taken from the literature).  Throughout this paper, we adopt a flat $\\Lambda$CDM cosmology with the matter density $\\Omega_{\\rm m}=0.3$, the cosmological constant $\\Omega_\\Lambda=0.7$, the Hubble constant $H_0=100 h$ km~s$^{-1}$~Mpc$^{-1}$ with  $h=0.7$.   \\begin{table*} \\caption{Summary of spectroscopic observations. (a) Cluster name in the IAU convention. (b) Field and catalogue number given in P1 (if listed). (c) The peak $\\kappa$ value. (d) The number density of galaxies with $18<mag<23$ (Vega mag) within 1 arcmin of the $\\kappa$ peak, and the mean value over the same Suprime-Cam FOV in parentheses. Here the luminosity is defined by the {\\it SExtractor} {\\tt MAG\\_AUTO} in the $R_C$-band, except for the COSMOS field where $i'$-band data is used. (e) Exposure time. (f) The date of the observation. (g) The number of galaxies for which a spectroscopic redshift was obtained. (h) The number of clusters identified (see \\S \\ref{sec:dynamics} for our cluster identification criteria) and, when no clusters were identified, the number of small galaxy concentrations in parentheses.} \\label{table:obs} \\begin{tabular}{llccllcc} \\hline Name$^{(a)}$ & Field-No.$^{(b)}$ & $\\kappa_{peak}$$^{(c)}$ & $n_g$$^{(d)}$ & $t_{\\rm exp}$$^{(e)}$ &obs date$^{(f)}$ & $N_{spec}$$^{(g)}$ & $N_{c}$$^{(h)}$ \\\\ {} & {} & {} & [arcmin$^{-2}$]& [sec] & {} & {}\\\\ \\hline SL~J0217.3$-$0524 & SXDS & 0.041 & 12(7.2) & $2 \\times 1200+900$ & 2005/12/23 & 24 & 1\\\\ SL~J0217.6$-$0530 & SXDS & 0.041 & 11(7.2) & $3 \\times 900$ & 2005/12/23 & 19 & 0(1)\\\\ SL~J0217.9$-$0452 & SXDS & 0.061 & 12(6.2) & $2 \\times 1200 + 900 $ & 2005/12/9 & 21 & 0(1)\\\\ SL~J0218.0$-$0444 & SXDS & 0.047 & 13(6.6) & $3 \\times 1200 $ & 2005/12/9 & 19 & 0(2)\\\\ SL~J0219.6$-$0453 & SXDS & 0.060 & 12(6.9) & $2 \\times 1200 + 900$ & 2005/12/23 & 25 & 1\\\\ SL~J0222.8$-$0416 & XMM-LSS-23 & 0.060 & 12(6.7) & $2 \\times 1500 + 1200$ & 2005/12/24 & 27 & 1\\\\ SL~J0224.4$-$0449 & XMM-LSS-02 & 0.065 & 14(6.0) & $2 \\times 1200+900 $ & 2005/12/23 & 25 & 1\\\\ SL~J0224.5$-$0414 & XMM-LSS-12 & 0.047 & 15(6.7) & $3 \\times 900 $ & 2004/12/9 & 25 & 1\\\\ SL~J0225.3$-$0441 & XMM-LSS-15 & 0.086 & 7.1(6.0) & $3 \\times 1200 $ & 2005/12/24 & 26 & 1\\\\ SL~J0225.4$-$0414 & XMM-LSS-22 & 0.083 & 14(6.9) & $3 \\times 1200 $ & 2004/12/9 & 20 & 1\\\\ SL~J0225.7$-$0312 & XMM-LSS-01 & 0.114 & 8.1(7.1) & $3 \\times 1200$ & 2005/12/24 & 25 & 1\\\\ SL~J0228.1$-$0450 & XMM-LSS-16 & 0.082 & 13(6.1) & $3 \\times 1200 $ & 2005/12/24 & 25 & 1\\\\ SL~J0850.5$+$4512 & Lynx-08 & 0.099 & 11(6.2) & $3 \\times 900$ & 2005/12/17 & 26 & 1\\\\ SL~J1000.7$+$0137 & COSMOS-00 & 0.092 & 14(7.0) & $3 \\times 900$ & 2004/12/9 & 32 & 1\\\\ SL~J1001.2$+$0135 & COSMOS & 0.081 & 13(8.8) & $2 \\times 900$ & 2004/12/9 & 32 & 2\\\\ SL~J1002.9$+$0131 & COSMOS & 0.047 & 18(8.8) & $2 \\times 1800$ & 2004/5/16 & 21 & 0(2)\\\\ SL~J1047.3$+$5700 & Lockman Hole-05 & 0.101 & 18(6.9) & $3 \\times 1200 $ & 2005/12/23 & 27 & 2\\\\ SL~J1048.1$+$5730 & Lockman Hole-15 & 0.077 & 10(7.8)  & $3 \\times 1200 $ & 2005/12/23 & 27 & 1\\\\ SL~J1049.4$+$5655 & Lockman Hole-06 & 0.087 & 14(6.9) & $3 \\times 1200 $ & 2005/12/23 & 26 & 1\\\\ SL~J1051.5$+$5646 & Lockman Hole-03 & 0.082 & 8.4(8.6)  & $2 \\times 1200 + 600 $ & 2005/6/2 & 15 & 0(2)\\\\ SL~J1052.0$+$5659 & Lockman Hole & 0.085 & 7.1(8.6) & $2 \\times 1200 $ & 2005/6/2 & 24 & 0(2)\\\\ SL~J1052.5$+$5731 & Lockman Hole & 0.075 & 11(8.6)  & $2 \\times 1200 + 600 $ & 2005/6/2 & 17 & 0(2)\\\\ SL~J1057.5$+$5759 & Lockman Hole-00 & 0.091 & 23(6.3) & $3 \\times 1200 $ & 2004/12/9 & 32 & 1\\\\ SL~J1135.6$+$3009 & GD140-00 & 0.117 & 10(6.7) & $2 \\times 900 $ & 2004/12/17 & 17 & 1\\\\ SL~J1201.7$-$0331 & PG1159-035-05 & 0.070 & 16(6.6) & $2 \\times 1800 $ & 2004/5/16 & 18 & 1\\\\ SL~J1204.4$-$0351 & PG1159-035  & 0.079 & 14(6.4) & $3 \\times 900 $ & 2004/12/17 & 20 & 1\\\\ SL~J1334.3$+$3728 & 13hr field-00 & 0.097 & 24(7.1) & $2 \\times 1800 + 420$ & 2004/5/16 & 31 & 1\\\\ SL~J1335.7$+$3731 & 13hr field-01 & 0.074  & 16(7.1) & $3 \\times 900$ & 2005/6/1 & 23 & 1\\\\ SL~J1337.7$+$3800 & 13hr field-04  & 0.082 & 14(7.3) & $3 \\times 1200$ & 2005/6/1 & 27 & 1\\\\ SL~J1601.6$+$4245 & GTO 2deg$^2$ & 0.107 & 14(6.2) & $3 \\times 1200 $ & 2005/7/31  & 29 &2\\\\ SL~J1602.8$+$4335 & GTO 2deg$^2$-00 & 0.099 & 22(7.0) & $3 \\times 900 $ & 2005/6/1 & 22 & 1\\\\ SL~J1605.4$+$4244 & GTO 2deg$^2$-09 & 0.067 & 11(7.8) & $3 \\times 1200 $ & 2005/7/31 & 20 & 1\\\\ SL~J1607.9$+$4338 & GTO 2deg$^2$ & 0.079 & 11(6.5) & $3 \\times 1200 $ & 2005/8/1 & 21 & 1\\\\ SL~J1634.1$+$5639 & CM DRA-06 & 0.104 & 5.1(6.3) & $3 \\times 900$ & 2005/6/1 & 16 & 1\\\\ SL~J1639.9$+$5708 & CM DRA-04 & 0.044 & 12(7.3)  & $3 \\times 1200$ & 2005/8/1 & 24 & 0(2)\\\\ SL~J1647.7$+$3455 & DEEP16-00 & 0.070 & 9.7(5.0) & $3 \\times 1200$ & 2004/12/9 & 32 & 1\\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{sec:summary}  We have presented the results of a multi-object spectroscopic campaign to target 36 cluster candidates located by the Subaru weak lensing survey (Miyazaki et al. 2007). We obtained the redshifts of $15-32$ galaxies within a few arcminutes of each cluster candidate. Our primary goals were to search for a spatial concentration of galaxies as an optical counterpart of each weak lensing density peak, and to determine the cluster redshifts. We found 31 galaxy concentrations containing more than five spectroscopic galaxies within a velocity of $\\pm 3000$km/s, and determined their redshifts. These included 25 detections of isolated clusters, and three systems (SL~J1000.7.3+0137, SL~J1047.3+5700, SL~J1601.6+4245) in which two galaxy clusters are projected at different redshifts along the same line-of-sight. This demonstrates that spectroscopic follow-up of weak lensing cluster candidates is a reliable way not only to identify their optical counterparts but also to distinguish superposed systems.  We have therefore identified secure optical counterparts of the weak lensing signal in 28 out of 36 targets. In 6 of the 8 unconfirmed cluster candidates, we found multiple small galaxy concentrations at different redshifts (each containing at least 3 spectroscopic galaxies). This suggests that the weak lensing signal in those cases may arise from the projection of small clusters along the same line-of-sight. However, it is also possible that a real, massive cluster is responsible for the weak lensing density peak, but was missed by our relatively sparse MOS observations. This is also the case for the final two unconfirmed candidates, where only a single small galaxy concentration was identified. In order to obtain a firm confirmation of the optical counterpart of such unconfirmed candidates, denser spectroscopic observations would be required.  We measured the mass of single cluster systems with known redshifts using two weak lensing methods: aperture densitometry and by fitting the shear profile to an NFW model. In most cases, the two mass estimators agree well: providing observational support for the NFW model. In the few clusters where the mass estimators did not agree, the weak lensing $\\kappa$ signal clearly deviates from spherical symmetry. This could account for the disagreement. It was also found, by eye, that the aperture mass profile of some clusters does not flatten even at a large radius of $\\theta\\sim 10$ arcmin. This can be accounted for by the mass contribution from surrounding structures. We found some candidates of super-cluster systems, whose weak lensing mass reconstructions show evidence of filamentary structure connecting the main cluster to surrounding systems.   We investigated statistical relations between clusters' weak lensing properties ($\\sigma_{\\rm sis}$ and $M_{\\rm vir}$ or $M_{200}$) and the velocity dispersion of their member galaxies ($\\sigma_{\\rm v}$), comparing our results to optically and X-ray selected cluster samples from the literature. Although our {\\it clean sample} contained only 12 clusters, we found our clusters to be consistent with $\\sigma_v = \\sigma_{\\rm sis}$, with a scatter as large as that of optically and X-ray selected samples. Therefore, as far as the relation between $\\sigma_v$ and $\\sigma_{\\rm sis}$ is concerned, no strong bias between the cluster selection techniques was identified. We also derived the empirical relation between the cluster virial mass and the galaxy velocity dispersion: $M_{\\rm vir}(1+z_c)^{1.08}=(13\\pm 2) \\times 10^{14}\\times (\\sigma_v/1000$km/s$)^{2.7\\pm 0.6} h^{-1}M_\\odot$. The derived $M_{\\rm vir}-\\sigma_{\\rm SIS}$ relation is similar to theoretical expectations from the SIS model, eq. (\\ref{eq:Msisvir}). It is important to note that, unlike the SIS model assumption, real cluster shear profiles (and density profiles) are {\\it not} single power-laws, so this result may depend upon details of the fitting technique. For comparison with numerical simulations, we also derived the $M_{200}-\\sigma_v$ relation and found $M_{200}~E(z) =(9.6\\pm 1.6)\\times10^{14}\\times(\\sigma_v/1000{\\rm km/s})^{2.7\\pm 0.6} h^{-1}M_\\odot$. This is in good agreement with predictions by Evrard et al.\\ (2008), demonstrating that the dynamical structure of galaxy clusters is similar to that expected in the standard CDM paradigm of structure formation.   \\bigskip We are very grateful to Subaru astronomers: Y. Ohyama, K. Aoki and T. Hattori for their dedicated supports of the FOCAS observing. Numerical computations presented in this paper were carried out on computer system at CfCA (Center for Computational Astrophysics) of the National Astronomical Observatory Japan. Data reduction and analysis were in part carried out on general common use computer system at ADAC (Astronomical Data Analysis Center) of the National Astronomical Observatory of Japan. This research was supported in part by the Grants--in--Aid from Monbu--Kagakusho and Japan Society of Promotion of Science: Project number 15340065 (TH\\&SM) and 17740116 (TH)."
    },
    "0808/0808.1033_arXiv.txt": {
        "abstract": "The young star cluster NGC~602 and its associated HII region, N90, formed in a relatively isolated and diffuse environment in the Wing of the Small Magellanic Cloud. Its isolation from other regions of massive star formation and the relatively simple surrounding HI shell structure allows us to constrain the processes that may have led to its formation and to study conditions leading to massive star formation. We use images from \\emph{Hubble Space Telescope} and high resolution echelle spectrographic data from the Anglo-Australian Telescope along with 21-cm neutral hydrogen (HI) spectrum survey data and the shell catalogue derived from it to establish a likely evolutionary scenario leading to the formation of NGC~602. We identify a distinct HI cloud component that is likely the progenitor cloud of the cluster and HII region which probably formed in blister fashion from the cloud's periphery. We also find that the past interaction of HI shells can explain the current location and radial velocity of the nebula. The surrounding Interstellar Medium is diffuse and dust-poor as demonstrated by a low visual optical depth throughout the nebula and an average HI density of the progenitor cloud estimated at 1~cm$^{-3}$. These conditions suggest that the NGC~602 star formation event was produced by compression and turbulence associated with HI shell interactions. It therefore represents a single star forming event in a low gas density region. ",
        "introduction": "The young star cluster NGC~602 (J2000 01:29:30, -73:34:0.0) and its associated nebula, N90 are located in the Wing of the Small Magellanic Cloud (SMC). This is the relatively diffuse southeastern region which transitions to the bridge connecting the SMC and the Large Magellanic Cloud (LMC). Massive star formation is naturally a topic of perennial interest and the range of conditions under which stars form at various epochs and host environments makes it a very diverse field. The NGC~602 cluster is an example of massive star formation in a region with diffuse interstellar medium (ISM), having the low metallicity characteristic of the SMC and without any apparent direct kinetic trigger such as a recent nearby supernova. Its isolation alone makes it a great laboratory for the study of star formation since the complexities introduced by kinetic and optical interaction with other star-forming regions are eliminated. This combined with low metallicity and relatively quiescent environment also makes it a candidate for comparison to theoretical work on primordial HII regions driven by Population III stars (e.g., \\citealp{abel2007}).  NGC~602 consists of a central star cluster surrounded by the N90 HII region. Its location in the eastern periphery of the SMC is indicated in Figure~\\ref{fig:NGC-602-and-} from the Magellanic Cloud Emission Line Survey (MCELS) \\footnote{SMC image credit: F. Winkler/Middlebury College, the MCELS Team, and NOAO/AURA/NSF. The image was obtained from the web page http://www.noao.edu/image\\_gallery/html/im0993.html.}. Also shown in the figure is a HI column density image of the NGC~602 neighborhood taken from the HI survey data used in our study described in \\S~\\ref{sec:Data}. The nominal outline of the H$\\alpha$ supergiant shell SMC-1, identified by \\citet{meaburn1980}, is shown and at its southern tip, the nebula N90 is outlined by a H$\\alpha$ contour. We adopt a distance to the SMC of 60~kpc \\citep{hilditch2005} and this results in a transverse scale of 0.29~pc~arcsec$^{-1}$.  The H$\\alpha$+{[}NII] \\emph{Hubble Space Telescope} (HST) Advanced Camera for Surveys (ACS) image of NGC~602 and N90 is shown in Figure \\ref{fig:NGC602-Halpha-Slits} from which the detailed structure of N90 can be appreciated, but does not do it full justice. A revealing color composite of all three HST ACS images used in this study has been published in \\citet{carlson2007}. High angular resolution reveals a very rich structure in the photodissociation regions (PDRs) along ridges and faces to the east, south and west in NGC~602. The nebula is host to many prototypical structures found in HII regions including ``elephant trunks'' and ``pillars of creation'' along the northwestern ridge and what is likely a molecular outcropping to the southeast. A catalogue and analysis of the many features found in NGC~602 will be the subject of a new paper.  Also notable are the many galaxies visible both outside and within the nebula, including a prominent grand design spiral seen nearly face-on north of the eastern HII face. The presence of distinct PDRs to the east, south and west, and the openness toward the north suggests a blister morphology. Since the largest dense condensation of gas, which also is in proximity to most of the young stars, is in the south-southeast part of the nebula, it is most likely that the cluster emerged from this region with the first breakout being in the projected northern and western directions. However, even to the south-southeast, there are galaxies visible indicating that there is no obvious remaining extended dense gas concentration there. This transparency suggests a lack of large-scale dust and molecular concentrations anywhere nearby.  The massive stars associated with NGC~602 are in two groups, one large group in the southeast and one smaller group to the northwest. The stellar population of NGC~602 has been studied by \\citet{westerlund1964}, \\citet{hodge1983}, \\citet{hutchings1991} and most recently by \\citet{carlson2007} and \\citet{schmalzl2008}. \\citeauthor{westerlund1964} estimated the maximum age at about 10~Myr. \\citeauthor{hodge1983} re-analyzed Westerlund's color-magnitude data and placed the age at 16$\\pm5$~Myr. \\citeauthor{hutchings1991} obtained optical and UV spectra of the OB stars in the cluster and estimated the age at 5~Myr. \\citeauthor{carlson2007}, used HST ACS optical and Spitzer IR images and estimated the age of the OB stars at 4 Myr and suggest that the young objects in the two ridges are a result of progressive star formation. Progressive star formation is also suggested by \\citet{gouliermis2007}.  The diffuse, quiescent nature of this environment and its low metallicity raises questions as to how massive stars evidently formed in what appears to be a relatively isolated region of dense gas. By analyzing optical images, we look for evidence of the molecular gas substructure left behind in the cluster's formation. In order to understand the kinematic relationship of NGC~602 to the larger scale environment, we compare radial velocities of the HII gas obtained through high resolution H$\\alpha$ longslit spectra to 21-cm HI SMC survey data. We also use the kinematics of HI shells in the region to develop a simple scenario that could have led to the formation of this cluster. ",
        "conclusions": "The young star cluster NGC~602 and its associated HII region N90 formed in relative isolation in the Wing of the SMC under conditions that are generally marginal for star formation, the ISM in this region being relatively diffuse and quiescent. A general transparency to galaxies throughout the nebula and its surroundings shows that there are no large-scale dust and molecular concentrations in or near the nebula to explain its formation through simple gravitational collapse. There is also no morphological evidence of violent events such as supernovae or strong stellar winds having directly triggered its initial formation as demonstrated by the low radial gas velocity gradient and dispersion measured across the nebula. Our measured low velocity gradient across the nebula also shows that NGC~602's formation and evolution has also been peaceful and, to our good fortune, has left the evidence of its birth relatively undisturbed. Star formation triggered by gas flow interaction might provide a consistent explanation for NGC~602's formation under these constraints.  Based on a tight, unambiguous neutral and ionized gas velocity correlation at $\\sim$180 $\\mbox{km\\, sec}^{-1}$, we determined that NGC~602 formed on the periphery of a low-density HI cloud component immediately to its south and blistered from it. The HI density of the natal cloud, designated {HI~J0130-7337+180H} was estimated at 1.3~cm$^{-3}$ and its size at 220~pc implying a mass of $3\\times10^{5}M_{\\odot}$. This mass is sufficient to provide NGC~602's approximate stellar mass with $\\leq0.7\\%$ star formation efficiency. Given the temperature range of cool cloud components in the SMC of 10~K~to~150~K, the cloud's size is comparable to the Jeans length for a simple gravitational collapse scenario. The timescale for undisturbed collapse is $\\sim$40~Myr, much longer than the dynamics of the associated HI shells in the SMC would allow. The interaction of large-scale gas components likely accelerated the formation process, producing local molecular concentrations from which the cluster formed. The pattern of transparency to galaxies in the region indicate the presence of fragments of these molecular components in and around the nebula connecting nebular features to the natal cloud.  A simple model of local HI shell evolution shows that two of the expanding shells would have begun interacting $\\sim$7~Myr ago and that the nominal path of this interaction region is consistent with the location of NGC~602. Turbulence at the intersection could have led to the formation of this cluster $\\sim$3~Myr later. Further study of the NGC~602 region thus offers an opportunity to explore star formation processes in a low density, relatively quiescent environment."
    },
    "0808/0808.4085_arXiv.txt": {
        "abstract": "Employing the perturbative treatment of gravitational clustering, we discuss possible effects of primordial non-Gaussianity on the matter power spectrum. As gravitational clustering develops, the coupling between different Fourier modes of density perturbations becomes important and the primordial non-Gaussianity which intrinsically possesses a non-trivial mode-correlation can affect the late-time evolution of the power spectrum. We quantitatively estimate the non-Gaussian effect on power spectrum from the perturbation theory. The potential impact on the cosmological parameter estimation using the power spectrum are investigated based on the Fisher-matrix formalism. In addition, on the basis of the local biasing prescription, non-Gaussian effects on the galaxy power spectrum are considered, showing that the scale-dependent biasing arises from a local-type primordial non-Gaussianity. On the other hand, an equilateral-type non-Gaussianity does not induce such scale-dependence because of weaker mode-correlations between small and large Fourier modes. ",
        "introduction": "Recent observations of cosmic microwave background (CMB) anisotropies as well as density perturbations in the large-scale structure strongly support the basic predictions of inflationary scenarios, in which primordial adiabatic fluctuations were produced during the accelerated phase of the cosmic expansion and their statistical properties are approximately described by Gaussian statistics with a nearly scale-invariant power spectrum (e.g., \\cite{Komatsu_WMAP5, Tegmark_etal2006}). Some specific inflationary models have been ruled out by observations, narrowing the constraints on the early stage of the universe. With precision measurements from future observations, we will detect clear signals that help us to discriminate between many candidate of inflationary models.   Amongst several signals accessible in the near future, departure from Gaussianity is an important clue to probe the generation mechanism for primordial perturbations as well as to discriminate between inflation models. While the simplest slow-roll inflation with single scalar field predicts a small departure from Gaussianity (e.g., \\cite{ABMR2003,Maldacena2003,BKMR2004}), models with a non-trivial kinetic term or late-time inflaton decay called curvaton scenario \\cite{LW2002,MT2001} can produce a large non-Gaussianity (e.g., \\cite{SL2005,LR2005,SVW2006, Assadullahi:2007uw}). There are also viable models motivated by string theory that produce large non-Gaussianity such as Dirac-Born-Infeld (DBI) inflation \\cite{AST2004, Chen:2006nt, Huang:2006eh, Huang:2006eh, Arroja:2008ga, Langlois:2008qf, AMK2008} and ekyprotic scenario \\cite{KOST2001, Creminelli:2007aq, KMVW2007, Buchbinder:2007at, Lehners:2007wc}. Although tentative detections of primordial non-Gaussianity have been reported very recently (\\cite{YW2008,MHLM2008}, but see also \\cite{Komatsu_WMAP5,HMCLHM2008}) and the result is broadly consistent with the standard prediction of slow-roll inflation, these detections are at relatively low statistical significance and more precise measurements are necessary for a definite detection of the non-Gaussian signals. Planned CMB surveys, namely the Planck mission \\cite{Planck}, will have much greater sensitivity and better detect non-Gaussian signals (e.g., \\cite{SC2008}). Further, large-scale structure data from future spectroscopic surveys will uncover the mass density fluctuations up to Giga parsec scales, which preserve the statistical properties of primordial fluctuations (e.g., \\cite{SSZ2004,SK2007}). Since these two measurements can probe different scales of primordial fluctuations, they are, in principle, complementary to each other.    In the present paper, we study the signature of primordial non-Gaussianity imprinted on the large-scale structure, especially focusing on the matter power spectrum. Naively, we expect that the signature of primordial non-Gaussianity basically appears in the statistical properties of higher-order quantities such as the bispectrum and trispectrum, and the power spectrum as a second-order statistic remains unchanged even if large non-Gaussianity has been produced. However, this is true only when the fluctuations of the mass density field are tiny and well within the linear regime. As the gravitational clustering develops, the coupling between different Fourier modes becomes important and the scale-dependent non-linear growth appears due to the mode-coupling effect. In the weakly non-linear regime, the strength of this mode-coupling sensitively depends on the initial condition. Since the non-Gaussian density field intrinsically possesses non-trivial mode-correlations, the evolution of the power spectrum may be altered by the primordial non-Gaussianity, at least in the weakly non-linear regime.   Here, employing the perturbative calculations, we study the effect of primordial non-Gaussianity on the matter power spectrum. The effect of non-Gaussianity on the matter power spectrum has been previously investigated by Ref.\\cite{SSZ2004} using perturbation theory, showing that the magnitude of this effect is roughly at the few percent level in the weakly non-linear regime. In this paper, adopting the two representative models of primordial non-Gaussianity described in Sec.~\\ref{sec:nonGaussian_model}, we quantitatively estimate the non-Gaussian effects on the power spectrum and elucidate their shape dependence (Sec.~\\ref{sec:non_GaussianPT}). In particular, we carefully examine the influence of primordial non-Gaussianity on cosmological parameter estimation, paying particular attention to the primordial spectral indices and the characteristic scale of the baryon acoustic oscillations as a cosmic standard ruler (Sec.~\\ref{sec:Fisher}). Also, we discuss the non-Gaussian effect on the galaxy biasing (Sec.~\\ref{sec:galaxy_biasing}). This issue has been recently studied by several authors in the context of the halo biasing prescription \\cite{DDHS2008,MV2008,SHSHP2008,McDonald2008,AT2008}, and local biasing prescription \\cite{McDonald2008}. Here, we adopt the local biasing scheme and consider the scale-dependent galaxy biasing arising from the primordial non-Gaussianity.   Throughout the paper, we assume a flat Lambda cold dark matter (CDM) model and the fiducial model parameters are chosen based on the five-year WMAP results (WMAP+BAO+SN ML, see Ref.\\cite{Komatsu_WMAP5}): $\\Omega_{\\rm b}h^2=0.02263$, $\\Omega_{\\rm c}h^2=0.1136$, $\\Omega_{\\rm K}=0\\,\\, (\\Omega_{\\Lambda}=0.724)$, $n_s=0.961$, $\\tau=0.080$, $\\Delta_{\\mathcal R}(k=0.002\\mathrm{Mpc}^{-1}) =2.42\\times10^{-9}$, $h=0.703$. ",
        "conclusions": "In this paper, we have comprehensively studied the effect of primordial non-Gaussianity on the power spectrum of large-scale structure. Using perturbation theory, we calculate the leading-order non-Gaussian corrections to the matter power spectrum through the non-linear mode coupling of the gravitational evolution. In the weakly non-linear regime, the signature of primordial non-Gaussianity in the matter power spectrum can be characterized by the primordial bispectrum. Adopting two representative models of the bispectrum (local and equilateral), we quantitatively estimate the non-Gaussian signals on the matter power spectrum. The primordial non-Gaussianity systematically enhances or suppresses the non-linear growth of the power spectrum amplitude at the $\\lesssim$ 1-2\\% level.   We then explore the potential impact on the cosmological parameter estimations from future dedicated surveys for BAO measurement. Under the assumption of scale-independent linear biasing, the Fisher-matrix analysis  reveals that while it is hard to detect the primordial non-Gaussianity to the level of current constraints, the inclusion of the non-Gaussian parameter $\\fnl$ significantly degrades the constraints on the primordial spectral index $n_s$ and the running of the index $\\alpha$. In this respect, the CMB prior information for $n_s$ and $\\alpha$ as well as the non-Gaussian parameter is very crucial. On the other hand, determination of the distance scale $D_A(z)$ is rather insensitive to the presence or absence of primordial non-Gaussianity, and thus the characteristic scale of BAOs is a robust standard ruler.   We have also considered the effects of primordial non-Gaussianity on the galaxy biasing. In the framework of local galaxy biasing, in which the number density of galaxies is described by a local function of the mass density field, we found that the non-linear mapping of galaxy biasing can modulate the power spectrum and this sensitively depends on the shape of non-Gaussianity. In the local model of primordial non-Gaussianity, mode-correlations between large scales and small scales are fairly strong, and the non-Gaussian correction induces the strong scale-dependent biasing on large-scales, while the scale-independent linear biasing is preserved to a good accuracy in the equilateral model. These remarkable properties may be very helpful in discriminating between the types of primordial non-Gaussianity, especially in connection with the single-field consistency relation. However, the biasing factor in the local model exhibits an apparent divergence, suggesting that the local biasing prescription may be incompatible with the local model of primordial non-Gaussianity. Indeed, the local biasing scheme is just a phenomenological parametrization and the validity of this prescription itself has not yet been tested enough in the presence of non-Gaussianity. Further study using simulations is necessary for quantitative prediction of the non-Gaussian signature on galaxy biasing."
    },
    "0808/0808.0203_arXiv.txt": {
        "abstract": "I review the theory of angular momentum acquisition of galaxies by tidal torquing, the resulting angular momentum distribution, the angular momentum correlation function and discuss the implications of angular momentum alignments on weak lensing measurements: Starting from linear models for tidal torquing I summarise perturbative approaches and the results from $n$-body simulations of cosmic structure formation. Then I continue to discuss the validity of decompositions of the tidal shear and inertia fields, the effects of angular momentum biasing, the applicability of parameterised angular momentum correlation models and the consequences of angular momentum correlations for shape alignments. I compile the result of observations of shape alignments in recent galaxy surveys as well as in $n$-body simulations. Finally, I review the contamination of weak lensing surveys by spin-induced shape alignments and methods for suppressing this contamination. ",
        "introduction": "This article reviews models of angular momenta of galaxies and angular momentum correlations in the cosmological large-scale structure, and the implications of angular momentum induced shape alignments of galaxies for gravitational lensing. The current paradigm states that galaxies acquire their angular momentum by a mechanism known as tidal torquing  \\citep{hoyle, 1969ApJ...155..393P, 1955MNRAS.115....2S, 1970Ap......6..320D, 1984ApJ...286...38W}, where tidal fields from the ambient large-scale structure excert a torquing moment onto the protogalactic object prior to gravitational collapse. The angular momentum, which is conserved during collapse, has important implications for the formation and orientation of the galactic disk. Tidal torquing explains naturally angular momentum correlations, as neighbouring galaxies form from correlated initial conditions and experience similar tidal fields. An important application of angular momentum correlations is gravitational lensing: The angular momentum-induced ellipticity alignments are a source of systematics, which, if uncorrected, deteriorates constraints on cosmological parameters and introduces estimation biases.  After a compilation of the key formul{\\ae} of cosmology, structure formation and perturbation theory in Sect.~\\ref{sect_cosmology}, I review angular momentum build-up by tidal torquing, which is a second order perturbative process arising in Lagrangian perturbation theory, in Sect.~\\ref{sect_torquing}. The distribution of angular momentum values and approximations used in the derivation of the angular momentum correlation function are discussed in Sects.~\\ref{sect_distribution} and~\\ref{sect_correlation}, respectively. Analytical derivations are contrasted with the results from $n$-body simulations of cosmic structure formation in Sect.~\\ref{sect_simulation} and to observations of angular momentum induced galaxy shape correlations in Sect.~\\ref{sect_observation}. Implications of such spin-induced shape correlations for weak gravitational lensing are derived in Sect.~\\ref{sect_lensing}, and a summary in Sect.~\\ref{sect_summary} concludes the review article.  The choice of specific values for the cosmological parameters is not as relevant for angular momentum correlations compared to other cosmological observations, because the angular momentum build-up of galaxies takes place at early times when the universe is matter dominated. Nevertheless I will attempt to keep the derivations as general as possible and use as a framework spatially flat dark energy cosmological models with Gaussian adiabatic initial perturbations in the cold dark matter density field. ",
        "conclusions": "\\label{sect_summary} This review article summarises models for galactic angular momenta and angular momentum correlations in the cosmological large-scale structure, and its implications for gravitational lensing. Starting from perturbative models and tidal torquing, angular momentum distributions and correlation functions are derived and simplifications and their applicability discussed. In general, linear tidal torquing is a very successful theory for explaining the angular momentum magnitude and direction, only the nonlinear stages become difficult to describe analytically, as well as the influence of baryonic dynamics. The many functional forms used show the ambiguity in defining an angular momentum correlation function. Interesting future developments might include higher-order statistics of the angular momentum field, which is intrinsically non-Gaussian even if it develops from Gaussian initial fluctuations, or for non-Gaussianities in the initial conditions, either primordial or due to translinear structure formation on small scales. Simulations find correlations between angular momenta and the strong tidal fields of nonlinear structures, illustrating the richness of the cosmological large-scale structure which can not be quantified using 2-point functions alone.  The primary application of angular momentum models are ellipticity correlations, which have been shown to be a serious contaminant of high-precision weak lensing measurements. A plethora of methods has been suggested for removing or suppressing intrinsic ellipticity correlations such that the power of weak lensing surveys for constraining cosmological parameters can be harnessed. Nevertheless one could expect futher development in predicting statistical properties of the intrinsic alignments, especially in comparing linear and quadratic alignment models, and numerical studies of how the ellipticity measured from the luminosity distribution is related to the angular momentum. The suppression and nulling techniques invented to suppress the intrinsic alignment can of course be used to achieve an amplification of the alignment signal over the weak lensing, and it might be possible to combine higher order image distortions such as weak flexions, which should be free of intrinsic alignments, with the weak shear signal, enhancing the understanding of intrinsic and weak lensing induced ellipticity correlations. \\spirou{Insight into the alignment of the stellar light distribution with respect to the dark matter host structure could possibly be gained using strong lensing \\citep[an overview is provided by][]{2006glsw.conf.....M}, by measuring the halo's potential well and combining with kinematical data.}"
    },
    "0808/0808.1080_arXiv.txt": {
        "abstract": "Recently there have been suggestions that the Type Ia supernova data can be explained using only general relativity and cold dark matter with no dark energy. In ``Swiss cheese'' models of the Universe, the standard Friedmann-Robertson-Walker picture is modified by the introduction of mass-compensating spherical inhomogeneities, typically described by the Lema\\^{i}tre-Tolman-Bondi metric.  If these inhomogeneities correspond to underdense cores surrounded by mass-compensating overdense shells, then they can modify the luminosity distance-redshift relation in a way that can mimic accelerated expansion.  It has been argued that this effect could be large enough to explain the supernova data without introducing dark energy or modified gravity. We show that the large apparent acceleration seen in some models can be explained in terms of standard weak field gravitational lensing together with insufficient randomization of void locations. The underdense regions focus the light less than the homogeneous background, thus dimming supernovae in a way that can mimic the effects of acceleration. With insufficient randomization of the spatial location of the voids and of the lines of sight, coherent defocusing can lead to anomalously large demagnification effects. We show that a proper randomization of the voids and lines of sight reduces the effect to the point that it can no longer explain the supernova data. ",
        "introduction": "Supernova cosmology \\cite{Riess, Perlmutter} has taught us that the Universe is not well described by a homogeneous and flat matter dominated model, with gravity governed by general relativity.  Type Ia supernovae at a given redshift appear dimmer than expected based on this simple model.  Many additions to the model have been developed to explain this discrepancy, including dark energy and modifications of general relativity on cosmological length scales.  Another suggestion is that the discrepancy results from large scale cosmological inhomogeneities, and arises from an unjustified fitting of supernova data to homogeneous models. It is well known that inhomogeneity in the cosmological matter distribution affects the redshifts and apparent magnitudes of supernovae \\cite{HG, CC, Lensing, V}, but could such effects be large enough to trick us into thinking that the Universe is accelerating when it is actually not? This issue has been vigorously debated in the literature; see Ref.\\ \\cite{Celerier} for a review.  One way to address the inhomogeneity issue in a fully nonlinear and relativistic fashion is with so-called ``Swiss cheese'' models \\cite{HW, Sugiura, CMY1, LTB}.  In these models, the homogeneous Friedmann-Robertson-Walker (FRW) metric is modified by the introduction of spherical regions of the spherically symmetric Lema\\^{i}tre-Tolman-Bondi (LTB) dust spacetime. In this paper we will restrict attention to such models where the FRW background is flat and matter dominated. The spherical inhomogeneities must be mass compensating and comoving in order to obtain a self consistent solution with the LTB and FRW regions evolving independently. The LTB regions can be of any size and can be located anywhere, so long as they do not overlap.  However, for a realistic model of the Universe, the sphere locations should be randomized, i.e.~they should not placed on a lattice of any sort.  In this paper, we show that the randomization issue is important and can have a dramatic effect on the model's predictions.  Marra, Kolb, and Matarrese (hereafter MKM) \\cite{MKM1, MKM2, Marra} use a Swiss cheese model where the LTB spheres correspond to underdense voids surrounded by overdense shells of matter.  They place these spheres on a regular lattice, and compute numerically the luminosity distance-redshift relation along a line passing through the spheres' centers. They show that the relation is modified enough to mimic a significant dark energy density.  We show that the results of MKM can be explained to $\\sim 10\\%$ accuracy with the formalism of standard weak field gravitational lensing.  The special line of sight that they consider leads to a large dimming because this is the line of sight with the lowest mean density, corresponding to less focusing.  We also show that if the void locations and lines of sight are properly randomized, then the mean effect is very small, in accordance with previous lensing studies \\cite{HW, Weinberg, Sereno2, KL}. Using a Monte Carlo simulation, we compute the statistical distribution of distance modulus shifts at fixed redshift in the more realistic randomized scenario.  This paper is organized as follows.  Section II reviews Swiss cheese models, focusing on the specific model constructed by MKM.  Section III shows that the apparent acceleration found by MKM is predominantly the result of gravitational lensing and of insufficient randomization of the void locations. In Sec. IV, we show that randomizing the void locations reduces the effect significantly, to the point that this model cannot explain the apparent acceleration seen in the supernova data.  In Sec. V, we provide our concluding remarks. ",
        "conclusions": "MKM found significant changes to the distance modulus of distant supernovae in a particular Swiss cheese model.  We have argued that the large size of their effect stems from the fact that their voids are on a regular lattice, with the line connecting observer and source passing straight through the void centers. Their result can be reproduced with $\\sim 10\\%$ accuracy using weak field gravitational lensing, and can be explained by noting that the integrated column density is lowest along lines of sight that pass through the centers of the voids. The lower column density leads to less focusing and thus dimmer supernovae.  We also showed that in a more realistic Swiss cheese model, the mean magnification due to gravitational lensing is very small. We simulated the distribution of distance modulus shifts obtained from a model where the light rays no longer pass only through the centers of the voids, but instead encounter them with random impact parameters.  For $N=5$ voids out to $z=1.8$, MKM obtained a coherent distance modulus shift of $0.34$ mag, whereas our randomized model gives a mean of $-0.003$ mag and a standard deviation of $0.11$ mag.  This mean is small and this standard deviation is comparable to the intrinsic scatter of Type Ia supernovae, and thus this particular model cannot explain the supernova data.  An analysis of the luminosity distance-redshift relation in a scenario with randomly-sized voids is the subject of future work.  While our results clarify the significance of the special setup underlying the specific model of MKM, they do not exclude the possibility that inhomogeneties may account for some of the apparent acceleration of the Universe. However, our work does suggest that coherent underdense structures on scales of order the horizon size generally are needed to obtain large apparent accelerations. Examples of models with such very large-scale structures are the spherical void models of Refs.\\ \\cite{Tomita, Alnes, VFW, Garfinkle, CMY2}. A general argument against such models is fine tuning: one would generically expect large anisotropies in the luminosity distance at a given redshift in any model with a large apparent acceleration, and the observational limits on anisotropy are small. Going beyond fine tuning, there are limits on large spherical void models from current observations of the cosmic microwave background spectrum \\cite{CMB} and low redshift Type Ia supernovae \\cite{Ia}, and further constraints could come from future observations of baryon acoustic oscillations \\cite{BAO} and time drifts of cosmological redshifts \\cite{zdot1, zdot2}.  We also note that in spherical void models, fitting luminosity distance data to spatially flat cosmological models would generally suggest the presence of dark energy.  If we allow more general fits to non flat models, then we will get a fit with non zero curvature and non zero dark energy, with the relative proportion depending on the details of the model. In linear theory, the shrinking mode of an LTB model is the same as the mode parameterized by the bang time function, and the growing mode is the mode parameterized by the energy function, which parameterizes curvature.  Therefore in a model with no important shrinking mode component, it is likely that the best fit will be have a higher proportion of curvature than dark energy."
    },
    "0808/0808.1049_arXiv.txt": {
        "abstract": "The effect of the QCD phase transition is studied for the mass-radius relation of compact stars and for hot and dense matter at a given proton fraction used as input in core-collapse supernova simulations. The phase transitions to the 2SC and CFL color superconducting phases lead to stable hybrid star configurations with a pure quark matter core. In supernova explosions quark matter could be easily produced due to $\\beta$-equilibrium, small proton fractions and nonvanishing temperatures. A low critical density for the phase transition to quark matter is compatible with present pulsar mass measurements. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.3085_arXiv.txt": {
        "abstract": "\\src\\ and \\srcxte\\ are considered the prototypical sources of the new class of High Mass X--ray Binaries, the Supergiant Fast X--ray Transients (SFXTs). These sources were observed  during bright outbursts on 2008 March 31 and 2008 April 8, respectively, thanks to an on-going monitoring campaign we are performing with {\\it Swift}, started in October 2007. Simultaneous observations with XRT and BAT allowed us to perform for the first time a broad band spectroscopy of their outbursts. The X--ray emission is well reproduced with absorbed cutoff powerlaws, similar to the typical spectral shape from accreting pulsars. \\src\\ shows a significantly harder spectrum during the bright flare (where a luminosity in excess of 10$^{36}$~erg~s$^{-1}$ is reached) than during the long-term low level flaring activity (10$^{33}$--10$^{34}$~erg~s$^{-1}$), while \\srcxte\\ displayed the same spectral shape, within the uncertainties, and a higher column density during the flare than in the low level activity. The light curves of these two SFXTs during the bright flare look similar to those observed during recent flares from other  two SFXTs, IGR~J11215--5952 and IGR~J16479--4514, reinforcing the connection among the members of this class of X--ray sources. ",
        "introduction": "}    \\src\\ and \\srcxte\\ are confirmed members of the new sub-class of High Mass X--ray Binaries, the Supergiant Fast X--ray Transients (SFXTs), whose members have been maily discovered with the INTEGRAL satellite (see e.g.\\ \\citealt{Sguera2005}). SFXTs are characterized by X--ray transient emission during short (a few hours long) flares reaching a few 10$^{36}$--10$^{37}$~erg~s$^{-1}$ and they are associated with blue supergiant companions (e.g.\\ \\citealt{Negueruela2005a} and  \\citealt{Smith2006aa}). The quiescent state in SFXTs have been observed only in a few sources and is characterized by a soft spectrum and X--ray luminosity at a level of 10$^{32}$~erg~s$^{-1}$, thus a very large dynamic range of about 1000--10000 has been observed. The short duration bright flares are part of a longer accretion phase at a lower level \\citep{Romano2007}. When not in outburst, these sources spend most of their lifetime in accretion at an intermediate (and flaring) level of X--ray luminosity, of 10$^{33}$--10$^{34}$~erg~s$^{-1}$ \\citep{Sidoli2008:sfxts_paperI}. The spectral properties are reminiscent of those of accreting pulsars, thus it is likely that several  members of the class are actually hosting neutron stars, although the spin period has been measured only in two SFXTs (AX~J1841.0--0536, \\citealt{Bamba2001}; IGR~J11215--5952, \\citealt{Swank2007:atel999}).  \\src\\ was discovered \\citep{Sunyaev2003} with IBIS/ISGRI on-board INTEGRAL on 2003 September 17  during a 2 hour flare reaching 160~mCrab (18--25~keV). During a $Chandra$ observation, both the quiescence level and the onset of an outburst was caught \\citep{zand2005}, observing a dynamic range as large as 10$^{4}$. The optical counterpart is an  O9Ib star \\citep{Pellizza2006} located at 3.6 kpc \\citep{Rahoui2008}. Several other bright flares have been observed with INTEGRAL in 2003, 2004 and 2005 (\\citealt{Grebenev2003:17544-2619}, \\citealt{Grebenev2004:17544-2619}, \\citealt{Sguera2006},  \\citealt{Walter2007},  \\citealt{Kuulkers2007}), with flare durations ranging from 0.5 to about 10 hours, reaching peak fluxes of 400~mCrab (20--40 keV). More recently, two new outbursts were detected with the {\\it Swift} satellite, on 2007 November 8  \\citep{Krimm2007:ATel1265} and on 2008 March 31 \\citep{Sidoli2008:atel1454}, 144 days apart. The flux at peak observed with {\\it Swift}/BAT was 165~mCrab (20--40~keV). The source was also active on  2007 September 21, with a fainter flaring emission up to 30--40 mCrab (20--60 keV), as observed with IBIS/ISGRI on-board INTEGRAL \\citep{Kuulkers2007:ATel1266}.   \\srcxte\\ was discovered with $RXTE$ after a short outburst in August 1997 \\citep{Smith1998}, and displayed a spectrum well fitted with a bremsstrahlung model with a temperature of $\\sim$22 keV, reaching a peak flux of 3.6$\\times$10$^{-9}$~erg cm$^{-2}$~s$^{-1}$   (2--25 keV). Later, several other short flares were observed with $RXTE$/PCA \\citep{Smith2006aa}. The optical counterpart is an  O8I star \\citep{Negueruela2006} located at 2.7 kpc \\citep{Rahoui2008}. Upper limits to the quiescent emission were placed with ASCA observations \\citep{Sakano2002} at a level of $<$1.1$\\times$10$^{-12}$~erg cm$^{-2}$~s$^{-1}$. Bright outbursts (up to 300~mCrab) were detected with IBIS/ISGRI in 2003 March, and 2004 March \\citep{Sguera2006}. Frequent flaring activity with INTEGRAL has been reported by \\citet{Walter2007}. Recently, it triggered the {\\it Swift} Burst Alert Telescope (BAT). An immediate slew allowed us to monitor the brightest part of a flare at soft energies \\citep{Romano2008:atel1466} with the {\\it Swift} X-ray Telescope (XRT). This outburst was also observed by the INTEGRAL/JEM-X monitor, which detected a flare starting 5 hours before the flares seen with {\\it Swift} \\citep{Chenevez2008}.  Here we report on the detailed analysis of the \\sw\\ data of two recent outbursts from these two prototypical SFXTs: the bright flares that occurred on 2008 March 31  \\citep{Sidoli2008:atel1454} from \\src\\ and  on 2008 April 8  from \\srcxte\\ \\citep{Romano2008:atel1466}. These observations are part of a monitoring campaign on four SFXTs with \\sw, which started on 2007 October 26. Results on the out-of-outburst emission of the earliest months of \\sw/XRT observations are reported in \\citet[][Paper I, see Fig.~\\ref{sfxts3:fig:lcv_history}]{Sidoli2008:sfxts_paperI}. The detailed analysis of the 2008 March 19 outburst of another SFXT in our monitoring program, IGR~J16479-4514, also caught by {\\it Swift}, is reported in \\citet[][Paper~II]{Romano2008:sfxts_paperII}. ",
        "conclusions": "}   Here we report on {\\it Swift} observations of \\src\\ and \\srcxte\\ during two bright flares, observed for the first time simultaneously in a broad energy range, from 0.3 to 50--60 keV (the highest energy where the spectroscopy is meaningful with BAT in these two sources). Indeed, before {\\it Swift}, the outbursts from these two SFXTs were mainly observed with INTEGRAL with the high energy detector (E$>$20 keV; e.g. \\citealt{Sguera2006}) or at softer energy bands (\\citealt{Smith2006aa}, \\citealt{zand2005}). The only SFXT previously observed simultaneously in a wide X--ray band was IGR~J16479--4514, during a flare caught with the {\\it Swift} satellite \\citep{Romano2008:sfxts_paperII}.   The X--ray spectroscopy shows that these two SFXTs, which are considered the prototypes of this new class of HMXBs, have different properties during the bright flares. \\src\\ is one order of magnitude less absorbed than \\srcxte, and displays a significantly flatter spectrum below 10 keV, with a XRT/WT spectrum well fitted  with a power-law with a photon index of 0.75$\\pm{0.11}$, compared with the \\srcxte\\ photon index, which lies in the range 1--2. The 1--10 keV spectral properties observed in \\src\\ during the flare are similar to what observed previously with $Chandra$ \\citep{zand2005}, where the absorbed powerlaw fit resulted in a photon index of 0.73$\\pm{0.13}$, a column density of (1.36$\\pm{0.22}$)$\\times$10$^{22}$~cm$^{-2}$, and a peak flux of $\\sim$3$\\times$10$^{-9}$~erg~cm$^{-2}$~s$^{-1}$.  The broad band analysis shows that \\src\\ displays a quite sharp cutoff at 18$\\pm{2}$~keV (when using the power law model with a high energy cut-off, {\\sc highecut} in Table~\\ref{tab:specigr}) or a well constrained temperature for the Comptonizing electrons (in the {\\sc comptt} model in XSPEC) at 4--5 keV. Instead, in \\srcxte, a single  power law (photon index of 2.2--2.5) can  describe the whole spectrum from soft to hard energies. Part of this difference could be explained by the much higher absorption towards the line of sight of \\srcxte, which does not allow to constrain well the low energy part of the power-law model.     \\begin{figure}[th!]% \\centerline{\\includegraphics[width=8.5cm,height=12cm,angle=0]{figure6.ps}} \\caption{Light curves of the outbursts of SFXTs followed by {\\it Swift}/XRT referred to their respective triggers. We show the 2005 outburst of IGR~J16479$-$4514 \\citep{Sidoli2008:sfxts_paperI}, which has a better coverage than the one observed in 2008 \\citep{Romano2008:sfxts_paperII}. The IGR~J11215$-$5952 light curve has an arbitrary start time, since the source did not trigger the BAT (the observations were obtained as a ToO; \\citet{Romano2007}. Note that where no data are plotted, no data were collected. For clarity, the time interval between two consecutive dashed vertical lines is one day. } [See the electronic edition of the Journal for a color version of this figure.] \\label{fig:comp} \\end{figure}     The observations we are reporting here are part of an on-going monitoring campaign of four SFXTs with {\\it Swift}  \\citep{Sidoli2008:sfxts_paperI}, which started on 2007 October 26. The two bright flares discussed here  are the first from these two SFXTs, since the start of the campaign, which could be simultaneously covered with both {\\it Swift} XRT and BAT. The results on the out-of-outburst X--ray emission (below 10 keV) have been reported by \\citet{Sidoli2008:sfxts_paperI}, where we find evidence that the accretion is still present, over long timescales of months, even  outside the bright outbursts. Both \\srcxte\\ and \\src\\ show evidence that they still accrete matter even outside the outbursts, at a much fainter (100--1000 times lower) level than during the flares, with still a large flux variability (at least one order of magnitude). A complete view of the different luminosity and spectral states of the monitored SFXTs will be clearer at the end of the campaign, but it is already possible to compare the average out-of-outburst emission properties with the spectra during the flares.  Regarding the 0.3--10 keV spectra (fitted with a simple absorbed power-law), \\srcxte\\ appears to be much more absorbed during the flare than during the out-of-outbust emission (see Fig.~\\ref{fig:contxte}), while the photon index is similar, within the large uncertainties \\citep{Sidoli2008:sfxts_paperI}. Similar changes in the absorbing column density of \\srcxte\\ have been observed before with $RXTE$/PCA and ASCA \\citep{Smith2006aa}, but during bright outbursts, where the $N_{\\rm H}$ ranged, from one bright flare to another, from 3 to 37$\\times10^{22}$~cm$^{-2}$. A $Chandra$ observation \\citep{Smith2006aa} displaying an unabsorbed 1--10~keV flux of $\\sim10^{-11}$~erg~cm$^{-2}$~s$^{-1}$, intermediate between the average out-of-outburst emission \\citep{Sidoli2008:sfxts_paperI} and the bright flare observed here, shows a hard power law spectrum with a photon index of 0.62$\\pm{0.23}$, absorbed with a column density    $N_{\\rm H}$=(4.2$\\pm{1.0}$)$\\times10^{22}$~cm$^{-2}$, which is compatible with that of the out-of-outburst emission. Thus, in \\srcxte, there does not seem to be a clear correlation between source intensity, spectral hardness and absorbing column density, to date.  Instead, \\src\\ shows a significantly harder spectrum during the flare, and a lower column density than during the out-of-outburst phase reported in \\citet{Sidoli2008:sfxts_paperI}, obtained summing together all the XRT data available from 2007 October 27 to 2008 February 28. During the out-of-outburst phase, the average observed flux was $\\sim$3$\\times10^{-12}$~erg~cm$^{-2}$~s$^{-1}$ (2--10 keV), the powerlaw photon index, $\\Gamma$, was 2.1$ ^{+0.6} _{-0.5}$, and the absorbing column density N$_{\\rm H}$=(3.2 $^{+1.2} _{-0.9}$)$\\times10^{22}$~cm$^{-2}$ \\citep{Sidoli2008:sfxts_paperI}. We  also compared the bright flare spectroscopy with the spectrum extracted from one of the observations obtained a few days before the bright flare from \\src\\ (dashed contours in Fig.~\\ref{fig:cont}). A hardening of the  \\src\\ spectrum during the flaring activity is evident. A similar behaviour was already suggested from the analysis of different XMM-Newton observations during a low level flaring activity \\citep{Gonzalez2004}, and from the analysis of the 2004 $Chandra$ observation \\citep{zand2005}. A proper comparison with the INTEGRAL results of a few outbursts from \\src\\ reported by Sguera et al. (2006) cannot be done since the energy range with INTEGRAL was limited to 20--60~keV. These authors fitted the 20--60~keV spectrum with a thermal bremsstrahlung, which is clearly not adequate to describe our XRT+BAT spectrum (reduced $\\chi^{2}_{\\nu}$ of 1.7, for 159 dof).   Different mechanisms have been proposed to explain the bright and short duration flaring activity in this new class of sources. Some models are related to the structure of the supergiant companion wind, involving spherically simmetric clumpy winds (see e.g. \\citealt{zand2005}, \\citealt{Negueruela2008}) or  anisotropic winds \\citep{Sidoli2007}; other models involve the interaction of the inflowing wind with the neutron star magnetosphere (see e.g. Bozzo et al. 2008). Sidoli et al. (2007) explain the outbursts as being due to enhanced accretion onto the neutron star when it crosses, moving along the orbit, an equatorial wind disk component from the supergiant  companion. Depending on the thickness and truncation of this supposed disk wind and on its inclination with respect to the orbital plane of the binary system, the compact object will cross once or twice in a periodic or quasi-periodic manner the disk, undergoing outbursts. In the framework of this model, the geometry, the structure of this disk wind and its inclination with respect to the line of sight could explain the variability in the local absorbing column density, even during different outbursts (as observed several times in \\srcxte) and compared with the low level activity. A lower column density during the out-of-outburst activity could be due to the fact the source is completely outside the denser equatorial wind from the companion. In the spherically symmetric clumpy winds model, the difference in the observed column density could be due to the accreting dense clumps. On the other hand, in this case the clump matter should remain neutral also in proximity of the neutron star during bright flares. We think it is more likely that the absorbing column density is not related with a neutral accreting matter, but with other clumps or wind structures located probably farther away from the compact source.  In Fig.~\\ref{fig:comp} we compare the light curves during bright flares from four SFXTs, all observed with {\\it Swift}: the two reported here from \\srcigr\\ and  \\srcxte, together with the one observed from IGR~J11215--5952 \\citep{Romano2007} and from IGR~J16479-4514 \\citep{Romano2008:sfxts_paperII}. All light curves during bright flares look similar, although they were observed with a different sampling. We postpone a more quantitative comparison between the four SFXTs (duty cycle, light curve rise time and decay times) to a final paper at the end of the on-going observing campaign. In any case, it is already evident that the behaviour of this sample of SFXTs in outburst is similar, and that their bright emission extends for more than a few hours, contrary to what originally thought at the time of the discovery of this new class of sources \\citep[e.g., ][]{Sguera2005}.  The wide band spectra during outbursts display high energy cut-offs (assuming the model with a power-law modified at high energy by a cutoft, {\\sc highecut} in XSPEC), although differently constrained in the two sources: in \\srcigr\\ it is at 18$\\pm{2}$~keV, in \\srcxte\\ it lies below 13~keV. These cut-off ranges are fully consistent with a neutron star magnetic field, B, of about a 2--3$\\times$10$^{12}$~G in the case of \\src\\ and of less than about 2$\\times$10$^{12}$~G for \\srcxte\\ \\citep{Coburn2002}. These estimates, although not based on a direct measurement of the magnetic field (which would be possible only in the case of detection of cyclotron lines), are already difficult to explain in the framework of the magnetar model recently proposed by \\citet{Bozzo2008}, where the magnetic field is at a level of 10$^{14}$~G. The same is true for other two SFXTs, IGR~J11215--5952 \\citep{Sidoli2007} and IGR~J16479--4514 \\citep{Romano2008:sfxts_paperII}.   {\\it Facilities:} \\facility{{\\it Swift}}."
    },
    "0808/0808.1878_arXiv.txt": {
        "abstract": "s{The NuMoon project uses the Westerbork Synthesis Radio Telescope to search for short radio pulses from the Moon. These pulses are created when an ultra high energy cosmic ray or neutrino initiates a particle cascade inside the Moon's regolith. The cascade has a negative charge excess and moves faster than the local speed of light, which causes coherent Cherenkov radiation to be emitted. With 100 hours of data, a limit on the neutrino flux can be set that is an order of magnitude better than the current one (based on FORTE). We present an analysis of the first 10 hours of data. } ",
        "introduction": "Above the Greisen-Zatsepin-Kuzmin (GZK) energy of $6\\cdot 10^{19}$~eV, cosmic rays (CRs) can interact with the cosmic microwave background photons, losing energy and producing pions when traversing distances of the order of 10 Mpc \\cite{g66,zk66}. Recent results of the Pierre Auger Observatory have confirmed a steepening in the cosmic ray spectrum at the GZK energy \\cite{Y07}. The pions will produce ultrahigh energy (UHE) neutrinos through weak decay. Observing these neutrinos is of great scientific interest as their arrival direction points back to sources at distances larger than 10 Mpc, in contrast to charged particles that are deflected in (extra-)galactic magnetic fields. In addition, while cosmic rays from faraway sources pile up at the GZK energy, information about the cosmic ray spectrum at the source is conserved in the GZK neutrino flux. Other possible sources of UHE neutrinos are decaying supermassive particles, such as magnetic monopoles or topological defects. This class of models is refered to as top-down (TD) models (see for example Stanev \\cite{s04} for a review).  Because of their small interaction cross section and low flux, the detection of cosmic neutrinos calls for extremely large detectors. Assuming the Waxman-Bahcall flux \\cite{wb01}, even at GeV energies the flux is not higher than a few tens of neutrinos per km$^2$ per year. Kilometer-scale detectors are not easily built but can be found in nature. For example, interaction of neutrinos in ice or water can be detected by the Cherenkov light produced by the lepton track or cascade. The now half finished IceCube detector \\cite{icecube} will cover a km$^3$ volume of South Pole ice with PMTs, while Antares \\cite{antares} and its successor KM3NET \\cite{KM3NET} exploit the same technique in the Mediterranean sea. Even larger volumes can be covered by observing large detector masses from a distance. The ANITA balloon mission \\cite{anita} monitors an area of a million km$^2$ of South Pole ice from an altitude of $\\sim 37$~km and the FORTE satellite \\cite{forte} can pick up radio signals coming from the Greenland ice mass. Alternatively, cosmic ray experiments like the Pierre Auger Observatory can possibly distinguish cosmic ray induced air showers from neutrino induced cascades at very high declinations where the atmosphere is thickest and only neutrinos can interact close to the detector.  For an even larger detector volume one can turn to the Moon. The negative charge excess of a particle shower inside a dense medium will cause the emission of coherent Cherenkov radiation in a process known as the Askaryan effect \\cite{a62}. This emission mechanism has been experimentally verified at accelerators \\cite{s01,g00} and extensive calculations have been performed to quantify the effect \\cite{zhs92,az97}. The idea to observe this type of emission from the Moon with radio telescopes was first proposed by Dagkesamanskii and Zheleznyk \\cite{dz89} and the first experimental endeavours in this direction were carried out with the Parkes telescope \\cite{parkes} and at Goldstone (GLUE) \\cite{glue}.  The NuMoon project uses the Westerbork Synthesis Radio Telescope (WSRT) to watch for the same flashes but at lower frequency, which has the distinct advantage that radio pulses have a much higher chance of reaching the observer, which will be explained in the next section. ",
        "conclusions": "We observe at a frequency window that offers an optimal sensitivity to lunar pulses. At the same time, for these low frequencies the effects of uncertainties concerning the lunar surface and interior are small. Because of the large spread in emission angle, we expect no systematic effect from surface irregularities. The detection efficiency is also largely independent from details in the structure of the (sub-)regolith \\cite{scholten}.  A possible complication is the limited dynamic range of the WSRT data. The individual dishes give a 2-bit signal, so the total dynamic range is $3N+1$ for $N$ dishes. When there is a lot of RFI background, as much as 90\\% of the total received power can be in narrow RFI lines. After background reduction the dynamic range is suppressed and detection of pulses may become impossible. Another effect of the limited dynamic range is that large pulses will be cut off to the maximum allowed value, potentially lowering the signal-to-noise. Here, the atmospheric dispersion contributes positively by spreading the power over more timebins.  Currently we are simulating the detection efficiency for various radio backgrounds and atmospheric conditions. The error on the TEC value, used for dedispersion, is also taken into account."
    },
    "0808/0808.0279_arXiv.txt": {
        "abstract": "{Using the atmosphere as a detector volume, Imaging Air Cherenkov Telescopes (IACTs) depend highly on the properties and the condition of the air mass above the telescope. On the Canary Island of La Palma, where the Major Atmospheric Gamma-ray Imaging Cherenkov telescope (MAGIC) is situated, the Saharan Air Layer (SAL) can cause strong atmospheric absorption affecting the data quality and resulting in a reduced gamma flux.}{To correlate IACT data with other measurements, e.g. long-term monitoring or Multi-Wavelength (MWL) studies, an accurate flux determination is mandatory. Therefore, a method to correct the data for the effect of the SAL is required.}{Three different measurements of the atmospheric absorption are performed on La Palma. From the determined transmission, a correction factor is calculated and applied to the MAGIC data.} {The different transmission measurements from optical and IACT data provide comparable results. MAGIC data of PG\\,1553+113, taken during a MWL campaign in July 2006, were analyzed using the presented method, providing a corrected flux measurement for the study of the spectral energy distribution of the source.}{} ",
        "introduction": "The Earth's atmosphere plays an important role in all ground-based observations because it is traversed by light from all objects observed. Scattering and absorption of light by the atmosphere is variable and reduces the flux measured at the telescope. On the Canary Island of La Palma, very strong atmospheric absorption can occur, in case the meteorologic phenomenon Saharan Air Layer (SAL) takes place. A warm air mass, consisting of mineral dust, travels from the Sahara westward; this becomes the so-called SAL, when it is undercut by a cool and moist air layer from the sea. Located between 1.5\\,km and 5.5\\,km above sea level (a.s.l.), the SAL, also known as Calima, can extend above the North Atlantic Ocean as far as the United States and the western Caribbean Sea \\citep{sal}.  Variable absorption can cause apparent flux variations, and the absolute flux level is in all cases reduced. To to correct for this reduction, the atmospheric absorption is measured. On La Palma, three measurements are available. The Carlsberg Meridian Telescope (CMT) provides nightly values of extinction. From the optical data of the KVA and also from the MAGIC data itself, the atmospheric transmission can be determined. The results of these measurements are discussed and compared for the nights between July 15${\\rm ^{th}}$ 2006 and July 28${\\rm ^{th}}$ 2006.  A correction method was developed for the MAGIC telescope: From the absorption measurements, a factor is calculated and used in the analysis to correct for the effect of SAL. ",
        "conclusions": "The different measurements of the atmospheric transmission from stellar photometry (KVA, CMT) and IACT data (MAGIC) agree well for nights with stable extinction. Since more than 90\\,\\% of the shower light originates above the SAL, the measurements can be used directly to correct the IACT data. By applying a correction factor in the calibration, the data of PG\\,1553+113 taken during the MWL campaign in July 2006 for example were corrected successfully. This shows that the loss of observation time due to atmospheric absorption can be reduced.  For additional checks, affected data of a steady source are required. When such data become available within scheduled MAGIC observations, the presented method can be tested further.  The correction should not be applied blindly to data affected by atmospheric absorption of another kind. For those, the vertical distribution of the atmospheric conditions should be studied e.g.\\ with a Lidar (light detection and ranging), i.e.\\ a device shooting a laser beam into the atmosphere and measuring the backscattered light. The possible use of information provided by a Lidar, however, still has to be investigated.  In cases of strong absorption, i.e.\\ above 40\\,\\%, the method is limited by the fluctuations in the night-sky background light and the effect cannot be corrected completely."
    },
    "0808/0808.2536_arXiv.txt": {
        "abstract": "The key contribution of the discovery of nuclear-powered pulsations from the accretion-powered millisecond pulsars (AMPs) has been the establishment of burst oscillation frequency as a reliable proxy for stellar spin rate.  This has doubled the sample of rapidly-rotating accreting neutron stars and revealed the unexpected absence of any stars rotating near the break-up limit. The resulting `braking problem' is now a major concern for theorists, particularly given the possible role of gravitational wave emission in limiting spin.  This, however, is not the only area where burst oscillations from the AMPs are having an impact.  Burst oscillation timing is developing into a promising technique for verifying the level of spin variability in the AMPs (a topic of considerable debate).  These sources also provide unique input to our efforts to understand the still-elusive burst oscillation mechanism. This is because they are the only stars where we can reliably gauge the role of uneven fuel deposition and, of course, the magnetic field. ",
        "introduction": "`History', as George Orwell once noted, `is written by the winners' \\citep{orw44} - or, in this case, by the workshop hosts.  When we gathered in Amsterdam in April 2008, it was ostensibly to celebrate ten years since the discovery of the first Accreting Millisecond X-ray Pulsar (AMXP).  We were, however, a full two years too late.  For the first AMXP was not SAX J1808.4-3658 \\citep{wij98}, but rather the far less well-known 4U 1728-34 \\citep{str96b}.  How on earth, you might ask, could such a slip go unnoticed?  The trick, of course, lies in the terminology.  Most astronomers (the author included) tend to think of the AMXPs as comprising only the {\\it accretion-powered} millisecond pulsars (AMPs), forgetting the equally large class of {\\it nuclear-powered} millisecond pulsars (NMPs) - the burst oscillation sources.  Most of this volume focuses on the AMPs, where persistent pulsations are generated as accreting material is channeled by the magnetic field onto magnetic polar caps that are offset from the rotational poles.  The NMPs, by contrast, show pulsations during Type I X-ray bursts (thermonuclear explosions on the stellar surface caused by rapid unstable burning of accreted material). The cause of the brightness asymmetry in the NMPs remains an open question \\citep{str06, gal08}, and to do full justice to NMP phenomenology would merit a much longer discussion.  In this article, however, I will focus on the small set of NMPs that are also AMPs.  These rare objects provide a unique insight into many current problems in neutron star astrophysics because, as suggested by my title, they are lighthouses with two different light sources.  The accretion-powered pulsations tell us how the material arrives on the stellar surface, while the nuclear-powered pulsations tell us what happens once it gets there.  Section \\ref{data} provides a brief overview of the relevant observational results.  The bulk of the review focuses, however, on the astrophysical questions where these sources have made or are making a major contribution to our understanding.  These include the spin distribution of AMXPs, torque modeling, and the burst oscillation mechanism. ",
        "conclusions": "\\label{conc}  The discovery of nuclear-powered pulsations from the AMPs has cemented the link between burst oscillation frequency and spin frequency. In addition to confirming the absence of rapid rotators (now a major problem for evolutionary models), this has imposed the strongest single constraint on candidate burst oscillation mechanisms. Despite this huge clue, the mechanism remains elusive: but continuing analysis of the AMPs is providing tantalising evidence that is driving the development of new theoretical models."
    },
    "0808/0808.2193_arXiv.txt": {
        "abstract": "We report on our early photometric and spectroscopic observations of the extremely luminous Type II supernova (SN) 2008es. With an observed peak optical magnitude of $m_V = 17.8$ and at a redshift $z = 0.213$, \\sn\\ had a peak absolute magnitude of $M_V$ = $-22.3$, making it the second most luminous SN ever observed. The photometric evolution of \\sn\\ exhibits a fast decline rate ($\\sim$0.042 mag d$^{-1}$), similar to the extremely luminous Type II-L SN~2005ap.  We show that \\sn\\ spectroscopically resembles the luminous Type II-L SN~1979C. Although the spectra of \\sn\\ lack the narrow and intermediate-width line emission typically associated with the interaction of a SN with the circumstellar medium of its progenitor star, we argue that the extreme luminosity of \\sn\\ is powered via strong interaction with a dense, optically thick circumstellar medium. The integrated bolometric luminosity of \\sn\\ yields a total radiated energy at ultraviolet and optical wavelengths of $\\ga$10$^{51}$ ergs.  Finally, we examine the apparently anomalous rate at which the Texas Supernova Search has discovered rare kinds of supernovae, including the five most luminous supernovae observed to date, and find that their results are consistent with those of other modern SN searches. ",
        "introduction": "Wide-field synoptic optical imaging surveys are continuing to probe the parameter space of time-variable phenomena with increasing depth and temporal coverage \\citepeg{2004ApJ...611..418B,2006MNRAS.371.1681M,2008MNRAS.386..887B}, unveiling a variety of transients ranging from the common \\citep{2008ApJ...682.1205R} to the unexplained (e.g., \\citealt{bdt+08}). Untargeted (``blind'') synoptic wide-field imaging surveys, such as the Texas Supernova Search (TSS; \\citealt{quimbyTSS}) conducted with the ROTSE-III 0.45-m telescope \\citep{akerlof03}, have uncovered a large number of rare transients, including the four most luminous supernovae (SNe) observed to date: SN~2005ap \\citep{quimby05ap}, SN~2008am \\citep{ATEL.1389}, SN~2006gy \\citep{ofek06gy, smith07-2006gy, smith08-2006gy}, and SN~2006tf \\citep{smith06tf}.  Observations of these very luminous events are starting to allow the detailed physical study of the extrema in core-collapse SNe. They appear to be powered in part by their interaction with a highly dense circumstellar medium (CSM; see \\citealt{smith06tf}, and references therein), though other possibilities have been advanced. Clearly, the discovery of more such events would allow an exploration of the variety of the phenomenology as related to the diversity of progenitors and CSM.  Recently, the TSS discovered yet another luminous transient on 2008 Apr.  26.23 (UT dates are used throughout this paper), which they suggested was a variable active galactic nucleus at a redshift $z = 1.02$ \\citep{ATEL.1515}. \\citet{ATEL.1524} then hypothesized that the transient was a flare from the tidal disruption of a star by a supermassive black hole. \\citet{ATEL.1576} first identified \\sn\\ as potentially an extremely luminous Type II SN (see also \\citealt{2008ATel.1578....1G}), and we later definitively confirmed this with further spectroscopic observations \\citep{ATEL.1644}; the event was assigned the name \\sn\\ by the IAU \\citep{2008CBET.1462....1C}. It is located at $\\alpha$ = 11$^h$56$^m$49.06$^s$, $\\delta$ = +54$^o$27$\\arcmin$24.77$\\arcsec$ (J2000.0).  Here we present our analysis of SN 2008es, which is classified as a Type II-Linear (II-L) SN based on the observed linear (in mag) decline in the photometric light curve \\citep{barbon79,doggett85}. At $z$ = 0.213, \\sn\\ has a peak optical magnitude of $M_V = -22.3$, among SNe second only to SN~2005ap. Aside from the extreme luminosity, \\sn\\ is of great interest since detailed ultraviolet (UV) through infrared (IR) observations provide a unique opportunity to study the mass-loss properties of an evolved post-main sequence massive star via its interaction with the surrounding dense CSM. A similar analysis of \\sn\\ has been presented by \\citet{gezari08}.  The outline of this paper is as follows. We present our observations in $\\S$2, and the photometric and spectroscopic analyses of this and public (NASA) data in $\\S$3 and $\\S$4, respectively. A discussion is given in $\\S$5, and our conclusions are summarized in $\\S$6. Throughout this paper we adopt a concordance cosmology of $H_0$ = 70 \\kms\\ Mpc$^{-1}$, $\\Omega_{\\rm M} = 0.3$, and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "We have reported on our early observations of \\sn, which at a peak optical magnitude of $M_V = -22.3$ is the second most luminous SN ever observed.  We argue that the extreme luminosity of this SN was likely powered via strong interaction with a dense CSM, and that the steep decline in the light curve, 0.042 mag d$^{-1}$, indicates that the radioactive decay of $^{56}$Co is likely not the dominant source of energy for this SN. Integration of the bolometric light curve of \\sn\\ yields a total radiated energy output of $\\ga$ 10$^{51}$ ergs.  The optical spectra of \\sn\\ resemble those of the luminous SN~1979C, but with an unexplained increase in the velocity of the H$\\beta$ absorption minimum over time.  We also examined the rate of discovery of extremely luminous SNe by the Texas Supernova Search and find that their discovery of the five most luminous observed SNe in the past four years is probably not a fluke; several more such detections are expected in the coming years.  Finally, what behavior can we expect from \\sn\\ at late times, roughly 1 yr or more after discovery?  Regardless of whether the peak luminosity was powered by radioactive decay or optically thick CSM interaction (see \\citealt{smith08-2006gy}), a SN can be powered by strong CSM interaction at late times if the progenitor had a sufficiently high mass-loss rate in the centuries before exploding. We have seen examples of both: SN~2006tf had strong H$\\alpha$ emission indicative of ongoing CSM interaction at late times \\citep{smith06tf}, whereas SN~2006gy did not \\citep{smith08-2006gy}.  SN~1979C was less luminous at peak than those two SNe, but it has been studied for three decades because its late-time CSM interaction is powering ongoing emission in the radio, optical, and X-rays \\citep{weiler81,fesen93, immler98}. With such an extraordinarily high peak luminosity, a late-time IR echo such as that seen in SN~2006gy \\citep{smith08-2006gy} is also likely if \\sn\\ has dust waiting at a radius of $\\sim$0.3 pc. Alternatively, if \\sn\\ were powered in whole or in part by radioactivity, a large mass of $^{56}$Ni should be evident in the late-time decline rate."
    },
    "0808/0808.2470_arXiv.txt": {
        "abstract": "We present observations of \\caii, \\znii\\, and \\crii\\ absorption lines in 16 damped Lyman alpha systems (DLA) and six subDLAs at redshifts $0.6 < z_{\\rm abs} < 1.3$, obtained for the dual purposes of: (i) clarifying the relationship between DLAs and absorption systems selected from their strong \\caii\\ lines, and (ii) increasing the still limited sample of Zn and Cr abundance determinations in this redshift range. We find only partial overlap between current samples of intermediate redshift DLAs (which are drawn from magnitude limited surveys) and strong \\caii\\ absorbers: approximately 25 per cent of known DLAs at these redshifts have an associated Ca\\,{\\sc ii}\\,$\\lambda 3935$ line with a rest-frame equivalent width greater than 0.35\\,\\AA, the threshold of the Sloan Digital Sky Survey sample assembled by Wild and her collaborators.  The lack of the strongest \\caii\\ systems (with equivalent widths greater than 0.5\\,\\AA) is consistent with these authors' conclusion that such absorbers are often missed in current DLA surveys because they redden and dim the light of the background QSOs.  We rule out the suggestion that strong \\caii\\ absorption is associated exclusively with the highest column density DLAs. Furthermore, we find no correlation between the strength of the \\caii\\ lines and either the metallicity or degree of depletion of refractory elements, although the strongest \\caii\\ absorber in our sample is also the most metal-rich DLA yet discovered, with [Zn/H] $\\simeq$ solar. We conclude that a complex mix of parameters must determine the strengths of the \\caii\\ lines, including the density of particles and ultraviolet photons in the interstellar media of the galaxies hosting the DLAs.  We find tentative evidence (given the small size of our sample) that strong \\caii\\ systems may preferentially sample regions of high gas density, perhaps akin to the DLAs exhibiting molecular hydrogen absorption at redshifts $z > 2$.  If this connection is confirmed, strong \\caii\\ absorbers would trace possibly metal-rich, H$_2$-bearing columns of cool, dense gas at distances up to tens of kpc from normal galaxies. ",
        "introduction": "\\label{Sec:Int} Quasar absorption lines are a powerful tool for the study of gaseous galactic structures over cosmic time, providing clues to the nature and evolution of galaxies and galactic halos.  The large columns of neutral hydrogen gas traced by damped Lyman alpha (DLA) and subDLA systems (conventionally defined to have \\hi\\ column densities $N$(H\\,{\\sc i}) $\\ge 10^{20.3}$\\,\\acm\\ and $10^{19} \\le N$(H\\,{\\sc i})\\,$ < 10^{20.3}$\\,\\acm\\ respectively) are of particular interest (Rao 2005; Wolfe, Gawiser, \\& Prochaska 2005), as they account for most of the neutral gas in the Universe and trace galaxies which are often difficult to detect directly. Thus, appropriate to their importance, much effort has been expended to understand the nature and incidence of the gas traced by DLAs and subDLAs.  While DLAs do contain metals (Wolfe et al.\\ 1986; Meyer, Welty, \\& York 1989; Pettini, Boksenberg, \\& Hunstead 1990) and ionised gas (Fox et al. 2007a), they are characterized by low metallicities (Kulkarni et al. 2005; Akerman et al. 2005; Prochaska et al. 2007) and large neutral fractions (Viegas 1995; Vladilo et al. 2001). Little, if any, evolution is detected in the co-moving mass density of \\hi, $\\Omega_{\\rm H\\,I}$, from $z \\simeq 4$ to $z \\approx 0.5$ (Rao, Turnshek,  \\& Nestor, 2006;  Lah et al.\\ 2007; although see Prochaska, Herbert-Fort, \\& Wolfe 2005) and of metals (Kulkarni et al. 2005). This is perhaps surprising, considering the putative role of DLAs as the reservoir of fuel for star formation. In fact, recent work (Wolfe \\& Chen 2006; Wild, Hewett, \\& Pettini 2007) has shown that \\textit{in situ} star formation in  DLAs is significantly less than expected on the basis of the present-day Kennicutt-Schmidt law (Kennicutt 1998), considering their large \\hi\\ surface densities. One possible explanation for the apparent lack of evolution in $\\Omega_{\\rm H\\,I}$ is that a significant fraction of DLAs with high star formation rates or metallicities  are missed from traditional magnitude limited optical QSO surveys due to chromatic extinction from large columns of dust. This possibility, however, is not supported by surveys of DLAs in radio-selected QSOs (Ellison et al. 2001; Jorgenson et al. 2006).  SubDLAs, on the other hand, contribute a smaller fraction of the total \\hi\\ in the Universe. However, they possess larger fractions of ionised gas than DLAs, and thus the total amount of hydrogen (\\hi\\ plus \\hii) in subDLAs may be comparable to that of DLAs (Fox et al. 2007b). Furthermore, subDLAs appear to have higher metallicities than DLAs (Kulkarni et al. 2007) and thus may be a large repository of the metals produced at high redshifts (see Pettini 2006). While DLAs and subDLAs both contain large amounts of hydrogen, they are likely to arise in different regions within galaxies and to trace different environments and properties.  Strong metal lines of Mg\\,{\\sc ii}\\,$\\lambda\\lambda 2796, 2804$ and \\feii\\,$\\lambda\\lambda 2587, 2600$ are ubiquitous in (sub)DLA systems. Their large oscillator strengths and near-ultraviolet (UV) rest wavelengths make them ideal for tracing large columns of neutral gas at redshifts $z \\la 1.7$, where the Ly$\\alpha$ line falls below the atmospheric cut-off at 3200\\,\\AA\\ and therefore requires observations from space. In this context, systems selected via strong Ca\\,{\\sc ii}\\,$\\lambda\\lambda 3935, 3970$ absorption have recently garnered attention (Wild \\& Hewett 2005). Extensive literature exists on the Ca\\,{\\sc ii} doublet in the interstellar media of the Milky Way and nearby galaxies, since these lines can be studied from the ground even at redshift $z=0$ (e.g. Marshall \\& Hobbs 1972; Welty, Morton, \\& Hobbs 1996 and references therein). This body of work has shown that in local interstellar environments Ca is a strongly refractory element, exhibiting some of the highest depletion factors (Savage \\& Sembach 1996). Furthermore, since the ionisation potential of Ca\\,{\\sc ii} is lower than that of hydrogen, Ca\\,{\\sc ii} is not the dominant ionisation stage of Ca in H\\,{\\sc i} regions, unlike Mg\\,{\\sc ii} and Fe\\,{\\sc ii}. Rather, the balance between Ca\\,{\\sc ii} and Ca\\,{\\sc iii} depends on the details of the physical conditions in the gas, primarily on temperature and on the densities of particles and far-UV photons. Therefore, while all DLAs exhibit strong \\mgii\\ and \\feii\\ lines, the same is not true of \\caii.  Wild, Hewett, \\& Pettini (2006; WHP06) recently conducted a survey for strong \\caii\\ systems, with $\\lambda 3935$ rest equivalent width \\wca\\,$\\geq 0.35$\\,\\AA, in QSO spectra from the Sloan Digital Sky Survey (SDSS). Based on the strengths of Zn\\,{\\sc ii}\\ and Cr\\,{\\sc ii}\\ lines, WHP06 concluded that, on average, strong Ca\\,{\\sc ii} absorbers have \\nhi\\ values above the DLA limit. They also showed that these absorbers are, on average, dustier than typical DLAs, introducing a reddening of $\\langle E(B-V)\\rangle \\ga 0.1$\\,magnitudes in the spectra of the background QSOs, compared to  $\\langle E(B-V)\\rangle \\la 0.02$ for the general population of SDSS DLAs (WHP06; Vladilo, Prochaska, \\& Wolfe 2008). This level of obscuration of the background QSOs is significant, as it causes flux-limited and/or colour-selected surveys to underestimate the numbers of strong Ca\\,{\\sc ii} systems. If \\caii-strong DLAs are being missed, and if they are preferentially metal-rich relative to DLAs with weaker \\caii\\ absorption, then this class of absorption system may contain a previously overlooked repository of both neutral gas and metals. In any case, if \\caii\\ systems select a population of absorbers with unique properties, they could hold the key to a better understanding of the detailed physical properties of DLAs and subDLAs.   Since the survey by WHP06, efforts have been devoted to better understanding the link between strong Ca\\,{\\sc ii} systems and galaxies. Using $K$-band imaging, Hewett \\& Wild (2007) reported an excess of luminous galaxies close to the sightlines to 30 QSOs with intervening Ca\\,{\\sc ii} systems at redshifts $0.7 < z < 1.2$. The galaxies appear to exhibit a luminosity-dependent cross section for \\caii\\ absorption and have a mean impact parameter of $\\approx 25$\\,kpc. Together with the incidence of absorbers from WHP06, this result implies a halo volume filling factor of $\\sim 10$ per cent out to $\\approx 35$\\,kpc from the absorbing galaxies. At lower redshifts, Zych et al. (2007) were able to associate luminous, metal-rich, star-forming spiral galaxies with four out of five strong Ca\\,{\\sc ii} absorbers at $z < 0.5$.  Although these studies represent progress in uncovering the nature of \\caii\\ absorbers, ignorance of \\nhi\\ in such systems limits the direct insight they provide into DLAs and subDLAs in general. While it seems very likely that \\caii\\ absorbers select large columns of neutral gas, direct measurements of \\nhi\\ are currently available for very few \\caii\\ systems. The reason is simple: by the time the SDSS data highlighted the potential importance of absorption systems selected via Ca\\,{\\sc ii}, no space-borne UV spectrograph was in operation to measure the associated column densities of H\\,{\\sc i}. This unfortunate state of affairs is about to change with the forthcoming installation of the Cosmic Origins Spectrograph (COS) on the refurbished \\textit{Hubble Space Telescope} (\\textit{HST}). In the meantime, the most immediate approach to this problem is to perform the complementary experiment of measuring the strength of \\caii\\ absorption in known DLAs at redshifts $z \\simlt 1.3$, where the \\caii\\ doublet lines are still accessible with optical spectrographs.  The samples of confirmed DLAs at these intermediate redshifts have increased considerably in recent years thanks to dedicated \\textit{HST} surveys (see Rao, Turnshek \\& Nestor 2006; RTN06). In this paper, we present new observations of QSOs from the RTN06 compilation designed to: \\textit{(i)} measure the equivalent widths of the Ca\\,{\\sc ii}\\,$\\lambda \\lambda 3935, 3970$ doublet lines in order to assess the overlap of known DLAs (and a few subDLAs) with the strong Ca\\,{\\sc ii} absorber sample of WHP06; and \\textit{(ii)} simultaneously determine the metallicities and dust content of intermediate redshift DLAs via observations of the associated Zn\\,{\\sc ii} and Cr\\,{\\sc ii} absorption lines (Pettini et al. 1990), thereby adding to the still somewhat limited statistics of these measures at $z < 1.5$ (Akerman et al. 2005; Kulkarni et al. 2007). In Section~\\ref{Sec:Obs} we describe our observations and equivalent width measurements, while Section~\\ref{sec:abundances} deals with the derivation of column densities and element abundances. The main results on the \\caii\\ properties of known DLAs at intermediate redshifts are presented in Section~\\ref{sec:caii_in_dlas}. We summarise our principal findings and conclusions in Section~\\ref{Sec:Sum}. ",
        "conclusions": "\\label{Sec:Sum} We have used the WHT telescope on La Palma to record at intermediate resolution the optical spectra of QSOs known to lie behind DLAs and subDLAs in the redshift range $0.6 < z_{\\rm abs} < 1.3$. Our principal aims were twofold: (i) to measure the strength of the Ca\\,{\\sc ii}\\,$\\lambda \\lambda 3935, 3970$ absorption lines in known DLAs, with a view to establishing to what extent DLAs and strong Ca\\,{\\sc ii} absorbers overlap; and (ii) to measure the \\znii\\,$\\lambda\\lambda 2026, 2063$ and \\crii\\,$\\lambda\\lambda 2056, 2062, 2066$ multiplets in the (sub)DLAs, thereby increasing the number of metallicity and depletion determinations in (sub)DLAs at low-redshift. By design, these complementary measurements allow us to test whether the strength of Ca\\,{\\sc ii} absorption is related to the metallicity of the gas, as measured by the abundance of Zn, and/or to the degree of dust depletion of refractory elements (both Ca and Cr are readily incorporated onto dust, whereas Zn is not). From consideration of such measurements in a sample of 16 DLAs and six subDLAs, augmented by similar data for three DLAs and four subDLAs from published studies, we reach the following conclusions.    \\begin{itemize}  \\item[1.] The new measurements presented here strengthen previous conclusions on the metallicities and dust depletions in DLAs at $z < 1.3$. We point out that, after accounting for the fact that at these redshifts DLAs have mostly been selected from Mg\\,{\\sc ii}-strong systems -- and may thus be biased in favour of metal-rich absorbers -- the already weak redshift evolution of the mean  DLA metallicity is reduced further: [$\\langle {\\rm Zn/H}\\rangle$] is only a factor of about two greater at  $0.6 < z < 1.3$ than at higher redshifts. We also confirm the dependence of dust depletion on metallicity uncovered by earlier studies.   \\item[2.] Most DLAs exhibit detectable Ca\\,{\\sc ii}\\,$\\lambda 3935$ absorption in our spectra.  Our new data indicate that selecting absorption systems above the threshold  \\wca\\,$ = 0.15$\\,\\AA\\ isolates systems with \\nhi\\,$\\ga 2 \\times 10^{20}$\\,cm$^{-2}$ (although with some contamination from subDLAs, as evidenced by detections reported in the literature). However, while it seems likely that strong Ca\\,{\\sc ii} absorption is associated with DLAs, the converse is not true: approximately one third of the DLAs in our sample have \\wca\\,$< 0.15$\\,\\AA, while adopting the criterion \\wca\\,$> 0.30$\\,\\AA\\ would miss $\\simeq 3/4$ of the DLAs.   \\item[3.] Only four out of the 19 DLAs in the combined sample meet the selection criterion \\wca\\,$\\geq 0.35$\\,\\AA\\ of the SDSS survey by WHP6, and only one out of 19 would be classed as a `strong' Ca\\,{\\sc ii} absorber by their definition (\\wca\\,$\\geq 0.5$\\,\\AA). Evidently, there is little overlap as yet between confirmed DLAs and the SDSS candidates identified via Ca\\,{\\sc ii}. This is not surprising, however, since -- as demonstrated by WHP06 -- dust associated with many of the strong Ca\\,{\\sc ii} absorbers dims the light from the background QSOs sufficiently for them to be excluded from the magnitude-limited surveys conducted with \\textit{HST} up to now.  \\item[4.] Overall, we find no convincing correlation of the strength  of Ca\\,{\\sc ii} absorption with the column density of neutral gas, the equivalent width of Mg\\,{\\sc ii}\\,$\\lambda 2796$, the metallicity of the gas (as measured from the abundance of Zn), or the degree of depletion of refractory elements (indicated by the [Cr/Zn] ratio). Despite the fact that, in the interstellar medium of the Milky Way, both Ca and Cr exhibit large and variable depletion factors from the gas-phase, DLAs with near-solar values of [Cr/Zn] do not show significantly stronger Ca\\,{\\sc ii} lines than those where most of the Cr is presumably in solid form. On the other hand, the finding that the only `strong' \\caii\\ absorber in our sample is the most metal-rich DLA yet discovered suggests that there may be some connection between the metallicity and \\caii\\ equivalent width.  \\item[5.] Presumably, a more complex mix of factors  than simply H\\,{\\sc i} column density, metallicity, and dust depletion, determines the value of \\wca\\ in DLAs.  We have presented preliminary evidence that the `joker in the pack' may be the variable fraction of Ca which is singly ionised in H\\,{\\sc i} regions, where Ca\\,{\\sc iii} is often  the dominant ion stage, depending on the densities of particles, far-UV photons, and on the temperature of the gas.  With the data at hand, we exploit the ionisation balance of Ca to place constraints on the ratio  of UV photons and particles.  If we further assume that the interstellar radiation field in the DLAs is of the same order as that at our location within the Milky Way -- as indeed found for DLAs at higher redshifts -- we can arrive at estimates of, or limits on,  the gas density $n$(H$_{\\rm TOT})$.  Our modelling, using the photoionisation code {\\sc cloudy}, indicates values  $n$(H$_{\\rm TOT}) \\simgt 1$\\,cm$^{-3}$ for most DLAs with \\wca\\,$ \\simgt 0.35$\\,\\AA.  Furthermore, for an `average' \\caii-absorber DLA (as described by WHP06) as well as the lone `strong' (as defined by those authors) \\caii\\ system in our sample, we deduce $n$(H$_{\\rm TOT}) \\simgt 10$\\,cm$^{-3}$.  Such high values are similar to those which seem to be typical of H$_2$-bearing DLAs at $z > 2$, raising the possibility that in the strong Ca\\,{\\sc ii} absorbers  selected by WHP06 at $z < 1.3$ we may be seeing the lower redshift counterparts of the subset of DLAs with a measurable molecular fraction.  Indeed, our models predict molecular fractions similar to those reported for relatively metal-rich DLAs at high-$z$, with the strongest \\caii\\ absorbers exhibiting values consistent with the highest reported measurements of $f_{\\rm H_2}$ in DLAs.   \\end{itemize}  If borne out by additional studies, the possible connection of strong \\caii\\ absorbers to metal- and relatively molecule-rich gas has intriguing consequences when considered in light of the imaging results of Hewett \\& Wild (2007).  It would mean that strong \\caii\\ absorbers select gas with these properties not only in galactic disks, but also at distances of tens of kpc from relatively bright galaxies.  Such a conclusion would be highly relevant to models of galactic outflows and star-formation feedback.  The currently observed population of DLAs are {\\it not}, in general, strong \\caii\\ absorbers.  However, if  most or all strong \\caii\\ absorbers are DLAs and the true incidence of strong \\caii\\ absorbers is as high as calculated by WHP06, these intriguing systems may represent a (potentially metal-rich) subset of DLA absorbers that are being missed from the current surveys.   It will be possible to make progress towards clarifying this and other questions raised by the present work with COS measurements of metal lines and \\nhi\\ in a sample of {\\it known} strong \\caii\\ absorbers.  Like the rest of the UV-spectroscopy community, we look forward with anticipation to the forthcoming \\textit{HST} servicing mission that will install this much needed instrument."
    },
    "0808/0808.0196_arXiv.txt": {
        "abstract": "A weakly interacting massive particle (WIMP) weighing only a few GeV has been invoked as an explanation for the signal from the DAMA/LIBRA experiment.   We show that the data from DAMA/LIBRA are now powerful enough to strongly constrain the properties of any putative WIMP.  Accounting for the detailed recoil spectrum, a light WIMP with a Maxwellian velocity distribution and a spin-independent (SI) interaction cannot account for the data. Even neglecting the spectrum, much of the parameter space is excluded by limits from the DAMA unmodulated signal at low energies. Significant modifications to the astrophysics or particle physics can open light mass windows. ",
        "introduction": "Introduction}  The DAMA/LIBRA NaI(Tl) scintillation experiment \\cite{Bernabei:2008yi} has used the annual modulation technique \\cite{Drukier:1986tm,Freese:1987wu} to search for dark matter (DM).  They now find a modulation of over 8$\\sigma$ with a period and phase consistent with a DM signal. However, in some models of DM, it is not trivial to square this positive result with the null results from other direct detection experiments.  Recent investigations \\cite{Foot:2008nw,Feng:2008dz,Petriello:2008jj,Bottino:2008mf}, updating the discussion of \\cite{Gondolo:2005hh} note that light DM, with mass of a few GeV  might reconcile the DAMA/LIBRA data with constraints from other experiments, e.g., \\cite{CDMS,XENON}.   In addition, a DM candidate with mass $\\sim$ GeV is tantalizing -- it might give insight into the ratio of the DM density to the density of baryons.   In this note, we point out that the statistics of the DAMA/LIBRA data are now sufficiently powerful that an explanation of the DAMA/LIBRA data must now go beyond just fitting the overall modulation rate.  Additional  self-consistency checks on the light WIMP scenario are now possible.  We discuss two such checks.  First,  DAMA/LIBRA has now measured the modulation rate as a function of the observed recoil energy.  This spectrum contains valuable information. Simple kinematics indicate that the spectra of nuclei recoiling against a WIMP are sensitive to the mass of the WIMP.   Thus, fitting the observed energy spectrum constrains the mass of the candidate WIMP particle.   Another constraint can be derived by looking at the total (unmodulated) rate of observed events with low energy recoils.  Some WIMP candidates will provide more events than the  {\\it total} number of observed events at low energies, despite a presumably sizable background.  These two constraints are powerful probes of the light WIMP region, and effectively exclude this interpretation if a Maxwellian velocity distribution with standard parameters is assumed.  Modifications to this assumption, in particular DM streams, can open up small regions of allowed parameter space. ",
        "conclusions": "Conclusions} The recent update from the DAMA/LIBRA collaboration has added significance to the modulated signature. Detailed spectral data should now be considered for any model proposed to explain the signal. As it appears that modulation extends up to at least 5 keVee, the spectrum places very strong constraints on the properties of a standard WIMP in a MB velocity distribution. Additionally, because the signal from a standard WIMP rises very rapidly at low recoil energies, one must insure that the {\\em unmodulated} signal predicted does not exceed the observed levels at DAMA/LIBRA.  Fitting to the spectral data, one finds two regions for a standard WIMP in a Maxwellian halo which explain the spectrum. The high mass region conflicts with other direct detection experiments and the unmodulated low-energy rate at DAMA/LIBRA. The light mass region gives a poorer fit to the spectrum, but is still acceptable and is largely consistent with limits from the DAMA unmodulated signal. However,  90\\% CL limits from CDMS-Si and XENON-10 exclude this region.  Modifications to astrophysics or particle physics can open light mass windows. Streams can open regions in the $m_\\chi \\sim 2 -5\\; \\gev$ range, although for very particular choices of parameters. Inelasticity can broaden areas at higher masses, but faces model-building challenges.   Future direct detection experiments should push to explore both of these regions."
    },
    "0808/0808.3476_arXiv.txt": {
        "abstract": "We consider the role of deuterium as a potential marker of location and ambient conditions during the formation of small bodies in our Solar system. We concentrate in particular on the formation of the regular icy satellites of Jupiter and the other giant planets, but include a discussion of the implications for the Trojan asteroids and the irregular satellites.  We examine in detail the formation of regular planetary satellites within the paradigm of a circum-Jovian subnebula. Particular attention is paid to the two extreme potential subnebulae -- ``hot'' and ``cold''. In particular, we show that, for the case of the \"hot\" subnebula model, the D:H ratio in water ice measured from the regular satellites would be expected to be near-Solar. In contrast, satellites which formed in a ``cold'' subnebula would be expected to display a D:H ratio that is distinctly over-Solar.  We then compare the results obtained with the enrichment regimes which could be expected for other families of icy small bodies in the outer Solar system -- the Trojan asteroids and the irregular satellites. In doing so, we demonstrate how measurements by Laplace, the James Webb Space Telescope, HERSCHEL and ALMA will play an important role in determining the true formation locations and mechanisms of these objects. ",
        "introduction": "\\label{intro}  The physical and chemical characteristics of the proto-planetary nebula from which our solar system formed can be inferred through the analysis of ``primitive'' objects such as meteorites, comets, and the giant planets themselves. We can obtain useful constraints on these processes by examining the degree to which fossil deuterium contained within the water in some of these objects is enriched when compared to the protosolar abundance, which can be measured in various objects within the Solar system. Within the Solar nebula, the main reservoir of deuterium was molecular hydrogen. However, isotopic exchange occurred between this hydrogen and other deuterated species, resulting in the formation of secondary reservoirs of deuterium. As a result of the high cosmic abundance of oxygen, the most important secondary reservoir in the nebula is water (HDO), either in gaseous or solid phase.  Calculations of the temporal and radial evolution of the D:H ratio in the primitive nebula (Drouart et al., 1999; Mousis et al., 2000) have been performed in order to reproduce existing data on comets (Balsiger et al., 1995, Eberhardt et al., 1995; Bockel{\\'e}e-Morvan et al., 1998; Meier et al., 1998; see Horner et al., 2007), and measurements taken from meteorites (Deloule et al., 1998). One particularly interesting result of these calculations is that the D:H ratio in water ice produced in the nebula varies by a significant amount as a function of the distance from the Sun at which the ice was formed. This variation can be seen clearly in the compilation of measurements given by Drouart et al. (1999), in Fig.1 of that work. Such results led Horner et al. (2007) to discuss how measurements of the D:H ratio in cometary bodies might prove helpful in answering the question of whereabouts in the Solar system the different cometary populations had formed. The study of D:H extends beyond the study of cometary bodies, however.  In order to consider the measured D:H enrichment within an object to be the direct result of its formation, it is clear that the deuterated reservoir incorporated within that object must have undergone little or no alteration since its formation within the nebula. Therefore, despite the fact that deuterium enhancement has been measured on the Earth, Venus and Mars, it is obvious that we cannot consider the resulting value to be primordial. Indeed, the value measured on the Earth could be the result of the combination of volatile material from a number of sources (Dauphas et al., 2000). Futhermore, the values measured on Venus and Mars are believed to be the result of strong atmospheric fractionation, which has occurred throughout the history of the planets (Donahue, 1999; Bertaux \\& Montmessin, 2001).  Even further from the Sun, observations of Titan reveal an unexpectedly high D:H ratio within the methane of the satellite's atmosphere (B{\\'e}zard et al., 2007).  The cause of this enhancement is still under some debate. It could be the result of isotopic exchange between methane and molecular hydrogen within the early Solar nebula, prior the formation of the icy planetesimals that were ultimately accreted by Titan (e.g. Mousis et al., 2002a). Alternatively, it could result from a photochemical process occurring within the atmosphere of the satellite (Lunine et al., 1999).  A number of ambitious projects, such as the proposed Laplace\\footnote{http://jupiter-europa.cesr.fr/} mission and the James Webb Space Telescope (scheduled for launch in 2013), will allow us to revisit the Jovian satellite system, and provide new measurements of the satellites of the giant planet in the coming years. As such, the time seems right to revisit the question of how these satellites formed, and what effect their formation would have on the quantities that could be observed by such a mission. In light of these proposed missions, such work takes on an interesting new aspect. How would the D:H ratio in the Jovian (or Saturnian) satellites be affected by their formation?  Could measurements of deuterium in these satellites help us to understand the final stages of the formation of their parent bodies?  Furthermore, it is important to examine such ideas during the period over which the missions are designed, to enable the construction of instruments which are fully capable of answering the questions posed.  In this work, we aim to detail the various properties which would have affected the regular satellites during their formation, highlighting how the fractionation of deuterium within their ices may represent a vital window into their formation. We also examine the benefits that observations of the degree of deuteration in the irregular satellites and the Trojan asteroids would have for our understanding of the origin of the volatiles they contain, together with presenting a discussion of how and when such observations may be made.  The simple picture portrayed for the comets, although applicable for other objects which formed freely in the Solar nebula (such as the asteroids), becomes more complicated when one wishes to understand the formation of the satellites of the giant planets. For the regular satellites, it is possible that material from the Solar nebula underwent a significant amount of additional processing within the planetary sub-nebula. In section 2, we will review the main results of prior discussions and calculations involving the D:H ratio in the Solar nebula, while in section 3 we show how further processing within planetary sub-nebulae would eventually lead to a reduction in the D:H ratio incorporated in the satellites forming therein when compared to gas at an equivalent distance in the Solar nebula. In section 4, we examine the cases of the irregular planetary satellites and the Jovian Trojan asteroids, two groups of object for which the study of deuteration may help untangle the minutae of their formation. The measurement of the D:H ratio in satellites and other objects by future space missions is discussed in section 5, and, in section 6, we conclude with a summary and discussion of our ideas, along with their implications for future missions to, and observations of, the giant planets and their satellite systems. ",
        "conclusions": "In previous work (Horner et al., 2007), we examined the role that variations in the D:H ratio through the Solar nebula would have on cometary objects observed to originate from different reservoirs in the outer Solar system. The goal of that work was to highlight that observations of the D:H ratio in such objects could prove a cornerstone in helping our understanding of the different regions in which the various cometary populations formed.  Here, we have extended our arguments to include the other populations of hydrated objects in the outer Solar system - the icy planetary satellites (both regular and irregular), and the Jovian and Neptunian Trojans. We have shown that, since these objects cover a wide range of formation scenarios, the study of the D:H incorporated in their water ice provides a useful tool to answer questions on their origin. Did the regular satellites form in a hot or cold subnebula?  Do the irregular planetary satellites truly represent a captured and then shattered population of objects? If so - from where were they captured? What was the origin of the Jovian and Neptunian Trojans?  We focus, in particular, on the formation of regular satellites in a circum-planetary disk of gas and dust around Jupiter -- the Jovian subnebula. Current theories which describe the formation of such satellites cover a wide range of initial conditions. However, the two extreme cases are the ``hot'' and ``cold'' subnebula models. In the former, the entire subdisk is sufficiently warm that icy material falling into the subnebula from the Solar nebula is entirely vaporised, allowing the exchange of deuterium between molecular hydrogen and water to continue in the gas phase, when the water involved would otherwise already be trapped as ice. The other extreme, the ``cold'' model, assumes that the subnebula was sufficiently cold that none of the infalling volatiles were vaporised prior to their accretion in the planetesimals which went on to form the satellites. Although the true formation scenario for the regular satellites undoubtably lies somewhere between these extremes, they allow us to constrain the behaviour of the ices which went on to form the satellites we currently observe. In the case of the ``hot'' subnebula model, we have shown that the effect of the re-vaporisation of the infalling water ice allows gas phase reactions between HDO and $H_2$ to occur, leading to the gradual depletion of deuterium within the initially D-rich water. As a result, satellites which formed in this way would be expected to exhibit D:H ratios in their ice which are close to the Solar value. By contrast, in the case where the satellites form in a ``cold'' subnebula, since vaporisation of the volatiles does not occur, the D:H ratio within the satellite ices would be significantly higher, representative of material within the Solar nebula at the heliocentric distance at which the satellite's parent planet formed. The measurement of the D:H ratios within the regular satellites therefore provides a key constraint for the discussion of their formation.  Although we do not go into specific detail, the models described above are equally applicable for satellites forming in the Saturnian subnebula.  In the future, new observatories and space missions will allow us to measure $f$ with a greater accuracy, and in many more objects, than ever before.  HERSCHEL, ALMA, Rosetta, and the JWST will allow the measurement of deuterium in objects as diverse as KBOs, comets, and the satelites of the giant planets, while the proposed Laplace Jupiter orbiter offers a unique opportunity to measure deuterium {\\it in-situ}, through use of a dust collector/ablation system, allowing the direct comparison of the regular and irregular Jovian satellites, and offering new insights into their formation.  Early discussion of the type of measurements required is vital for those involved in the planning stages of these projects, in order that appropriate instrumention is constructed to make the required measurements.  Another factor which may have influenced the final value of $f$ obtained by objects (particularly the giant planets and their regular satellites) is that the giant planets are thought to have migrated over a significant distance during their formation and subsequent evolution. In the most extreme examples, it has even been suggested that Uranus and Neptune formed between the orbits of Jupiter and Saturn (Levison et al., 2004)!  In these cases, it is likely that the value of $f$ shown by the regular satellites of these planets will provide a doubly useful tool in determining the true history of the outer Solar system, even though it seems likely that these satellites have been destroyed and reassembled since their formation (e.g. Banfield \\& Murray, 1992). If it is the case that Uranus and Neptune formed far closer to the Sun than their current location, then it seems likely that the material incorporated in their regular satellites could reflect their formation and migration. It must be pointed out, here, that current measurements of the $f$-value in the giant planets themselves provides no useful constraint on their formation location, due to the fact that the measured value is heavily affected by the gaseous hydrogen incorporated during their formation, in addition to water obtained from planetesimals. Given a sufficiently good interior model (e.g. Feuchtgruber et al. 1999), it may be possible to extract the native $f$-value of the accreted icy fraction of the planet, which would clearly prove very useful in the study of the planetary migration. However, the satellites currently offer a simpler solution to the problem, since they formed solely from the accretion of planetesimals.  Studies of these objects could provide detailed information on their $f$-values, which can constrain their migration and formation histories. Were the satellites present prior to the migration? Did they form afterward, or during the process?  Similarly, should it turn out that the ``cold subnebula'' model is the fairest representation of the formation of the regular Jovian and Saturnian satellites, then it is possible that these objects contain, trapped within their ices, a record of the migration of their parent planets through the Solar nebula. Clearly, such measurements could even be used to place constraints on the distance over which the planets migrated."
    },
    "0808/0808.3195_arXiv.txt": {
        "abstract": "Electrostatic behavior of a collisionless plasma in the foot region of high Mach number perpendicular shocks is investigated through the two-dimensional linear analysis and electrostatic particle-in-cell (PIC) simulation. The simulations are double periodic and taken as a proxy for the situation in the foot. The linear analysis for relatively cold unmagnetized plasmas with a reflected proton beam shows that obliquely propagating Buneman instability is strongly excited. We also found that when the electron temperature is much higher than the proton temperature, the most unstable mode is the highly obliquely propagating ion two-stream instability excited through the resonance between ion plasma oscillations of the background protons and of the beam protons, rather than the ion acoustic instability that is dominant for parallel propagation.  To investigate nonlinear behavior of the ion two-stream instability, we have made PIC simulations for the shock foot region in which the initial state satisfies the Buneman instability condition. In the first phase, electrostatic waves grow two-dimensionally by the Buneman instability to heat electrons. In the second phase, highly oblique ion two-stream instability grows to heat mainly ions. This result is in contrast to previous studies based on one-dimensional simulations, for which ion acoustic instability further heats electrons.  The present result implies that overheating problem of electrons for shocks in supernova remnants is resolved by considering ion two-stream instability propagating highly obliquely to the shock normal and that multi-dimensional analysis is crucial to understand the particle heating and acceleration processes in shocks. ",
        "introduction": "The discovery of thermal and synchrotron X-rays from young supernova remnants (SNRs) provides the evidence that electrons are heated up to a few keV and that a portion of them are accelerated to highly relativistic energy in SNR shocks \\citep{koy95}. Because SNR shocks are collisionless, not only particle acceleration mechanisms but also electron heating mechanisms in SNR shocks are not so simple. Previous studies have given an important key to the formation mechanism of perpendicular collisionless shocks. When the Alfv${\\rm \\acute{e}}$n Mach number $M_{\\rm A}$ is larger than the critical Mach number, about 3, a perpendicular shock reflects some of the incident ions to the upstream, where a foot region forms on a spatial scale of the ion gyroradius \\citep{ler83}. The plasma in the foot region consists of incident ions and electrons and reflected ions and returning ions which are made from reflected ions and move to the shock after a gyration. As for the electron heating machanism, \\citet{pap88} proposed that when the Mach number is larger than $0.5(m_{\\rm p}/m_{\\rm e})^{1/2}\\sim 20$, incident electrons and reflected ions excite electrostatic waves by the Buneman instability \\citep{bun58} because the relative velocity between them is large compared with the electron thermal velocity. They also suggest that after electrons are  heated by electrostatic waves induced by the Buneman instability, ion acoustic instability is triggered because the electron temperature becomes much higher than the proton temperature. As a result, electrons are strongly heated by the Buneman instability and the ion acoustic instability. \\citet{car88} performed a one-dimensional hybrid simulation and demonstrated that strong electron heating actually occurs. They concluded that an $M_{\\rm A}=500$ shock heats electrons by a factor of $10^5$ across the shock. This means that if the upstream electron temperature is 1 eV, the downstream electron temperature becomes 100 keV. This value of the downstream temperature is much larger than the recent observational one for SNRs, a few keV \\citep{sta06}. This discrepancy has been an open issue to be resolved for a long time.  On the other hand, \\citet{shi00} and \\cite{hos02} performed one-dimensional full particle-in-cell (PIC) simulations to investigate the electron acceleration at perpendicular shocks. Their simulation solves a whole region of a collisionless perpendicular shock and makes reflected ions self-consistently by employing a small proton to electron mass ratio. Their results showed that electrons are not only heated at the foot region but also significantly accelerated by surfing acceleration mechanism. However, this acceleration is valid only for the one-dimensional case because the surfing acceleration strongly depends on the structure of the electrostatic potential. In our first paper \\citep{ohi07}, we performed two-dimensional electrostatic PIC simulations to solve for the two-dimensinal structure of the electrostatic potential excited by the Buneman instability. We employed the real mass ratio but the simulation region is limited to the foot region. Our results showed that oblique modes grow as strongly as the modes parallel to the beam direction, that the potential structure becomes two-dimensional and that no efficient surfing acceleration occurs, while electron heating occurs. Thus, the problem of electron acceleration has been back to the start again. In that paper, we concentrated on the stage of the Buneman instability and did not follow the long time scale evolution after the Buneman nstability has saturated.  In this paper, we study the time evolution of electrostatic collisionless plasma instabilities in the foot region by making linear analysis and by performing two-dimensional electrostatic PIC simulation. We perform simulations with a higher resolution, a larger simulation box and a longer simulation time than in \\citet{ohi07}. Especially, we focus on the evolution after electrostatic waves excited by the Buneman instability have decayed. Our simulation substantially improves most previous works that are one-dimensional, employ an artifically small proton electron mass ratio or impose rather strong magnetic field. It is obvious that a multi-dimensional analysis is necessary as discussed above \\citep{blu60, lam74, ohi07}. Our motivation for employing the real mass ratio is as follows. For a small mass ratio, the foot region in the simulation is shorter than the realistic one and the time scale on which electrons stay in the foot region in the simulation is also shorter than the realistic one because the size of the foot region is about the ion gyroradius $m_{\\rm p}v_{\\rm d}c/eB$, where $v_{\\rm d}$ and $B$ are the drift velocity of reflected protons and the magnetic field, respectively. Because reflected ions have a large free energy, we expect that more energy is transported to electrons through collective instabilities with the realistic mass ratio in the foot region. The drift velocity is not large enough to excite electromagnetic waves, so that electrostatic waves are more important.  In \\S 2 we perform linear analysis for two-dimensional electrostatic modes. In \\S 3 we describe the initial setting of the PIC simulations and numerical results, followed by a discussion in \\S 4. ",
        "conclusions": "Now we discuss the final outcome of the ion two-stream instability. It becomes unstable when two conditions are satisfied; one is the resonance condition, $k_{x} v_{\\rm d} \\sim \\omega_{\\rm pi}$ and the other is that the wave length be between the ion Debye length and the electron Debye length. Thus, the ion two-stream instability becomes stabilized when $T_{\\rm p} \\sim T_{\\rm e}$. At the end of our simulation, the $y$-component of electric field has not decayed still completely. So, we expect $T_{\\rm e}/T_{\\rm p} < 10$ in the final stage, and probably ions will be heated up to $T_{\\rm i} \\sim m_{\\rm e} v_{\\rm d} ^2 \\sim 4m_{\\rm e} v_{\\rm sh}^2$ at the foot region in high Mach number perpendicular shocks. Of course, because the proton temperature of the drift direction is still cold in the present simulation, we must check the isotropilazation process of ion velocity distribution by doing longtime full PIC-simulations.  As for the electrons, the electron temperature in the foot region is also about $m_{\\rm e} v_{\\rm sh}^2$. In the later stage, the growth of the two-stream instability dominates over the ion acoustic instability and little electron heating occurs. Hence, as mentioned in \\citet{ohi07}, if other electron heating mechanisms do not exist, after passing the shock front, electrons will undergo the adiabatic heating and finally, the electron temperature in the downstream becomes \\begin{equation} T_{\\rm e}\\sim 4\\times \\frac{1}{2}m_{\\rm e}(2v_{\\rm sh})^2 = 0.41{\\rm keV} \\left(\\frac{v_{\\rm sh}}{0.01c}\\right)^2, \\end{equation} where we assume that the compression ratio is 4. This has a very important implication for the electron heating process of the SNR shocks. The proton temperature in the downstream is $T_{\\rm p}=3m_{\\rm p}v_{\\rm sh}^2/16$, hence the ratio of two temperatures is \\begin{equation} T_{\\rm e}/T_{\\rm p} \\sim \\frac{128}{3}\\frac{m_{\\rm e}}{m_{\\rm p}} \\sim 0.023. \\end{equation} This value is close to the observed value as long as the shock velocity $v_{\\rm sh}$ is larger than $1500{\\rm km/s}$ \\citep{ade08}. Namely, we expect that the overheating problem of electrons raised by \\citet{car88} can be solved by the ion two-stream instability.   In this paper, we prepare three ion beams as the initial condition. We have also performed simulations for other initial conditions such that there exist upstream electrons, upstream and reflected protons. Electrons and reflected protons have drift velocities to satisfy the vanishing current condition. This situation corresponds to the initial phase of the shock reformation phenomena. The results turn out to be basically the same as those in the results presented in this paper.   Our simulation does not include electromagnetic modes. In the linear stage of the Buneman and ion two-tream instabilities, the effects are negligible because the growth rates of  electromagnetic modes are smaller than those of electrostatic modes for $v_{\\rm d} \\sim 0.01c$. In the nonlinear stage, electromagnetic modes may be important. One of the reasons is that the electrostatic waves with a wave vector almost perpendicular to the drift direction may make currents because the charge fluctuation can not be shielded by electrons in this situation, where the fluctuation scale is smaller than the electron Debye length. Consequently the current may make the magnetic field. It is an interesting speculation that the magnetic field might be amplified more rapidly by the ion two-stream instability than the ion Weibel instability. The other reason is that anisotropic ion heating caused by the highly oblique ion two-stream instability excites the Weibel instability due to ion temperature anisotropy. At the end of our simulations, the ion temperature of $x$-direction is almost the same as initial one and the ratio of the ion temperature of $y$-direction to that of $x$-diretion $T_{\\rm iy}/T_{\\rm ix}$ is about 100. These two features may lead to magnetic field amplification and accompanying particle acceleration and heating in the shock foot region. We will make full-PIC simulations in future work to investigate these issues."
    },
    "0808/0808.1968_arXiv.txt": {
        "abstract": "We present the first results on galaxy metallicity evolution at z$>$3 from two projects, LSD (Lyman-break galaxies Stellar populations and Dynamics) and AMAZE (Assessing the Mass Abundance redshift Evolution). These projects use deep near-infrared spectroscopic observations of a sample of $\\sim$40 LBGs to estimate the gas-phase metallicity from the emission lines. We derive the mass-metallicity relation at z$>$3 and compare it with the same relation at lower redshift. Strong evolution from z=0 and z=2 to z=3 is observed, and this finding puts strong constrains  on the models of galaxy evolution. These preliminary results show that the effective oxygen yields does not increase with stellar mass, implying that the simple outflow model does not apply at z$>$3. ",
        "introduction": "Metallicity is one the most important property of galaxies, and its study is able to shed light on the detail properties of galaxy formation. It is an integrated property, not related to the present level of star formation the galaxy, but rather to the whole past history of the galaxy. In particular, metallicity is sensitive to the fraction of baryonic mass already converted into stars, i.e., to the evolutive stage of the galaxy. Also, metallicity is effected by the presence of inflows and outflows, i.e., by feedback processes and the interplay between the forming galaxy and the intergalactic medium.  It is well known that local galaxies follow a well-defined mass-metallicity relation, where galaxies with larger stellar mass have higher metallicities (\\cite{tremonti04,lee06}). The origin of the relation is uncertain because several effects can be, and probably are, active. It is well known that in the local universe starburst galaxies eject a significant fraction of metal-enriched gas into the intergalactic medium because of the energetic feedback from exploding SNe, both core-collapse and, possibly, type Ia (\\cite{mannucci06}). Outflows are expected to be more important in low-mass galaxies, where the gravitational potential is lower and a smaller fraction of gas is retained. As a consequence, higher mass galaxies are expected to be more metal rich (see, for example, \\cite{edmunds90,garnett02}). A second possibility is related to the well know effect of ``downsizing'' (e.g., \\cite{cowie96}), i.e., lower-mass galaxies form their stars later and on longer time scales. At a given time, lower mass galaxies have formed a smaller fraction of their stars, therefore are expected to show lower metallicities. Other possibilities exist, for example some properties of star formation, as the initial mass function (IMF), could change systematically with galaxy mass (\\cite{koppen07}).  All these effects have a deep impact on galaxy formation, and the knowledge of their relative contributions is of crucial importance. Different models have been built to reproduce the shape of the mass-metallicity relation in the local universe, and different assumptions produce divergent predictions at high redshifts (z$>$2). To explore this issue several groups have observed the mass-metallicity relation in the distant universe, around z=0.7 (\\cite{savaglio05}) and z=2.2 (\\cite{erb06}). They have found a clear evolution with cosmic time, with metallicity for a given stellar mass decreasing with increasing redshift.  For several reasons, it is very interesting to explore even higher redshifts. The redshift range at z$\\sim$3--4 is particularly interesting:  it is before the peak of the cosmic star formation density (see, for example, \\cite{mannucci07}), only a small fraction ($~$15\\%, \\cite{pozzetti07}) of the total stellar mass have already been created, the number of mergers among the galaxies is much larger than at later times (\\cite{conselice07}). As a consequence, the prediction of the different models tend to diverge above z=3, and it is important to sample this redshift range observationally. The observations are really challenging because of the faintness of the targets and the precision required to obtain a reliable metallicity. Nevertheless, the new integral-field unit (IFU) instruments on 8-m class telescopes are sensitive enough to allow for the project. ",
        "conclusions": ""
    },
    "0808/0808.4126_arXiv.txt": {
        "abstract": "The carrier of the 2175$\\Angstrom$ interstellar extinction feature remains unidentified since its first detection over 40 years ago. In recent years carbon buckyonions have been proposed as a carrier of this feature, based on the close similarity between the electronic transition spectra of buckyonions and the 2175$\\Angstrom$ interstellar feature. We examine this hypothesis by modeling the interstellar extinction with buckyonions as a dust component. It is found that dust models containing buckyonions (in addition to amorphous silicates, polycyclic aromatic hydrocarbon molecules, graphite) can closely reproduce the observed interstellar extinction curve. To further test this hypothesis, we call for experimental measurements and/or theoretical calculations of the infrared vibrational spectra of hydrogenated buckyonions. By comparing the infrared emission spectra predicted for buckyonions vibrationally excited by the interstellar radiation with the observed emission spectra of the diffuse interstellar medium, we will be able to derive (or place an upper limit on) the abundance of interstellar buckyonions. ",
        "introduction": "In the interstellar medium (ISM), the strongest spectroscopic extinction feature is the 2175$\\Angstrom$ bump which is characterized with a stable peak wavelength and a width variable from one sightline to another (Fitzpatrick \\& Massa 2007). Since Stecher (1965) first detected this ultraviolet (UV) extinction feature through rocket observations, the origin of this feature and the nature of its carrier(s) are still an enigma. Many candidate materials, including graphite, amorphous carbon, graphitized (dehydrogenated) hydrogenated amorphous carbon, nano-sized hydrogenated amorphous carbon, quenched carbonaceous composite, coals, polycyclic aromatic hydrocarbon (PAH), and OH$^{-}$ ion in low-coordination sites on or within silicate grains have been proposed, while no single one is generally accepted (see Li \\& Greenberg 2003 for a review).  Recently, Chhowalla et al.\\ (2003) measured the UV-visible photoabsorption spectra of carbon buckyonions (BOs)\\footnote{% Early pioneering experimental studies of the UV absorption properties of BOs and their association with the 2175$\\Angstrom$ interstellar extinction feature include those of de Heer \\& Ugarte (1993), Ugarte (1995), and Wada et al.\\ (1999). } composed of spherical concentric fullerene shells (so far the largest BOs produced in laboratory have $N\\sim 100$ shells; Iglesias-Groth et al.\\ 2003). They found that the plasmon-like feature of BOs due to a collective excitation of the $\\pi$ electrons closely fits the 2175$\\Angstrom$ interstellar extinction feature. More recently, Ruiz et al.\\ (2005) theoretically simulated the photoabsorption spectra of BOs.\\footnote{% Early pioneering theoretical studies of the UV absorption properties of BOs and their association with the 2175$\\Angstrom$ interstellar extinction feature include those of Wright (1988), Henrard et al.\\ (1993, 1997) and Lucas et al.\\ (1994). } They found that the calculated absorption spectra of BOs are in close agreement with the experimental data of Chhowalla et al.\\ (2003) and the observed 2175$\\Angstrom$ interstellar feature. That the $\\pi$-plasmon absorption band of BOs exhibits a stable peak position at $\\simali$5.70\\,eV ($\\lambda\\approx 2175\\Angstrom$, $\\lambda^{-1}\\approx 4.6\\mum^{-1}$) and a variable bandwidth (Chhowalla et al.\\ 2003, Ruiz et al.\\ 2005) suggests that BOs may be a promising candidate material for the 2175$\\Angstrom$ interstellar extinction feature.  Indeed, it is shown by Chhowalla et al.\\ (2003) and Ruiz et al.\\ (2005) that the photoabsorption spectra of BOs {\\it alone} very well reproduce the entire interstellar extinction curve at $\\lambda^{-1}$\\,$\\simali$3.2--7.3$\\mum^{-1}$. This requires $\\simali$190\\,ppm (parts per million) C/H to be locked up in BOs (Ruiz et al.\\ 2005). However, it is well recognized that in the ISM, in addition to the 2175$\\Angstrom$ extinction carrier, there must exist other dust components as well -- there must be a population of amorphous silicate dust, as indicated by the strong, ubiquitous 9.7 and 18$\\mum$ interstellar absorption features; there must be a population of aromatic hydrocarbon dust (presumably PAH molecules), as indicated by the distinctive set of ``unidentified'' infrared (UIR) emission bands at 3.3, 6.2, 7.7, 8.6, and 11.3$\\mum$ ubiquitously seen in the ISM; there must also exist a population of aliphatic hydrocarbon dust, as indicated by the 3.4$\\mum$ C--H absorption feature which is also ubiquitously seen in the diffuse ISM of the Milky Way and external galaxies (see Li 2004).  Although buckyonions are able to closely reproduce the 2175$\\Angstrom$ extinction feature, it is not clear if dust models with amorphous silicates, PAHs, and other carbon dust species (e.g. amorphous carbon, hydrogenated amorphous carbon, organic refractory, and graphite) incorporated (in addition to BOs) are still capable of fitting the 2175$\\Angstrom$ extinction feature. One may intuitively expect that the almost perfect fit to the $\\lambda^{-1}$\\,$\\simali$3.2--7.3$\\mum^{-1}$ interstellar extinction would easily be distorted by the addition of other dust components. It is the purpose of this {\\it Letter} to examine this issue. To this end, we consider the extinction of dust models consisting of multi-modal grain populations with BOs as an interstellar grain component.  \\begin{figure} \\begin{center} \\includegraphics{f1.eps} \\end{center} \\caption{\\label{fig:mrn_solar} Comparison of the interstellar extinction (dot-dashed line) with the model extinction (solid line with oscillatory features) obtained by summing up the contributions of buckyonions (C/H\\,=\\,115\\,ppm), PAHs (C/H\\,=\\,60\\,ppm), graphite (C/H\\,=\\,35\\,ppm), and amorphous silicate (Si/H\\,=\\,35\\,ppm). We assume a MRN-type size distribution for the silicate and graphite grains. } \\end{figure} ",
        "conclusions": "It is encouraging that the dust models consisting of multi-modal grain populations (including BOs) fit the interstellar extinction curve (including the 2175$\\Angstrom$ extinction feature) very well while satisfying the interstellar abundance constraints. The fit to the 2175$\\Angstrom$ feature is mainly affected by the silicate dust component. The smaller amount of silicate dust is included in the model, the better is the fit. We would achieve a closer match to the 2175$\\Angstrom$ extinction if we take the interstellar Si abundance to be subsolar (Snow \\& Witt 1996) like that of B stars $\\sidust=18\\ppm$ (Sofia \\& Meyer 2001) while having C/H\\,=\\,140\\,ppm in BOs and C/H\\,=\\,10\\,ppm in graphite.\\footnote{% If both C and Si are subsolar (say, 2/3 of that of solar, see Snow \\& Witt 1996), we will not be able to fit the interstellar extinction (also see Li 2005). }  BOs are a promising candidate material for the 2175$\\Angstrom$ interstellar extinction feature: (1) their $\\pi$-plasmon absorption profile closely resembles the 2175$\\Angstrom$ extinction feature and is stable in peak wavelength position while varies in width;\\footnote{% Ruiz et al.\\ (2005)'s theoretically simulated photoabsorption spectra of BOs showed that the width of the 2175$\\Angstrom$ feature $\\gammabump$ varies with size for small BOs, however, the width becomes independent of size when the number of shells $N\\simgt 6$ ($\\gammabump \\approx 1.03\\mum^{-1}$; see their Fig.\\,4). For BOs to explain the 2175$\\Angstrom$ interstellar features broader than $\\gammabump = 1.03\\mum^{-1}$ (some sightlines have $\\gammabump$ up to $\\simali$1.2$\\mum^{-1}$), clustering of individual BOs (de Heer \\& Ugarte 1993, Rouleau et al.\\ 1997), imperfect growth of BOs (e.g. mixing with amorphous carbon impurities; see Chhowalla et al.\\ 2003), and coating BOs with a layer of PAHs (Mathis 1994) may play an important role. } (2) the almost perfect fit to the interstellar feature is not significantly distorted by the inclusion of other dust components (e.g. silicate, PAHs, graphite) which are required to account for other interstellar phenomena (e.g. the 9.7, 18$\\mum$ absorption features, the ``UIR'' bands); and (3) BOs are highly stable molecules that they can survive under intense UV radiation and are highly resistant to destruction by collisions.  BOs can be formed by electron beam irradiation of carbon soot (Ugarte 1992) and by heat treatment (annealing at temperatures $T\\simgt 700\\,^{\\circ}$C) and electron beam irradiation of nanodiamond (Kuznetsov et al.\\ 1994, Tomita et al.\\ 2002). The generation of BOs by annealing nanodiamonds is astrophysically relevant. Presolar nanodiamonds are identified in primitive carbonaceous meteorites based on their isotopic anomalies (Lewis et al.\\ 1987). The possible existence of nanodiamonds in dense clouds and circumstellar dust disks or envelopes have also been suggested (Allamandola et al.\\ 1992, Guillois, Ledoux, \\& Reynaud 1999, van Kerckhoven, Tielens, \\& Waelkens 2002). The generation of BOs through annealing nanodiamonds can occur in the ISM where nanodiamonds are stochastically heated to temperatures as high as $\\simali$1000$\\K$ by the interstellar radiation field (Jones \\& d'Hendecourt 2000). Indeed, BOs have been found in meteorites with anomalous isotopic compositions (Smith \\& Buseck 1981, Bernatowicz et al.\\ 1996, Harris et al.\\ 2000).  Recently, BOs have also been invoked to explain other interstellar phenomena: the mysterious diffuse interstellar bands (DIBs) resulting from the electronic transitions of BOs (Iglesias-Groth 2004, 2007),\\footnote{% To date, $>$300 DIBs have been detected in the Galactic and extragalactic ISM. These mysterious optical-to-near-IR absorption lines (broader than the narrow lines from gas atoms, ions, and small molecules) still remain unidentified. Fullerenes are known to have strong structural resonances at this wavelength range (Dresselhaus et al.\\ 1996, Iglesias-Groth 2004). Iglesias-Groth (2007) argued that BOs may be responsible for the strongest optical DIB at 4430$\\Angstrom$ and possibly also the 6177$\\Angstrom$ and 6284$\\Angstrom$ DIBs. } the 10--100\\,GHz Galactic anomalous microwave emission resulting from the electric dipole emission of rotationally excited BOs (Iglesias-Groth 2005, 2006), and the broad feature around 100$\\mum$ of the diffuse emission from two active star-forming regions (the Carina Nebula and Sharpless 171) resulting from the small particle surface resonance (Onaka \\& Okada 2003).  Experimentally, all BOs are found to belong to the C$_{60N^2}$ icosahedral family with $N=1,2,...$ (i.e. their shells are consecutive elements of the icosahedral C$_{60N^2}$ fullerene family, with C$_{60}$ being the smallest shell\\footnote{% Kroto et al.\\ (1985) first proposed that C$_{60}$ could be present in the ISM with a considerable quantity. This molecule and its related species were later proposed as the carriers of the 2175$\\Angstrom$ extinction hump, the DIBs, the ``UIR'' bands, and the extended red emission (see Webster 1991, 1992, 1993a,b). C$_{60}$ and C$_{70}$ are unlikely a major contributor to the 2175$\\Angstrom$ extinction feature since they have a characteristic doublet absorption in this wavelength region. Foing \\& Ehrenfreund (1994) attributed the two DIBs at 9577$\\Angstrom$ and 9632$\\Angstrom$ to C$_{60}^{+}$. However, attempts to search for these molecules in the UV and IR were unsuccessful (Snow \\& Seab 1989; Somerville \\& Bellis 1989; Moutou et al.\\ 1999; Herbig 2000). These molecules are now estimated to consume at most $<$0.7$\\ppm$ carbon (Moutou et al.\\ 1999). Therefore, C$_{60}$ is at most a minor component of the interstellar dust family. } and the intershell distance being very close to the interplanar separation in graphite, often taken to be 3.55$\\Angstrom$, the C$_{60}$ radius; Yoshida \\& 1993, Ruiz et al.\\ 2004). So far,  all the experimentally generated and analyzed BOs are in the nanometer size range ($\\simali$3--15\\,nm in diameter; see de Heer \\& Ugarte 1993, Cabioc'h et al.\\ 1997, Chhowalla et al.\\ 2003).   The fullerenes produced in laboratory are usually accompanied by tubular and other non-spherical graphite-like structures of which the theoretical UV spectra exhibit some similarity to that of BOs. In the ISM, tubular graphite and PAHs can be catalytically formed at $T\\simgt 1000\\K$ on Fe nanoparticles and are probably the carrier of some of the DIBs (Zhou et al.\\ 2006). They may also contribute to the 2175$\\Angstrom$ interstellar extinction feature.  Finally, we should note that a powerful test of the BOs hypothesis would be in the IR. Because of their small sizes (and therefore small heat capacities), in the ISM, BOs will be stochastically heated by single UV photons (Draine \\& Li 2001). With their surface shell hydrogenated, they will emit in the near- to mid-IR through their characteristic C--H stretching and bending bands, and C--C stretching bands. The detection (or non-detection) of these bands will allow us to derive (or place an upper limit on) the abundance of BOs. Note that the BO models require an appreciable amount of C/H to be in BOs ($>110\\ppm$), nearly twice as much as PAHs. Unfortunately, little is known about the positions and strengths of these vibrational bands. We call for urgent experimental measurements and/or theoretical calculations of the IR vibrational spectra of hydrogenated BOs. Future far-IR/submm space missions (e.g. Herschel) will also be useful for studying their far-IR vibrational spectra (see Iglesias-Groth \\& Bret\\'on 2000). Moreover, it would be interesting to see if the silicate-graphite-PAH-BO multi-component dust model is able to reproduce the $\\simali$2--3000$\\mum$ overall IR emission of the Galactic ISM."
    },
    "0808/0808.0914_arXiv.txt": {
        "abstract": "We present Space Telescope Imaging Spectrograph observations of 14 nearby low-luminosity active galactic nuclei, including 13 LINERs and 1 Seyfert, taken at multiple parallel slit positions centered on the galaxy nuclei and covering the H$\\alpha$ spectral region. For each galaxy, we measure the emission-line velocities, line widths, and strengths, to map out the inner narrow-line region structure, typically within $\\sim$100 pc from the galaxy nucleus. There is a wide diversity among the velocity fields: in a few galaxies the gas is clearly in disk-like rotation, while in other galaxies the gas kinematics appear chaotic or are dominated by radial flows with multiple velocity components. In most objects, the emission-line surface brightness distribution is very centrally peaked. The [\\ion{S}{2}] line ratio indicates a radial stratification in gas density, with a sharp increase within the inner 10--20 pc, in the majority of the Type 1 (broad-lined) objects. The electron-density gradients of the Type 1 objects exhibit a similar shape that is well fit by a power law of the form $\\emph{n}_{\\mathrm{e}} = \\emph{n}_0(r/1~ \\mathrm{pc})^{\\alpha}$, where $\\alpha = -0.60 \\pm{0.13}$. We examine how the [\\ion{N}{2}] $\\lambda 6583$ line width varies as a function of aperture size over a range of spatial scales, extending from scales comparable to the black hole's sphere of influence to scales dominated by the host galaxy's bulge. For most galaxies in the sample, we find that the emission-line velocity dispersion is largest within the black hole's gravitational sphere of influence, and decreases with increasing aperture size toward values similar to the bulge stellar velocity dispersion measured within ground-based apertures. We construct models of gas disks in circular rotation and show that this behavior can be consistent with virial motion, although for some combinations of disk parameters we show that the line width can increase as a function of aperture size, as observed in NGC 3245. Future dynamical modeling in order to determine black hole masses for a few objects in this sample may be worthwhile, although disorganized motion will limit the accuracy of the mass measurements. ",
        "introduction": "\\label{sec:intro}  With its $\\sim$ 0\\farcs1 resolution, the Space Telescope Imaging Spectrograph (STIS) aboard the \\emph{Hubble Space Telescope (HST)} has provided the exceptional ability to spatially resolve gas kinematics within the gravitational sphere of influence of the putative central black hole in nearby galaxies. Resolving the gas kinematics very close to the black hole allows for the determination of black hole mass in a relatively straightforward manner, provided that the gas is in Keplerian rotation in a thin, disk-like structure. STIS has thus had a major impact on the field of supermassive black hole detection in galactic nuclei, and data from STIS have contributed greatly toward establishing the correlations between black hole mass and host-galaxy properties, such as those with the stellar velocity dispersion \\citep{Ferrarese_2000,Gebhardt_2000,Tremaine_2002} and bulge luminosity \\citep{Kormendy_Gebhardt_2001}.  More recently, ground-based integral field unit (IFU) observations have provided two-dimensional (2D) kinematic information on the ionized gas in galaxy centers \\citep{Sarzi_2006, Dumas_2007}. Additionally, ground-based IFU observations with the assistance of adaptive optics have produced gas dynamical black hole mass measurements \\citep{Hicks_Malkan_2007, Neumayer_2007}. Without the use of adaptive optics, however, ground-based observations are usually unable to resolve scales comparable to the black hole radius of influence in nearby galaxies. Although STIS is a long-slit instrument, 2D information similar to IFU data can be obtained with STIS by using multiple slit positions.  Even with the past decade of high-resolution observations, only about 30 dynamical black hole mass measurements have been made \\citep{Ferrarese_Ford_2005}. Consequently, searching for regular gas velocity fields dominated by gravitational motion is essential for extending the number of black hole mass measurements and further establishing the connections between black holes and host galaxies. While searching for regular gas velocity fields is clearly important, it appears, however, that they are fairly rare. Instead, complicated and chaotic emission-line velocity fields are more commonly found near the centers of galaxies \\citep{Ho_2002, Atkinson_2005}. Although only upper limits on the black hole mass can be estimated using gas dynamical modeling in these situations \\citep{Sarzi_2002}, subarcsecond-resolution observations can still be useful for studying the dynamical state of the gas and searching for radial motions that might be related to active galactic nucleus (AGN) fueling or outflows. Modeling AGN outflows on subarcsecond scales \\citep{Capetti_1996, Crenshaw_Kraemer_2000} is important for furthering our understanding of the connection between AGN feedback and host-galaxy properties. Of particular interest are the connection to star formation, the growth of the supermassive black hole, and how gas may be driven out from the host galaxy as predicted in some feedback scenarios \\citep{Springel_2005, Hopkins_2006, Sijacki_2007}.  Moreover, the narrow-line region (NLR) velocity dispersion ($\\sigma_{g}$), often measured from the width of the [\\ion{O}{3}] $\\lambda 5007$ line, may provide a reasonable estimate of the host-galaxy stellar velocity dispersion ($\\sigma_\\star$), albeit with substantial scatter \\citep{Nelson_Whittle_1996, Greene_Ho_2005}. Detailed NLR studies are essential in understanding the origin of the scatter in the $\\sigma_{g}$--$\\sigma_\\star$ relationship. This relationship can be useful for distant and luminous AGNs where measurement of the stellar velocity dispersion is not feasible \\citep{Salviander_2007, Shields_2003}. Using the narrow emission-line widths as a proxy for stellar velocity dispersion is only meaningful if the NLR gas is dominated by virial motion in the host-galaxy bulge and not by the gravitational influence of the black hole or nongravitational motions. Therefore, outflows and similar nongravitational motions, which can disrupt viral motion, are one potential cause of the scatter in the $\\sigma_{g}$--$\\sigma_\\star$ relationship.  While nongravitational motions are clearly important in setting the narrow-line widths in AGNs with strong outflows, there can be an important contribution to the line widths from nongravitational motion even in objects that are not outflow-dominated. From \\emph{HST} STIS observations, \\cite{VerdoesKleijn_2006} found that circumnuclear gas in Fanaroff-Riley 1 radio galaxies typically has an intrinsic velocity dispersion in excess of that expected from purely gravitational motion in the host-galaxy potential. They concluded that this excess line width is driven by the injection of energy by the AGN, although there was no obvious relationship between the AGN radiative luminosity and the magnitude of the nongravitational line dispersion.  Other origins for the scatter were studied by \\cite{Greene_Ho_2005} using a large sample of Sloan Digital Sky Survey (SDSS) type 2 AGNs. They looked for correlations of nuclear properties (radio power, AGN luminosity, Eddington ratio), as well as global properties (host-galaxy morphology, local environment, star-formation rate), with the excess line width, defined to be the difference between the dispersion of the [\\ion{O}{3}] $\\lambda 5007$ line and the stellar velocity dispersion. They found a strong correlation between the excess line width and Eddington ratio, where larger excess line widths, indicative of outflowing gas, were seen in the higher Eddington ratio objects.  Another potential explanation for the scatter in the $\\sigma_{g}$--$\\sigma_\\star$ relationship is the effect of observational aperture size. \\cite{Rice_2006} investigated the relationship between line width and aperture size in a sample of nearby Seyferts using single-slit STIS observations. They found a wide range of behavior in how the widths of the [\\ion{O}{3}] $\\lambda 5007$ and [\\ion{S}{2}] $\\lambda\\lambda$6716, 6731 lines varied as a function of radius, but the most common trends were either a roughly constant width or an increase in width with increasing aperture size. Their analysis showed that the line widths measured within the largest STIS aperture were systematically smaller by 10\\%--20\\% than both the stellar velocity dispersion and the line widths measured within a ground-based sized aperture. They concluded that the smaller line width measured within the randomly oriented STIS slit is the result of sampling a smaller portion of the velocity field than the larger ground-based aperture.  Beyond the $\\sigma_{g}$--$\\sigma_\\star$ relationship, we can address other open questions concerning low-luminosity AGNs. \\cite{Laor_2003} pointed out that while the NLR in luminous AGNs occurs on scales of tens of parsecs, where its dynamics are dominated by the bulge, the NLR in low-luminosity AGNs is more compact, and its dynamics could be dominated by the black hole. It is therefore interesting to ask whether the narrow-line widths in low-luminosity AGNs are set by the black hole or by the bulge potential.  The NLR emission lines also provide information on ionization mechanisms. Spectra of low-ionization nuclear emission-line regions \\citep[LINERs;][]{Heckman_1980} can be reproduced with photoionization models involving a nonstellar continuum and low ionization parameters \\citep{Ferland_Netzer_1983, Halpern_Steiner_1983}, but high electron densities, as well as a large range of densities, are needed in order to reproduce the strength of the [\\ion{O}{3}] $\\lambda 4363$ line \\citep{Filippenko_Halpern_1984, Filippenko_1985, Ho_1993}. \\cite{Filippenko_Halpern_1984} demonstrated that such conditions exist in the LINER NGC 7213, and that there is a correlation between line width and the critical density for collisional deexcitation. A similar relationship between line width and ionization potential was also previously seen in Seyferts \\citep{Pelat_1981}. This implied that the NLR is not a homogeneous environment, but must have radial gradients in velocity and density. In ground-based spectra these spatial gradients were generally unresolved, but using STIS it is possible to directly resolve the density gradients in the inner NLR \\citep{Barth_2001a, Shields_2007}. It is therefore of interest to determine what fraction of LINERs show spatially resolved density gradients, and whether the type 2 LINERs, where a central photoionizing source is not always seen, exhibit density gradients.  In this paper, we present multislit \\emph{HST} STIS observations of 14 low-luminosity AGNs, including 13 LINERs and 1 Seyfert. We use the high spatial resolution of STIS to map out the 2D kinematic structure of the NLR within the inner $\\sim 100$ pc, and search for velocity fields that show signs of regular rotation for which gas dynamical modeling may be performed. The subarcsecond resolution STIS data also allow us to detect electron-density gradients through the [\\ion{S}{2}] $\\lambda 6716/\\lambda 6731$ line ratio diagnostic. Finally, we examine how the [\\ion{N}{2}] line width varies as a function of aperture size over spatial scales ranging from those comparable to the black hole sphere of influence to scales dominated by the bulge. ",
        "conclusions": "\\label{sec:conclusions}  We have used STIS data to measure detailed velocity, velocity dispersion, flux, and line-intensity ratio information within $\\sim 100$ pc from the nucleus of 14 nearby galaxies. All of the galaxies harbor low-luminosity AGNs, 13 of which are classified as LINERs and one of which is a Seyfert. We see a large diversity in the velocity fields of the 14 galaxies, where simple disk rotation is not a common state. Only a few galaxies have definite disk rotation, while some others have possible disk rotation but the emission-line S/N is too low to clearly determine the kinematic state of the gas. Some velocity fields are dominated by irregular motions, some show an overall rotation but with large random components, and others clearly exhibit outflows. Dynamical modeling of NGC 2911 and NGC 4594 to determine black hole masses may be possible and worthwhile, although chaotic and disorganized motion will limit the accuracy of mass measurements.  Through the [\\ion{S}{2}] $\\lambda 6716$/[\\ion{S}{2}] $\\lambda 6731$ diagnostic, we were able to detect significant electron-density gradients in four (NGC 1052, NGC 3998, NGC 4278, and NGC 4579) of the five LINERs having a sufficient S/N. Moreover, the electron-density gradients in these objects share a similar shape that can be described by a power law with a slope of $-0.60 \\pm {0.13}$. A similar increase near the nucleus was seen in the [\\ion{O}{1}]/[\\ion{S}{2}] line ratio for the first three of these galaxies. These results demonstrate that density gradients within the inner $\\sim 20$ pc are common in type 1 LINERs. This is consistent with expectations from photoionization models of LINERs and from ground-based detections of correlations between line width and \\emph{n}$_{\\mathrm{crit}}$ in LINERs.  We find that the density gradients persist in NGC 1052 and over a portion of the NLR in NGC 3227, even in the presence of strong outflows that clearly have a large impact on the NLR gas as seen through the disorganized velocity fields with multiple velocity components. We also see evidence for correlations between line width and \\emph{n}$_{\\mathrm{crit}}$ in most LINERs even on the 0\\farcs1 scales probed by STIS.  Measurements of the STIS spectra additionally provided a means to study the [\\ion{N}{2}] $\\lambda 6583$ narrow-line width variation with spectrograph aperture size. For most galaxies in our sample, the emission-line velocity dispersion peaks within the black hole sphere of influence, then decreases in larger apertures, approaching the bulge stellar velocity dispersion. This trend is consistent with, but does not uniquely imply, virial motion. The increasing line width with aperture size seen in NGC 1052 and NGC 3227 is due to the outflows that dominate the NLR kinematics. In NGC 3245, the increase is the result of the combination of the galaxy's fairly flat emission-line surface brightness profile, the galaxy's massive bulge, and the highly inclined gaseous disk. NGC 3245 demonstrates that even gas in pure disk rotation can exhibit an increasing line width as a function of aperture size.  While work by \\cite{Boroson_2005} and \\cite{Greene_Ho_2005} provide evidence for a connection between the Eddington ratio and kinematic disturbance in the NLR, which may reflect an underlying trend of outflow-dominated NLRs being more predominant at high AGN luminosities, our sample of galaxies is too small and does not span a sufficiently wide luminosity range to arrive at such a conclusion. Since such a relationship would be helpful in understanding AGN feedback and the effects on host-galaxy properties, further investigation, ideally with IFU data on a much larger sample of AGNs spanning a wide range of luminosities, is needed."
    },
    "0808/0808.2250_arXiv.txt": {
        "abstract": "Pre-main-sequence stars are observed to be surrounded by both accretion flows and some kind of wind or jet-like outflow. Recent work by Matt and Pudritz has suggested that if classical T Tauri stars exhibit stellar winds with mass loss rates about 0.1 times their accretion rates, the wind can carry away enough angular momentum to keep the stars from being spun up unrealistically by accretion. This paper presents a preliminary set of theoretical models of accretion-driven winds from the polar regions of T Tauri stars. These models are based on recently published self-consistent simulations of the Sun's coronal heating and wind acceleration. In addition to the convection-driven MHD turbulence (which dominates in the solar case), we add another source of wave energy at the photosphere that is driven by the impact of plasma in neighboring flux tubes undergoing magnetospheric accretion. This added energy, determined quantitatively from the far-field theory of MHD wave generation, is sufficient to produce T Tauri-like mass loss rates of at least 0.01 times the accretion rate. While still about an order of magnitude below the level required for efficient angular momentum removal, these are the first self-consistent models of T Tauri winds that agree reasonably well with a range of observational mass loss constraints. The youngest modeled stellar winds are supported by Alfv\\'{e}n wave pressure, they have low temperatures (``extended chromospheres''), and they are likely to be unstable to the formation of counterpropagating shocks and clumps far from the star. ",
        "introduction": "Our current state of knowledge about how stars and planets are formed comes from an intertwined web of observations (spanning the electromagnetic spectrum) and theoretical work. The early stages of low-mass star formation comprise a wide array of inferred physical processes, including disk accretion, various kinds of outflow, and magnetohydrodynamic (MHD) activity on time scales ranging from hours to millennia (see, e.g., Lada 1985; Bertout 1989; Appenzeller \\& Mundt 1989; Hartmann 2000; K\\\"{o}nigl \\& Pudritz 2000; McKee \\& Ostriker 2007; Shu et al.\\  2007). A key recurring issue is that there is a great deal of {\\em mutual interaction and feedback} between the star and its circumstellar environment. This interaction can determine how rapidly the star rotates, how active the star appears from radio to X-ray wavelengths, and how much mass and energy the star releases into its interplanetary medium.  A key example of circumstellar feedback is the magnetospheric accretion paradigm for classical T Tauri stars (Lynden-Bell \\& Pringle 1974; Uchida \\& Shibata 1984; Camenzind 1990; K\\\"{o}nigl 1991). Because of strong ($\\sim$1 kG) stellar magnetic fields, the evolving equatorial accretion disk does not penetrate to the stellar surface, but instead is stopped by the stellar magnetosphere. Accretion is thought to proceed via ballistic infall along magnetic flux tubes threading the inner disk, leading to shocks and hot spots on the surface. The primordial accretion disk is dissipated gradually as the star enters the weak-lined T Tauri star phase, with a likely transition to a protoplanetary dust/debris disk.  Throughout these stages, solar-mass stars are inferred to exhibit some kind of wind or jet-like outflow. There are several possible explanations of how and where the outflows arise, including extended disk winds, X-winds, impulsive (plasmoid-like) ejections, and ``true'' stellar winds (e.g., Paatz \\& Camenzind 1996; Calvet 1997; K\\\"{o}nigl \\& Pudritz 2000; Dupree et al.\\  2005; Edwards et al.\\  2006; Ferreira et al.\\  2006; G\\'{o}mez de Castro \\& Verdugo 2007; Cai et al.\\  2008). Whatever their origin, the outflows produce observational diagnostics that indicate mass loss rates exceeding the Sun's present mass loss rate by factors of $10^3$ to $10^6$. It is of some interest to evaluate how much of the observed outflow can be explained solely with {\\em stellar} winds, since these flows are locked to the star and thus are capable of removing angular momentum from the system. Recent work by Matt \\& Pudritz (2005, 2007, 2008) has suggested that if there is a stellar wind with a sustained mass loss rate about 10\\% of the accretion rate, the wind can carry away enough angular momentum to keep T Tauri stars from being spun up unrealistically by the accretion.  Despite many years of study, the dominant physical processes that accelerate winds from cool stars have not yet been identified conclusively. For many stars, the large-scale energetics of the system---i.e., the luminosity and the gravitational potential---seem to determine the overall magnitude of the mass loss (Reimers 1975; 1977; Schr\\\"{o}der \\& Cuntz 2005). Indeed, for the most luminous cool stars, radiation pressure seems to provide a direct causal link between the luminosity $L_{\\ast}$ and the mass loss rate $\\dot{M}_{\\rm wind}$ (e.g., Gail \\& Sedlmayr 1987; H\\\"{o}fner 2005). However, for young solar-mass stars, other mediating processes (such as coronal heating, waves, or time-variable magnetic ejections) are more likely to connect the properties of the stellar interior to the surrounding outflowing wind. For example, magnetohydrodynamic (MHD) waves have been studied for several decades as a likely way for energy to be transferred from late-type stars to their winds (Hartmann \\& MacGregor 1980; DeCampli 1981; Airapetian et al.\\  2000; Falceta-Gon\\c{c}alves et al.\\  2006; Suzuki 2007).  In parallel with the above work in improving our understanding of stellar winds, there has been a great deal of progress toward identifying and characterizing the processes that produce the {\\em Sun's} corona and wind. It seems increasingly clear that closed magnetic loops in the low solar corona are heated by small-scale, intermittent magnetic reconnection that is driven by the continual stressing of their footpoints by convective motions (e.g., Aschwanden 2006; Klimchuk 2006). The open field lines that connect the Sun to interplanetary space, though, appear to be energized by the dissipation of waves and turbulent motions that originate at the stellar surface (Tu \\& Marsch 1995; Cranmer 2002; Suzuki 2006; Kohl et al.\\  2006). Parker's (1958) classic paradigm of solar wind acceleration via gas pressure in a hot ($T \\sim 10^{6}$ K) corona still seems to be the dominant mechanism, though waves and turbulence have an impact as well. A recent self-consistent model of turbulence-driven coronal heating and solar wind acceleration has succeeded in reproducing a wide range of observations {\\em with no ad-hoc free parameters} (Cranmer et al.\\  2007). This progress on the solar front is a fruitful jumping-off point for a better understanding of the basic physics of winds and accretion in young stars.  The remainder of this paper is organized as follows. {\\S}~2 presents an overview of the general scenario of accretion-driven MHD waves that is proposed here to be important for driving T Tauri mass loss. In {\\S}~3 the detailed properties of an evolving solar-mass star are presented, including the fundamental stellar parameters, the accretion rate and disk geometry, and the properties of the clumped gas in the magnetospheric accretion streams. These clumped streams impact the stellar surface and create MHD waves that propagate horizontally to the launching points of stellar wind streams. {\\S}~4 describes how self-consistent models of these wind regions are implemented, and {\\S}~5 gives the results. Finally, {\\S}~6 contains a summary of the major results of this paper and a discussion of the implications these results may have on our wider understanding of low-mass star formation. ",
        "conclusions": "The primary aim of this paper has been to show how accretion-driven waves on the surfaces of T Tauri stars may help contribute to the strong rates of atmospheric heating and large mass loss rates inferred for these stars. The ZEPHYR code, which was originally developed to model the solar corona and solar wind (Cranmer et al.\\  2007), has been applied to the T Tauri stellar wind problem. A key aspect of the models presented above is that the only true free parameters are: (1) the properties of MHD waves injected at the photospheric lower boundary, and (2) the background magnetic geometry. Everything else (e.g., the radial dependence of the rates of chromospheric and coronal heating, the temperature structure of the atmosphere, and the wind speeds and mass fluxes) emerges naturally from the modeling process.  For solar-mass T Tauri stars, time-steady stellar winds were found to be supportable for all ages older than about 0.45 Myr, with accretion rates less than $7 \\times 10^{-8}$ $M_{\\odot}$ yr$^{-1}$ driving mass loss rates less than $4 \\times 10^{-10}$ $M_{\\odot}$ yr$^{-1}$. Still younger T Tauri stars (i.e., ages between about 13 kyr and 0.45 Myr) may exhibit time-variable winds with mass loss rates extending up several more orders of magnitude to $\\sim 2 \\times 10^{-8}$ $M_{\\odot}$ yr$^{-1}$. The transition between time-steady and variable winds occurs when the critical point of the flow migrates far enough past the Alfv\\'{e}n point (at which the wind speed equals the Alfv\\'{e}n speed) such that the Alfv\\'{e}n wave amplitude begins to decline rapidly with increasing radius. When this happens, the outward wave-pressure acceleration is quickly ``choked off;'' i.e., parcels of gas that make it past the critical point cannot be accelerated to infinity, and stochastic collisions between upflowing and downflowing parcels must begin to occur.  The maximum wind efficiency ratio $\\dot{M}_{\\rm wind}/\\dot{M}_{\\rm acc}$ for the T Tauri models computed here was approximately 1.4\\%, computed for ages of order 0.1 Myr. This is somewhat smaller than the values of order 10\\% required by Matt \\& Pudritz (2005, 2007, 2008) to remove enough angular momentum from the young solar system to match present-day conditions. It is also well below the observational ratios derived by, e.g., Hartigan et al.\\  (1995, 2004), which can reach up to 20\\% for T Tauri stars between 1 and 10 Myr old (see Fig.~12). These higher ratios, though, may be the product of both stellar winds and disk winds (possibly even dominated by the disk wind component). Additionally, it is possible that future observational analysis will result in these empirical ratios being revised {\\em upward} with more accurate (lower) values of $\\dot{M}_{\\rm acc}$ (S.\\  Edwards 2008, private communication).  The accretion-driven solutions for $\\dot{M}_{\\rm wind}$ depend crucially on the properties of the waves in the polar regions. It is important to note that the calculation of MHD wave properties was based on several assumptions that should be examined in more detail. The relatively low MHD wave efficiency used in equation (\\ref{eq:EA}) is an approximation based on the limiting case of waves being far from the impact site. A more realistic model would have to contain additional information about both the nonlinearities of the waves themselves and the vertical atmospheric structure through which the waves propagate. It seems likely that a better treatment of the wave generation would lead to larger wave energies at the poles. On the other hand, the assumption that the waves do not damp significantly between their generation point and the polar wind regions may have led to an assumed wave energy that is too high. It is unclear how strong the waves will be in a model that takes account of all of the above effects. In any case, the ZEPHYR results presented here are the first self-consistent models of T Tauri stellar winds that produce wind efficiency ratios that even get into the right ``ballpark'' of the angular momentum requirements.  Additional improvements in the models are needed to make further progress. For example, the effects of stellar rotation should be included, both in the explicit wind dynamics (e.g., ``magneto-centrifugal driving,'' as recently applied by Holzwarth \\& Jardine 2007) and in the modified subsurface convective activity that is likely to affect photospheric amplitudes of waves and turbulence. Young and rapidly rotating stars are likely to have qualitatively different convective motions than are evident in standard (nonrotating, mixing-length) models. It is uncertain, though, whether rapid rotation gives rise to larger (K\\\"{a}pyl\\\"{a} et al.\\  2007; Brown et al. 2007; Ballot et al. 2007) or smaller (Chan 2007) convection eddy velocities at the latitudes of interest for T Tauri stellar winds. In any case, rapid rotation can also increase the buoyancy of subsurface magnetic flux elements, leading to a higher rate of flux emergence (Holzwarth 2007). Also, the lower gravities of T Tauri stars may give rise to a larger fraction of the convective velocity reaching the surface as wave energy (e.g., Renzini et al.\\  1977), or the convection may even penetrate directly into the photosphere (Montalb\\'{a}n et al.\\ 2004).  Future work must involve not only increased physical realism for the models, but also quantitative comparisons with observations. The methodology outlined in this paper should be applied to a set of real stars, rather than to the idealized evolutionary sequence of representative parameters. Measured stellar masses, radii, and $T_{\\rm eff}$ values, as well as accretion rates, magnetic field strengths, and hot spot filling factors ($\\delta$), should be used as constraints on individual calculations for the stellar wind properties. It should also be possible to use measured three-dimensional magnetic fields (e.g., Donati et al.\\  2007; Jardine et al.\\  2008) to more accurately map out the patterns of accretion stream footpoints, wave fluxes, and the flux tubes in which stellar winds are accelerated.  This work helps to accomplish the larger goal of understanding the physics responsible for low-mass stellar outflows and the feedback between accretion disks, winds, and stellar magnetic fields. In addition, there are links to more interdisciplinary studies of how stars affect objects in young solar systems. For example, the coronal activity and wind of the young Sun is likely to have created many of the observed abundance and radionuclide patterns in early meteorites (Lee et al.\\  1998) and possibly affected the Earth's atmospheric chemistry to the point of influencing the development of life (see, e.g., Ribas et al.\\  2005; G\\\"{u}del 2007; Cuntz et al.\\  2008). The identification of key physical processes in young stellar winds is important not only for understanding stellar and planetary evolution, but also for being able to model and predict the present-day impacts of solar variability and ``space weather'' (e.g., Feynman \\& Gabriel 2000; Cole 2003)."
    },
    "0808/0808.2585_arXiv.txt": {
        "abstract": "The evolution of star clusters in the Magellanic Clouds has been the subject of significant recent controversy, particularly regarding the importance and length of the earliest, largely mass-independent disruption phase (referred to as `infant mortality'). Here, we take a fresh approach to the problem, using a large, independent, and homogeneous data set of $UBVR$ imaging observations, from which we obtain the cluster age and mass distributions in both the Large and Small Magelanic Clouds (LMC, SMC) in a self-consistent manner. We conclude that the (optically selected) SMC star cluster population has undergone at most $\\sim$30\\% (1$\\sigma$) infant mortality between the age range from about 3--10 Myr, to that of approximately 40--160 Myr. We rule out a 90\\% cluster mortality rate per decade of age (for the full age range up to $10^9$ yr) at a $>$6$\\sigma$ level. Using a simple approach, we derive a `characteristic' cluster disruption time-scale for the cluster population in the LMC that implies that we are observing the {\\it initial} cluster mass function (CMF). Preliminary results suggest that the LMC cluster population may be affected by $<10$\\% infant mortality. ",
        "introduction": "One of the most important diagnostics used to infer the formation history, and to follow the evolution of an entire star cluster population is the `cluster mass function' (CMF; i.e., the number of clusters per constant logarithmic cluster mass interval, ${\\rm d}N/{\\rm d}\\log m_{\\rm cl}$). The {\\it initial} cluster mass function (ICMF) is of particular importance. The debate regarding the shape of the ICMF, and of the CMF in general, is presently very much alive, both observationally and theoretically. This is so because it bears on the very essence of the star-forming processes, as well as on the formation, assembly history, and evolution of the clusters' host galaxies on cosmological time-scales. Yet, the observable at hand is the cluster {\\it luminosity} function (CLF; i.e., the number of objects per unit magnitude, ${\\rm d}N/{\\rm d}M_V$).  The discovery of star clusters with the high luminosities and the compact sizes expected for (old) globular clusters (GCs) at young ages facilitated by the {\\sl Hubble Space Telescope (HST)} has prompted renewed interest in the evolution of the CLF (and CMF) of massive star clusters. Starting with the seminal work by Elson \\& Fall (1985) on the young Large Magellanic Cloud (LMC) cluster system (with ages $\\lesssim 2 \\times 10^9$ yr), an ever increasing body of evidence seems to imply that the CLFs of young massive clusters (YMCs) are well described by a power law of the form ${\\rm d}N \\propto L^{1+{\\alpha}} {\\rm d}\\log L$, equivalent to a cluster luminosity spectrum ${\\rm d}N \\propto L^{\\alpha} {\\rm d} L$, with a spectral index $-2 \\lesssim \\alpha \\lesssim -1.5$ (e.g., Whitmore \\& Schweizer 1995; Elmegreen \\& Efremov 1997; Miller et al. 1997; Whitmore et al. 1999; Whitmore et al. 2002; Bik et al. 2003; de Grijs et al. 2003; Hunter et al. 2003; Lee \\& Lee 2005; see also Elmegreen 2002). Since the spectral index, $\\alpha$, of the observed CLFs resembles the slope of the high-mass regime of the (lognormal) old GC mass spectrum ($\\alpha \\sim -2$; McLaughlin 1994), this observational evidence has led to the popular theoretical prediction that not only a power law, but {\\it any} initial CLF (and CMF) will be rapidly transformed into a lognormal distribution because of (i) stellar evolutionary fading of the lowest-luminosity (and therefore lowest-mass, for a given age) objects to below the detection limit; and (ii) disruption of the low-mass clusters due to both interactions with the gravitational field of the host galaxy, and internal two-body relaxation effects leading to enhanced cluster evaporation (e.g., Elmegreen \\& Efremov 1997; Gnedin \\& Ostriker 1997; Ostriker \\& Gnedin 1997; Fall \\& Zhang 2001; Prieto \\& Gnedin 2007).  However, because of observational selection effects it is often impossible to probe the CLFs of YMC systems to the depth required to fully reveal any useful evolutionary signatures. As such, the young populous cluster systems in the Magellanic Clouds are as yet the best available calibrators for the canonical young CLFs that form the basis of most theoretical attempts to explain the evolution of the CLF and CMF.\\footnote{We point out, however, that with the latest {\\sl HST}/Advanced Camera for Surveys observations of the Antennae interacting galaxies, we are now finally getting to the point where the young (power-law) CLF shape appears to be confirmed independently down to sufficient photometric depths and for larger samples containing more massive clusters than in the Magellanic Clouds (B. Rothberg, priv. comm.).} It is therefore of paramount importance to understand the Magellanic Cloud cluster systems in detail. ",
        "conclusions": ""
    },
    "0808/0808.3898_arXiv.txt": {
        "abstract": "Cosmological simulations consistently predict specific properties of dark matter halos, but these have not yet led to a physical understanding that is generally accepted. This is especially true for the central regions of these structures. Recently two major themes have emerged. In one, the dark matter halo is primarily a result of the sequential accretion of primordial structure (ie `Nature'); while in the other, dynamical relaxation (ie `Nurture') dominates at least in the central regions. Some relaxation is however required in either mechanism. In this paper we accept the recently established scale-free  sub-structure of halos as an essential part of  both mechanisms. Consequently; a simple model for the central relaxation based on a self-similar cascade of tidal interactions, is contrasted with a model based on the accretion of adiabatically self-similar, primordial structure.  We conclude that a weak form of this relaxation is present in the simulations, but that is normally described as the radial orbit instability. ",
        "introduction": "In Henriksen (2006b, H06; 2007, H07) a theory of dark matter relaxation has been proposed that is based on a temporally convergent series solution for the Distribution Function (true phase-space density or DF). The relaxation is effected in a non-mechanistic fashion by maximizing the local Boltzmann function calculated from this DF. This procedure determines all of the parameters in the DF and allows the density distribution $\\rho(r)$ and the pseudo phase-space density $\\phi(r)\\equiv \\rho(r)/\\sigma(r)^3$ to be calculated, as well as any other quantity such as specific angular momentum. Studying these results one finds that, starting from the position and corresponding slopes found in the simulations, either the density flattens rapidly or the pseudo density steepens rapidly within the next decade of smaller radius. There is at present no convincing evidence for either of these trends in the simulations, although unfortunately this region is at or near the current resolution limit. This paper attempts to understand the proposed relaxation more intuitively by suggesting a simple mechanistic model. Such a view may be briefly summarized as `nurture'.  In Salvador-Sol\\'e et al. (2007, S07) a rather different explanation for the simulated structures is proposed. In effect one uses the Press-Schechter (1974) formalism as expressed by Lacy and Cole (1993) to calculate the instantaneous merger rate. Since most of the halo mass is added by many small objects, this is approximated by being smooth. This smoothness allows one to accrete the dark matter halo from post-recombination large scale structure in a sort of `layer cake' fashion, where each layer may be deduced from the  accretion flow at the epoch when it was added. The innermost layers correspond to the earliest times so that `inside-out' growth is established.  The predicted density profile compares well to NFW (Navarro, Frenk and White, 1997; Navarro et al. 2004) profiles for various masses over appropriate ranges and even better over the whole range of scales to an Einasto (Navarro et al., 2004) or S\\'ersic profile (e.g. S07). The latter profile has a finite density at the centre but an infinite slope. Similar comparisons may be made for velocity dispersion and angular momentum (Gonzales-Casado et al., 2007). Such a view may be briefly summarized as `Nature', although there is also a kind of relaxation that leads to adiabatic self-similarity (e.g. Henriksen 2006b, H07) in the asumption of ``smoothness''.  Unless it is also hidden in the asumption of `smoothness', the latter picture leaves little room for `thermodynamic relaxation' (ie maximum entropy) as an element in the simulations. Nevertheless, whether or not such relaxation is present in reality in order to explain possible  density cores remains an open question. Moreover the remarkable results of the `Via Lactea' simulation (Diemand, Kuhlen and Madau, 2006-DKM06; Madau, Diemand, and Kuhlen, 2008-MDK08) show elaborate sub-structure that is characterized by a definite mass spectrum. Such a hierarchy should be interacting tidally and by dynamical friction and so it provides a mechanism for relaxation (see e.g. the discussions in El-Zant et al., 2004, H07, H06 and  Henriksen, 2006a).  One indication of the possible presence of such relaxation has been provided by Hoffman et al. (2007). They observe that the halo surface density at the NFW scale radius $r_s$  (that is $\\rho_sr_s$) remains constant during the evolution of an individual halo. This together with virialization within $r_s$ allows them to conclude that $\\phi_s\\propto r_s^{-5/2}$. Some time ago, virialization together with strong dynamical inter-scale coupling were shown to be equivalent to constant surface density and virialization throughout a cascade of structures. This was in a completely different context (Henriksen and Turner, 1984; Henriksen, 1991 and references therein) where it was studied in the context of star formation. It was proposed that a hierarchy of molecular clouds slowly evolved by collisional and tidal interactions into the stellar Initial Mass Function (IMF) as the result of a kind of `ballistic turbulence'. The observation of Hoffman et al (2007) together with the power law substructure (DKM06) suggests a similar argument may apply to dark matter halos.  In the stellar case it has been difficult to find any evidence for such an initial fragmented state, and so ironically  the picture may be more relevant to dark halos than it is to luminous stars!  However Hoffman et al. have not applied the idea of a dynamical cascade to all scales below $r_s$. This will be the subject of the next section. In the third section we give a naive calculation that imitates the layer cake approach of (S07), by simply assuming adiabatic self-similarity as dictated by the primordial perturbation spectrum. Finally in the conclusions we discuss the liklihood of either sort of relaxation. We conclude  that a weak form of cascade relaxation is present in the simulations that is not distinguishable from the radial orbit instability. We conclude further that the pure cascade evolution should lead to a sharp break in the  density and pseudo-density trends as simulated  currently in the next decade or so of resolution. Otherwise we are left only with the weak form of relaxation that is described in this paper as either `nature' or adiabatic self-similarity, and which seems to be equivalent to the radial orbit instability. ",
        "conclusions": "The preceding sections allow us to conclude that the simulated dark matter halos are, up to the presently available resolution, consistent with the primarily `nature' (S07) adiabatically self-similar smooth growth from the CDM spectrum (e.g. \\ref{fig:nature}). In our  semi-analytic treatment of the previous section, we imitate this growth using the notion of adiabatic  self-similarity (Henriksen, 2006b) to connect it to the CDM spectrum. In S07 however this was done using the assumption of smooth inside-out or `layer-cake' growth according to the Lacey and Cole (1993) prescription. This employs a Press-Schecter (1974) type argument for the CDM spectrum, together with the statistical treatment of Bond et al. (1991). We conclude that this may be done more simply using the notion of adiabatic self-similarity, which does  involve a measure of relaxation. Moreover, although it happens slowly, this approach predicts a central core ($n\\rightarrow -3$) rather than a cusp.  What we have termed variously `nurture',`thermodynamic' or `maximum entropy' relaxation based on a cascade of interacting structure (H07 and above) agrees on the qualitative trends with the `nature' view . However it evolves at small scales  towards a flat density and a steep pseudo-density more rapidly than is found either in the `nature' discussion or in the simulations if it is assumed to extend close to $r_s$ (\\ref{fig:evolve}). This discrepancy is less pronounced when the relaxation extends only to smaller scales, however (e.g. figure(\\ref{fig:K2evolve})). If we believe that cascade relaxation is playing a r\\^ole in the halo evolution, then  equation (\\ref{beta}) and the initial $\\beta\\approx -1$ tell us  that the relaxation  should indeed be most effective at small scales initially. In time, as $\\beta$ becomes positive, the relaxation should move to large scales in agreement with equation(\\ref{tJ}).  The only relaxation that is visible in the simulations at present is in the density profile within two decades or so of $r_s$.  The phase space pseudo-density shows no relaxation over this same range. This behaviour can be fitted over  this limited range by either of the above relaxation mechanisms, but  both mechanisms predict  stronger small scale relaxation for which there is as yet no evidence in the simulations.  One wonders whether the adiabatic self-similar relaxation that seems to fit  the current simulations well, is in fact due to a form of cascade relaxation. Recently a strong case has been made  (MacMillan, Widrow and Henriksen, 2006 (MWH06) and references therein) that the relaxation in this region near $r_s$ is due to the Radial Orbit Instability (ROI), so we should reconsider this in the present context.  The onset of this instability was found in the MWH06 paper to coincide with the development of the mean square specific angular momentum at $r$ into the Keplerian form $\\propto GM(r)r$. Let us suppose that the angular momentum perpendicular to the radial direction takes the form $\\ell_r^2\\sigma^2(\\ell_r)$, where we use the same notation as in equation (\\ref{coupling}). Then the onset of the ROI is marked by the condition $\\ell_r^2\\sigma^2(\\ell_r)\\approx GM(r)r$ which on taking $\\ell_r\\approx r$ becomes essentially the cascade coupling condition of equation (\\ref{coupling}). The explanation offered in the MWH06 paper is in terms of a resonant interaction between a `bar-like' density perturbation and `particle' (perhaps a sub-structure) orbits. In the present context we may see this as the largest asymmetric sub-structure at $r$. We  expect it to be composed of sub-structures and to be interacting with such objects in the environment.  It s interesting to note incidentally that this interaction is a kind of coupling between the radial infall and the transverse orbital motion, which delivers free-fall energy to the top of the cascade. The Cascade is really  gravitationally driven turbulence. Ome might say that the mode $k\\approx 1/r$ parallel to the infall is coupling to the transverse mode $k_r\\approx 1/\\ell_r$ perpendicular to the radial infall to produce this turbulence.  Of course these are all  speculative, order of magnitude, arguments. Nevertheless we are inclined to suggest that the simulated halos, hence also the adiabatic self-similarity  relaxation that works well for the current simulations, is due to the ROI. This is in turn a mechanism for the conversion of radial infall into transverse orbital motion at $r$. This transverse mode represents the top of the cascade at each radius. We thus identify the ROI with a weak form of cascade relaxation at large scales. This weakness is in accord with equations (\\ref{beta}) and (\\ref{tJ}) that both predict ($\\beta<0$ in the first case) stronger relaxation at small scales. The stronger cascade is expected to be at scales below the resolution of the current simulations.  We therefore conclude that cascade evolution may be at work both in reality and in the simulations. We conclude tentatively  that `Nature' type adiabatic self-similarity is the weak form of the cascade relaxation and that it is equivalent to the ROI. Our predictions that would confirm the presence of cascade evolution are that there is a halo density core rather than a cusp and that the pseudo-density power law should break at higher resolution in both scale and mass. In the Via Lactea run this region is just inside their reported convergence radius at about 1kpc. There may be weak evidence in figure 1 of DKM06 that a density flattening is not excluded.  Should this density and pseudo-density behaviour not appear down to say 100pc in a Milky Way type halo, then the strong cascade relaxation is not present in the simulations. It is then probably not present in reality, unless for some reason the necessary interactions at a distance between sub-structures is discriminated against in the simulations. DKM06 do observe that at lower resolution than that of the Via Lactea run, the small sub-halos are poorly resolved."
    },
    "0808/0808.3917_arXiv.txt": {
        "abstract": "{It has been reported that exoplanet host stars are lithium depleted compared to solar-type stars without detected massive planets. } {We investigate whether enhanced lithium depletion in exoplanet host stars may result from their rotational history. } {We have developed rotational evolution models for slow and fast solar-type rotators from the pre-main sequence (PMS) to the age of the Sun and compare them to the distribution of rotational periods observed for solar-type stars between 1~Myr and 5~Gyr.} {We show that slow rotators develop a high degree of differential rotation between the radiative core and the convective envelope, while fast rotators evolve with little core-envelope decoupling. We suggest that strong differential rotation at the base of the convective envelope is responsible for enhanced lithium depletion in slow rotators. } {We conclude that lithium-depleted exoplanet host stars were slow rotators on the zero-age main sequence (ZAMS) and argue that slow rotation results from a long lasting star-disk interaction during the PMS. Altogether, this suggests that long-lived disks ($\\geq$ 5~Myr) may be a necessary condition for massive planet formation/migration.} ",
        "introduction": "As the discovery rate of exoplanets has steadily increased over the years, now reaching nearly 300 detections, interest has grown in the properties of their host stars (Santos et al. 2003; Kashyap et al. 2008). Israelian et al. (2004) report that solar-type stars with massive planets are more lithium-depleted than their siblings without detected massive planets. This result has recently been confirmed by Gonzalez (2008). Lithium over-depletion in massive exoplanet hosts appears to be a generic feature over a restricted T$_{eff}$ range from 5800 to 5950~K, independent of the planet's orbital properties. Metallicity effects and/or the early accretion of planetesimals have been ruled out as the origin of lithium over-depletion in exoplanet hosts (Castro et al. 2008). Instead, it has been suggested that enhanced lithium burning could stem from the tidal effect of the giant planet on the host star, as it induces rotationally-driven mixing in the stellar interior (Israelian et al. 2004). However, the mass of giant exoplanets is much lower than the mass of the convective envelope of solar-type stars ($\\geq$0.02$M_\\odot$). In addition, enhanced lithium depletion is seen for stars with giant planets that span a wide range of semi-major axes, up to about 2~AU. Whether the tidal interaction between a distant massive planet and the host star can be strong enough to induce rotationally-driven mixing remains to be investigated.  Rotationally-driven mixing may instead be related to the rotational history of the star (Takeda et al. 2007; Gonzalez 2008). We revisit here the link between lithium depletion and the rotational history of solar-type stars, and discuss its relationship with massive planet formation. In Section 2, we present models for the rotational evolution of solar-type stars from the early PMS to the age of the Sun. We use these models to fit the most recent measurements of rotational periods obtained for solar-type stars in star forming regions and young open clusters. We discuss the properties of 2 extreme models: one reproduces the evolution of fast rotators over time, and the other the evolution of slow rotators. We find that slow rotators develop a high degree of differential rotation between the inner radiative zone and the outer convective envelope, while fast rotators evolve with little core-envelope decoupling. From these results, we investigate in Section 3 the evolution of lithium abundances in slow and fast rotators. We suggest that the large velocity shear at the core/envelope interface in slow rotators leads to more efficient Li-depletion. In Section 4, we relate the Li-depletion pattern of massive exoplanet host stars to their rotational history. We argue that both are dictated by their initial angular momentum and the lifetime of their protoplanetary disk. We conclude that long-lived disk may be a necessary condition for massive planet formation/migration around young solar-type stars. ",
        "conclusions": "Based on what we currently know of the rotational properties of young stars, of the lithium depletion process in stellar interiors, and of the angular momentum evolution of solar-type stars, it seems likely that the lithium-depleted content of massive exoplanet host stars is a sequel to their specific rotational history. This history is predominantly dictated by star-disk interaction during the PMS. Using simple rotational models, it is suggested here that the long disk lifetimes of initially slow rotators lead to the development of a large velocity shear at the base of the convective zone. The velocity gradient in turn triggers hydrodynamical instabilities responsible for enhanced lithium burning on PMS and MS evolution timescales. Rotationally-driven lithium depletion in exoplanet host stars can thus be at least qualitatively accounted for by assuming PMS disk lifetimes of 5-10 Myr. Such long-lived disks may be a necessary condition for planet formation and/or migration around young solar-type stars, at least for the class of giant exoplanets detected so far (Ida \\& Lin 2008).  Admittedly, the scenario outlined here is based on rotational models that do not incorporate a detailed physical description of the processes at work. The conclusion therefore remains qualitative and has to be ascertained by the development of more sophisticated angular momentum evolution models. Most important, the transport processes that lead to a strong core-envelope coupling in fast rotators and a weak one in slow rotators still have to be identified. While this goal is well beyond the scope of the present study, we hope that the simplified models presented here may provide useful guidelines for further refinements."
    },
    "0808/0808.3310_arXiv.txt": {
        "abstract": "The asymmetric time dependence and various statistical properties of polarity reversals of the Earth's magnetic field are utilized to infer some of the most essential parameters of the geodynamo, among them the effective (turbulent) magnetic diffusivity, the degree of supercriticality, and the relative strength of the periodic forcing which is believed to result from the Milankovic cycle of the Earth's orbit eccentricity. A time-stepped spherically symmetric $\\alpha^2$-dynamo model is used as the kernel of an inverse problem solver in form of a downhill simplex method which converges to solutions that yield a stunning correspondence with paleomagnetic data. ",
        "introduction": "The hydromagnetic dynamo in the Earth's outer core converts gravitational and thermal energy into magnetic energy \\cite{MERRILL}. One of the most impressive feature of the geomagnetic field is the irregular occurrence of polarity reversals. Averaged over the last few million years the mean rate of reversals is approximately 4-5 per Myr, although the last reversal occurred approximately 780000 years ago. At least two, but very likely  three \\cite{GALLET,COURTILLOT} superchrons have been identified as ''quiet'' periods of some tens of millions of years showing no reversal at all.  Knowledge on ancient magnetic field data is mainly obtained from paleomagnetic measurements of permanent magnetization from (frozen) lava and sedimentary rocks. However, appropriate paleomagnetic sites are rare, unevenly distributed across the Earth's surface and, furthermore,  actual dating methods allow only a rather rough time-resolution. Hence only few reversal characteristics can be evaluated as robust \\cite{MERRILL2}. One of the commonly accepted features of reversals is their pronounced temporal asymmetry  with the initial decay of the dipole being much slower than the subsequent recreation of the dipole with opposite polarity \\cite{VALET}.  Recent numerical simulations have been successful in reproducing not only the dominance of the axial dipole and the spectrum of the geomagnetic field, but also the irregular occurrence of polarity reversals \\cite{ROGLA,WICHTOLSON,AUBERT}. Polarity reversals were also observed in one \\cite{BERHANU} of the recent liquid sodium dynamo experiments which have flourished during the last decade \\cite{RMP}.  It is important to note, however, that neither in simulations nor in experiments it is possible to accommodate all dimensionless parameters of the geodynamo \\cite{GUBBINS,ZAMM}, and many of them are not even well known \\cite{ROGLA2}.  A complementary way to acquire knowledge about the geodynamo is to use available magnetic field data for constraining the source of dynamo action. Most famous among those attempts is the frozen-flux approximation \\cite{ROSCO,HOLMEOLSEN} which allows to infer the tangential flow at the core-mantle-boundary from secular variation of the geomagnetic field. Going beyond this frozen-flux approximation and trying to infer properties of the geodynamo in the depth of the Earth's core has been less successful so far. Such trials to  ''look inside the dynamo'' by utilizing spectral data  worked nicely for simplified dynamo models \\cite{AN,PEPI,PRE} but are  hardly applicable for real world dynamos.  With this sobering experience in mind, in the present paper we undertake a rather uncommon attempt to constrain some of the most significant parameters of the geodynamo by various characteristics of paleomagnetic reversal records. Most prominent among those characteristics is the above mentioned temporal asymmetry  of reversals. Two further characteristics reflect some sort of ordering in the otherwise irregular reversal sequence. The first one is the clustering property of reversals which was discovered only recently \\cite{CARBONE,LUCA}. Clustering of reversals is manifested in an enhanced probability of a consecutive reversal shortly after a first reversal has occurred,  which results in a deviation from the Poisson distribution that would hold for uncorrelated events. The second one is the appearance of a $\\sim$100 kyr periodicity in the distribution of the residence times (RTD) between reversals which is believed to result from the Milankovic cycle of the Earth's orbit eccentricity \\cite{CONSOLINI,LORITO}.   Based on these three input features we will examine in this paper a simplifying $\\alpha^2$-model of the geodynamo for which we estimate the degree of supercriticality, the noise level, the relative strength of the periodic forcing, and the effective (turbulent) magnetic diffusivity of the Earth's outer core. Actually, a few dependencies on individual parameters were already published in preceeding papers. In \\cite{EPSL} the dependence of typical time scales on the supercriticality of the dynamo was studied, in \\cite{LUCA} some dependencies of the clustering property on the supercriticality and the noise level were shown, and in \\cite{EPJB} the influence of the diffusion time scale on the RTD between reversals was touched upon. What is new in the present paper is that we take all three reversal features together and  try to infer from them some essential parameters of the geodynamo.  We will start with a presentation of the forward dynamo problem for which we will use a rather simple, spherically symmetric mean-field dynamo of the $\\alpha^2$ type. This simple model had turned out to be  quite helpful for understanding the basic principle of the reversal process as a noise-induced relaxation-oscillation in the vicinity of an exceptional point of the spectrum of the non-selfadjoint dynamo operator \\cite{EPSL,PRL,GAFD}. This exceptional point, at which two real eigenvalues coalesce and continue as a complex conjugated pair of eigenvalues,  is associated with a nearby local maximum of the growth rate situated at a slightly lower magnetic Reynolds number. It is the negative slope of the growth rate curve between this local maximum and the exceptional point that makes stationary dynamos vulnerable to  noise. Then, the instantaneous eigenvalue is driven towards the exceptional point and beyond into the oscillatory branch where the sign change of the dipole polarity happens.  After having delineated the simplified mean-field dynamo model we will present the solution method for the inverse problem and the main results.  The paper concludes with a summary and a speculation on the possible consequences of our findings for the general understanding and the numerical simulations of the geodynamo. ",
        "conclusions": "With version 2, we have obtained the best reproduction of the paleomagnetic input data for a 10 times supercritical dynamo, a relative strength of the periodic forcing of some 15 per cent, and an effective magnetic diffusion time of approximately 65 kyr which is by a factor 3.5 smaller than the value that would result from the molecular conductivity. The latter is perhaps the most important result of the inversion, for the following reason: The conductivity enhancing effect of turbulence is a subject of ongoing debate. Only very few estimations exists for the turbulent diffusivities within the Earth's outer core. However, their values are of extraordinary importance for the key parameters of the geodynamo (e.g. Ekman number or the magnetic Prandtl number) as well as for the estimation of turbulent transport properties and the destructive influence of turbulence on magnetic field generation through turbulent field diffusion. The molecular values of the dimensionless parameters are extremely small and cannot be realized in simulations so that usually enhanced values are applied that shall resemble the effective (i.e. turbulent) quantities. A more detailed information on the turbulent diffusivities therefore delivers important information on the significance of the parameter space accessible in global MHD simulations for the geodynamo. Furthermore, reality is more complicated, as it is very likely that a conductivity reduction due to turbulence would be anisotropic, because the small scale turbulence in a fast rotating object like the Earth is subject to preferred directions parallel to the rotation axis (and also along the dominating field component) so that the convection cells (that determine the turbulent transport of physical quantities) are oriented parallel to the rotation axis and elongated along the magnetic field. It is well known that an anisotropic conductivity could have a tremendous effect on the selection between equatorial and axial dipole solution \\cite{TILGNER,ELSTNER}. Roughly speaking, axially aligned rotating columns tend to decrease the conductivity for horizontal currents less  than the conductivity for vertical currents, and this effect leads to a preference of the axial dipole compared to the equatorial dipole. This could be an extremely important point since the conclusions of many geodynamo related papers rely  on an isotropic conductivity when looking for criteria for the selection of axial or equatorial dipoles.    The stunning agreement of paleomagnetic and numerical reversal characteristics as shown in figures 1, 2 and 3 gives support to our hypothesis that reversals are indeed noise triggered relaxation oscillations in the vicinity of an exceptional point of the spectrum of the dynamo operator. In this respect it is important to note that in particular the time asymmetry and the clustering property are intrinsic and robust properties of the model that appear for very wide regions of parameter (if not for all). By solving the inverse problem we have not ''produced'' them, but have only fine-tuned the dynamo parameters to fit optimally the paleomagnetic data.  We have carefully tried not to over-interpret our simple model by  focusing  only on those parameters to be determined, and those functionals to be minimized, that refer to the  temporal properties of reversal sequences, and not to any spatial features. This makes us optimistic that the results will prove robust when inversions of this kind will later be repeated using more realistic dynamo models. Given that one downhill simplex run for our simple model takes already one week, one can imagine that corresponding runs with better dynamo models will lead to significant computational costs.       \\ack This work was supported by Deutsche Forschungsgemeinschaft in frame of SFB 609."
    },
    "0808/0808.1301_arXiv.txt": {
        "abstract": "Motivated by the possibility that the fundamental ``constants'' of nature could vary with time, this paper considers the long term evolution of white dwarf stars under the combined action of proton decay and variations in the gravitational constant. White dwarfs are thus used as a theoretical laboratory to study the effects of possible time variations, especially their implications for the future history of the universe.  More specifically, we consider the gravitational constant $G$ to vary according to the parametric relation $G = G_0 (1 + t/\\tG)^{-p}$, where the time scale $\\tG$ is the same order as the proton lifetime $t_{P}$.  We then study the long term fate and evolution of white dwarf stars.  This treatment begins when proton decay dominates the stellar luminosity, and ends when the star becomes optically thin to its internal radiation. ",
        "introduction": "One of the overarching but unresolved questions in physics is whether or not the fundamental constants are truly constant, or if they could vary with time. A related question is whether these constants could, in principle, have other values in far-away regions of space-time (i.e., in effectively other universes).  Current experiments place limits on the time variations of the fundamental constants, e.g., the gravitational constant $G$ and the fine-structure constant $\\aem$ (Uzan 2003 and references therein). These constraints can be expressed in terms of corresponding time scales, $\\tg = G / {\\dot G}$ and $\\tem = \\aem / {\\dot \\aem}$, which are found in the range $10^{10} - 10^{12}$~yr (Uzan 2003).  Although these time scales are safely longer than the current age of the universe ($t_0$ = 13.7 Gyr; e.g., Spergel et al. 2007), these values are much shorter than some time scales that are experimentally accessible; as one example, the proton lifetime has a measured lower bound of $10^{33}$ yr (Super-K 1999).  As a result, over the vast expanses of time available to the universe in the future, time variations in the physical constants could change the projected future history of the universe.  This paper addresses this issue by considering possible time variations of the gravitational constant and its corresponding effects on the long term evolution of white dwarf stars. In this context, we are using white dwarfs as a theoretical laboratory to consider the effects of time-varying $G$. This choice is justified because white dwarfs are among the simplest and hence most well understood stellar objects (starting with Chandrasekhar 1939) and because they play an important role in the future history of the universe (Adams \\& Laughlin 1997, hereafter AL97; see also Cikovic 2003, Dyson 1979, Islam 1977).  For example, almost all stars turn into white dwarfs after their nuclear burning phase, and hence a sizable fraction of the accessible baryonic content of the universe is locked up in white dwarfs after stellar evolution has run its course (for completeness, note that a sizable fraction of the baryons remain in the medium between galaxies in dusters, e.g. Nagamine \\& Loeb 2004).  To address this issue, we must define both the time scales and the functional form of the time variations under consideration.  Although an enormous variety of time variations are possible, we narrow the scope by allowing the gravitational constant $G$ to vary according to the parametric relation \\be G(t) =  G_0 (1 + t/\\tG)^{-p} \\, , \\ee where $G_0$ is the current value, and where $\\tG$ and $p$ are parameters.  This functional form is perhaps the simplest type of allowed time variation --- the value of $G$ reduces to the current value at the current cosmological epoch and decreases as a power-law in the long term future.  Note that current experimental limits imply that $\\tg = G/{\\dot G} > 10^{12}$ yr (e.g., Table IV of Uzan 2003; see also Barrow 1996) so that $\\tG \\gg t_0$. In the long term future, proton decay is one of the most important energy sources for white dwarfs. This problem thus has two time scales of interest: the proton lifetime $t_P$ and the time scale $\\tG$ for variations in $G$. In the limit $t_\\star \\gg t_P$, white dwarf evolution closely follows previous results obtained using a fixed gravitational constant (AL97, Adams et al. 1998; Dicus et al. 1982).  As a result, this paper primarily considers the limit $t_\\star \\ll t_P$, where white dwarf evolution is dominated by changes in gravity, and the case where $\\tG$ is comparable to $t_P$.  Although this work will focus on time variations in the gravitational constant $G$, for completeness we note that time variations in other physical quantities are possible.  For example, the masses of the fundamental particles or, equivalently, the ratio of electron to proton mass $\\mu = m_e / m_P$, could be time dependent (Calmet \\& Fritzsch 2006).  We note that the quark masses, which play a role in determining the proton mass, are actually the fundamental constants of the underlying theory; however, the ratio $\\mu$ has been experimentally constrained (see Uzan 2003 and references therein).  In this paper we consider the variation of $G$ by itself in order to isolate the effects of its possible time evolution.  However, fundamental theories (string theory, M theory) suggest that the forces are unified and hence may vary in strength in a coupled manner (Uzan 2003 provides a brief review of the possibilities).  Unfortunately, we do not have a definitive working theory of how such coupled time variations are expected to occur, and we leave such a more comprehensive treatment for future work.  This paper is organized as follows. In \\S 2, we outline a basic model for white dwarf structure, including our treatment of time variations in $G$ and proton decay. One interesting complication is that proton decay leads to a downward cascade for the atomic weights of the constituent nuclei; this issue is addressed in some detail.  In \\S 3 we outline the basic results of the paper, including the H-R diagrams for the long term evolution of these degenerate stars and their corresponding chemical evolution. We then conclude in \\S 4 with a summary and discussion of our results. ",
        "conclusions": "This paper has constructed a model for the structure and evolution of white dwarf stars under the action of both proton decay and time variations in the gravitational constant (section 2). This model includes variations in chemical composition (Figure 2 and section 2.3) which affects the equation of state.  The model also includes separate accounting for the inner degenerate regions (which obey the polytropic equation of state [1]) and the outer non-degenerate regions.  Our main findings can be summarized as follows:  In the absence of baryon decay, white dwarfs grow larger as the strength of gravity decreases (Figure 4). The radius reaches a maximum value when the entire star becomes non-degenerate; at this time, the radius is $\\sim 100$ times larger than its starting value.  With the inclusion of baryon decay, white dwarfs lose mass with time and follow tracks in the H-R diagram as shown in Figure \\ref{fig:HRC}. In all cases, after an initial transient phase, the stars grow dimmer and redder with time, and hence move to the lower right in the H-R diagram.  While the stars remain (primarily) degenerate, the tracks have a shallow slope, with $L \\propto T^{12/5}$.  After the stars lose enough mass to become primarily non-degenerate, the tracks become much steeper, with $L \\propto T^{12}$. These slopes would be exact in the absence of changes in the chemical composition; the chemical variations and the time variations in $G$ lead to small departures from this behavior.  During the degenerate regime of evolution, the radii of these white dwarf stars grow larger due to mass loss from proton decay and due to the weakening of gravity. During the later, non-degenerate phase, the radii become smaller. The maximum value of the radius occurs near the transition between the degenerate and non-degenerate regimes, with a value that is typically $\\sim 10$ times the starting radius (see Figure \\ref{fig:RvstC}).  In addition to understanding the long term fate and evolution of white dwarfs, one motivation for this work is to understand the larger issue of the future of the universe.  Previous projections of the future (e.g. AL97, Dyson 1979) are predicated on the assumption that the laws of physics are known and will not change with time.  One important issue is thus the question of whether or not the constants of nature change with time, and how such variations would change projections of our future history.  In the case of white dwarf evolution, we find that time variation leads to quantitative --- but not qualitative --- changes in the future timeline.  Time varying $G$ leads to variations in the exact mass, size, and composition of white dwarfs as they become non-degenerate rock-like objects.  However, the starting states and final states are essentially the same, so time variations in gravity do not lead to fundamental changes in the overall picture. White dwarfs start as degenerate objects with the current value of $G$, and hence are constrained to begin their evolution in nearly the same states.  At the other end of time, white dwarfs cease to act as stars when they become optically thin to their internal radiation.  At this epoch, the ``stars'' are essentially large rocks of Hydrogen ice, and their structure is non-degenerate and largely independent of gravity.  As a result, time variations in $G$ primarily affect the intermediate states, and this work shows that the modifications are relatively modest (see Figures \\ref{fig:rhovtC} - \\ref{fig:HRC})."
    },
    "0808/0808.1889_arXiv.txt": {
        "abstract": "Some early-type galaxies show a correlation between color and integrated magnitude among the brighter metal-poor globular clusters (GCs). This phenomenon, known as the blue tilt, implies a mass-metallicity relationship among these clusters. In this paper we show that self-enrichment in GCs can explain several aspects of the blue tilt, and discuss predictions of this scenario. ",
        "introduction": "Several recent photometric studies of extragalactic globular cluster (GC) systems have discovered a novel feature. The blue, metal-poor population of GCs shows a color-magnitude relation at bright magnitudes: the more luminous GCs are redder. This ``blue tilt\" was first noticed in the GC systems of massive elliptical (E) galaxies (Harris \\etal~2006; Strader \\etal~2006), but has since been seen in the Sa/S0 NGC 4594 (Spitler \\etal~2006), fainter early-type galaxies in Virgo (Mieske \\etal~2006), and even in normal spirals (Strader \\etal~2008, in preparation). Figure 1 shows the blue tilt in M87 (Strader \\etal~2006; see also Mieske \\etal~2006). The blue tilt is common but not ubiquitous---the massive E NGC 4472 does not have it, for example, and among galaxies that do show a tilt, the slope varies. There is no evidence for a tilt in the red, metal-rich subpopulation of GCs.  Many explanations for this phenomenon have been proposed. These include (i) self-enrichment, (ii) ``pollution\" of the GC sequence with objects such as dwarf galaxy nuclei or merged star clusters (e.g., Bekki \\etal~2007), and (iii) gravitational capture of metal-rich field stars by GCs (Mieske \\etal~2006). A number of other scenarios are discussed in Mieske \\etal~(2006) but are thought by them to be much less probable.  Option (ii) is considered by both Mieske \\etal~(2006) and Strader \\etal~(2006), who conclude that it is unlikely. While the nuclei of Virgo dwarfs show their own color-magnitude relation (Lotz \\etal~2001), the red mean colors of the luminous GCs are not due solely to the addition of bright red objects, but instead come from systematic shifts in the GC color distribution with magnitude. In addition, the number of nuclei needed is much larger than the handful of stripped nuclei that have been found in either the Virgo or Fornax clusters (see also discussion in \\S 4).  Except for a few objects, the sizes of the blue tilt GCs are generally small, inconsistent with the expectations of merged clusters. Merged clusters should be larger as the orbital energy of the binary cluster is converted into kinetic energy in the merger product. The best candidate for a merged cluster, NGC 1846 in the Large Magellanic Cloud (Mackey \\& Broby Nielsen 2007), has an unusually large projected half-light radius of $\\sim 15$ pc (D.~Mackey, private communication). While it seems likely that such exotic objects contribute to the GC systems of massive galaxies at some level, especially at the brightest magnitudes (Harris \\etal~2006; Strader \\etal~2006), they are unlikely to be primarily responsible for the blue tilt.  Option (iii) was investigated using collisional N-body simulations by Mieske \\& Baumgardt (2007). They report that field star capture is inefficient even in optimistic circumstances, and conclude that this is not the cause of the blue tilt.  Our working assumption is that self-enrichment remains the most likely explanation for the blue tilt. In this case the basic explanation is that the color-magnitude relation reflects a physical mass-metallicity relation for metal-poor GCs. As discussed in Strader \\etal~(2006), there are a variety of GC self-enrichment models and simulations in the literature (e.g., Morgan \\& Lake 1989; Brown \\etal~1991; Parmentier \\etal~1999; Parmentier \\& Gilmore 2001; Parmentier 2004; Recchi \\& Danziger 2005). Here we refer to models that focus on self-enrichment as the primary determinant of cluster metallicity, and not to the larger body of work on the origin of abundance anomalies in Galactic GCs that primarily involve light elements such as C, N, O, and Mg. Among GC self-enrichment models, there is no consensus on the dominant physical processes or expected level of enrichment as a function of cluster mass. The lack of a tilt in the red GCs might be expected if the self-enriched metals were a constant fraction of the stellar mass, since the overall metal content of red GCs is a factor of $\\sim 10$ higher than the blue GCs (Strader \\etal~2006; Mieske \\etal~2006). We illustrate this point further in \\S 3.  A possible relationship between GC self-enrichment and a mass-metallicity relation among GCs was discussed for the oldest halo Galactic GCs by Parmentier \\& Gilmore (2001), prior to the discovery of the blue tilt in external galaxies. However, Parmentier \\& Gilmore were attempting to explain an anticorrelation between GC metallicity and mass, the opposite of the blue tilt.  Here we do not attempt detailed modeling of the self-enrichment process (see the studies cited in the previous paragraphs for such models). Our model differs from many of those previously cited in that we invoke a single generation of star formation; within the context of self-enrichment, we expect a spread in metallicity in individual star clusters. This is inconsistent with the chemical homogeneity observed in most Galactic GCs, but the Galaxy does not have a blue tilt (see additional discussion in \\S 8).  Our approach is to use scaling relation assumptions about the putative enrichment to explore the conditions under which the observed GC mass-metallicity relationships may be reproduced. These conditions may then be used as additional constraints on GC formation. ",
        "conclusions": "We have examined a number of scenarios for the blue tilt, a mass-metallicity relationship for metal-poor GCs observed in many external galaxies. The model that seems to best reproduce current observations is self-enrichment in proto-GC clouds; in this model, star formation is governed by supernova feedback, and the efficiency is a function of the protocloud mass. Certain efficiency scalings produces no mass-radius relation for the resultant GC (in accordance with observations), and are equivalent to an energy balance condition for star formation. We have also speculated as to how the observed variations in the slope of the blue tilt and the break mass might be reproduced. Contrary to the comments in Strader \\etal~(2006), a dark matter halo is not necessarily required for effective self-enrichment, at least given the assumptions of the current model.  Many of the color-metallicity relations used are uncertain, and may be the source of some of the observed scatter in blue tilt slopes. Since it is unlikely that these relations will see significant improvement soon, it would be useful to make photometric observations of different galaxies with a common color (e.g., $B-I$ or $g-z$) so that the intercomparisons are at least consistent. Another route is to directly estimate the metallicities of individual clusters using spectroscopy, or with near-IR imaging, which is more sensitive to metallicity than optical colors. Most of the break masses are poorly determined, so deeper imaging of selected galaxies would be beneficial.  The assumption of a variable formation efficiency can be tested directly by measuring cloud and cluster masses in nearby systems of young massive clusters, as suggested by Ashman \\& Zepf (2001). With ALMA, accurate cloud masses and radii can be determined in such systems, testing both the efficiency of cluster formation and the initial mass-radius relationship for clouds as a function of external parameters such as pressure.  Self-enrichment has been separately proposed to explain a number of abundance anomalies in Galactic GCs, primarily involving light elements such as C, N, O, and Mg. The culprits are generally argued to be an unknown combination of massive main-sequence stars and intermediate-mass AGB stars. A natural expectation of self-enrichment models for the blue tilt would be the appearance of abundance anomalies associated with blue tilt GCs. The integrated abundances of C, N, and Mg can all be estimated in extragalactic GCs from low-resolution spectra. For example, Cenarro \\etal~(2007) find CN line strengths for blue tilt GCs in NGC 1407 that are much larger than in Galactic GCs at comparable metallicities; unfortunately, less luminous GCs in NGC 1407 have not been observed. Studying the abundance patterns of GCs both above and below the blue tilt break mass could constrain the timing and duration of the hypothesized self-enrichment events.  The self-enrichment model makes another prediction that was described in \\S 1 and \\S 3. Unless star formation in the GC is oddly segregated by mass, such that the high mass stars form, explode as supernovae, and enrich the intracluster gas before any low-mass stars form, then the blue tilt GCs would be expected to have internal dispersions in not just the lighter elements such as C through Mg, but also in heavier elements produced by massive star supernovae. In the Milky Way the only GC known to have a large star-to-star spread in the $\\alpha$-elements heavier than Mg is $\\omega$ Cen (e.g., Freeman \\& Rodgers 1975; Smith et al. 2000), a cluster that has been viewed as a test case of self-enrichment and chemical evolution (e.g., Carraro \\& Lia 2000; Ikuta \\& Arimoto 2000; Platais et al. 2003; Marcolini et al. 2007), and possibly M22 (Norris \\& Freeman 1983; Lehnert et al. 1991; Anthony-Twarog et al. 1995). The majority of Galactic GCs are homogeneous to a high degree in the $\\alpha$-elements.  However, the Milky Way GC system is thought not to exhibit a blue tilt (although see Parmentier \\& Gilmore 2001), so there is no conflict with the self-enrichment scenario for this phenomenon. However, the challenge in other galaxies could be two-fold: (i) if blue tilt GCs are chemically homogeneous, then the self-enrichment picture may be less compelling; (ii) if blue tilt GCs are typically inhomogeneous in heavy elements, why is this not true in Galactic GCs of comparable mass? Regarding (i), it is intriguing that four of the most massive GCs in M31 have large internal metallicity spreads (Fuentes-Carrera \\etal~2008). Forthcoming homogeneous photometric and spectroscopic datasets should allow the detection of a blue tilt in M31, if it exists. From a theoretical standpoint, self-enrichment models such as those of Brown \\etal~(1991) attempt to reconcile supernova-induced enrichment with chemical homogeneity. With regard to (ii), we suggest reluctantly that it might be necessary to invoke fundamental differences between the parent clouds of Galactic GCs and those in blue tilt galaxies."
    },
    "0808/0808.1071_arXiv.txt": {
        "abstract": "{The equilibrium rotation of tidally evolved ``Earth-like'' extra-solar planets is often assumed to be synchronous with their orbital mean motion. The same assumption persisted for Mercury and Venus until radar observations revealed their true spin rates. As many of these planets follow eccentric orbits and are believed to host dense atmospheres, we expect the equilibrium rotation to differ from the synchronous motion. Here we provide a general description of the allowed final equilibrium rotation states of these planets, and apply this to already discovered cases in which the mass is lower than 12 $M_{\\oplus}$. At low obliquity and moderate eccentricity, it is shown that there are at most four distinct equilibrium possibilities, one of which can be retrograde. Because most presently known ``Earth-like'' planets present eccentric orbits, their equilibrium rotation is unlikely to be synchronous.} ",
        "introduction": "After a significant number of discoveries of new extra-solar gaseous giant planets, a new barrier has been passed with the detections of several planets in the Neptune and even Earth-mass ($M_{\\oplus}$) regime: $5-12$ $M_{\\oplus}$ \\citep{Rivera_etal_2005,Lovis_etal_2006,Udry_etal_2007,Bonfils_etal_2007}. If the commonly accepted core-accretion model can account for the formation of these planets, resulting in a mainly icy/rocky composition, the fraction of the residual He-H$_2$ atmospheric envelope accreted during the planet migration is not tightly constrained for planets more massive than the Earth \\citep[e.g.][]{Alibert_etal_2006}. A minimum mass of below 10 $M_{\\oplus}$ is usually considered to be the boundary between terrestrial and giant planets, but \\citet{Rafikov_2006} found that planets more massive than 6 $M_{\\oplus}$ could have retained more than 1 $M_{\\oplus}$ of the He-H$_2$ gaseous envelope. For comparison, masses of Earth's and Venus' atmosphere are respectively $\\sim 10^{-6}$ and $10^{-4}$ times the planet's mass. Despite significant uncertainties, these discoveries of ``super-Earths'' provide an opportunity to test some properties that could be similar to those of our more familiar telluric planets.  Because some of these planets are potentially in the ``habitable zone'' \\citep{Udry_etal_2007,Selsis_etal_2007}, their spin state is an important factor in understanding their climate. For planets with a negligible atmosphere such as Mercury, solid tides induced by the host star are expected to despin the planet until an equilibrium position or a capture in a spin-orbit resonance, depending on the orbital eccentricity and the permanent quadrupole moment of inertia \\citep{Goldreich_Peale_1966,Correia_Laskar_2004}. % However, for planets with dense atmosphere such as Venus, thermal atmospheric tides, driven by solar insolation, have a profound influence on the spin, and may destabilize the previous tidal equilibrium state. \\citet{Correia_Laskar_2001,Correia_Laskar_2003I} investigated the combined effect of gravitational and atmospheric tides on Venus and showed that the planet could evolve into two different final states, the existing retrograde motion being the most probable one. This study was  performed with the small eccentricity approximation that may no longer be adequate for  ``Earth-like'' planets, which exhibit a wide range of eccentricities, orbital distances, or central star types. Following the work of \\citet{Laskar_Correia_2004}, we generalize here these previous  studies and investigate the possible equilibrium rotation states of ``Earth-like'' planets in the presence of significant eccentricity. Although our knowledge of these planets is restricted to their orbital parameters and minimum masses, we attempt to place new constraints on their surface rotation rate, assuming that they have a dense atmosphere. ",
        "conclusions": "We have derived a simple model that permits the determination of the final equilibrium rotation of ``Earth-like'' planets. Our model contains some uncertain parameters related to the dissipation within the planets, but we were able to gather all this information in a single parameter, $\\omega_s$. By varying this, we can then cover all possibilities for the rotation of terrestrial planets. We demonstrated that for a planet of moderate eccentricity and low obliquity, at most four final equilibrium positions are possible. For eccentricities higher than $ e \\sim 0.5 $, terms of higher degree in $e$ should be considered in Eq.(\\ref{eq02}) that may generate additional equilibrium positions.  Based on the present rotation of Venus, we provided an estimate of $\\omega_s$ for different environments (Eq.\\,\\ref{eq60}). An important consequence is that the ratio $ \\omega_s / n $ increases rapidly with the semi-major axis and mass of the star because $ \\omega_s /n \\propto (a M_*)^{2.5} $. The effect of the atmosphere on the equilibrium rotation is therefore more relevant for planets that orbit Sun-like stars at not close distances. Unfortunately, this situation is precisely the opposite of that for the discovered Earth-like planets (Table\\,\\ref{Tab1}), since the radial velocity technique is more sensitive to the detection of short-period planets. {\\bfi As a consequence, unlike Venus, none of these planets can be stabilized with a rotation rate $\\omega < n $, because for $ e > 0.1 $ these final states only exist if $ a M_* > 0.2 $\\,AU\\,$M_\\odot$ (Fig.\\ref{Fig2}).} The effect of atmospheric tides is extremely small ($ \\omega_s / n \\sim 0 $) and the equilibrium rotation is essentially driven by the tidal gravitational torque. However, their significant eccentricity ($0.1 < e < 0.2$) moves the equilibrium position away from the synchronous motion. Capture in this resonance is unlikely, but capture in a higher order spin-orbit resonance cannot be rejected, in particular for {\\small GJ}\\,581\\,d and {\\small GJ}\\,674\\,b, depending on the asymmetry of their mass distribution."
    },
    "0808/0808.0077_arXiv.txt": {
        "abstract": "Strong size and internal density evolution of early-type galaxies between $z\\sim 2$ and the present has been reported by several authors.  Here we analyze samples of nearby and distant ($z\\sim 1$) galaxies with dynamically measured masses in order to confirm the previous, model-dependent results and constrain the uncertainties that may play a role.  Velocity dispersion ($\\sigma$) measurements are taken from the literature for 50 morphologically selected $0.8<z<1.2$ field and cluster early-type galaxies with typical masses $\\mvir = 2\\times10^{11}~\\msol$.  Sizes ($\\reff$) are determined with Advanced Camera for Surveys imaging.  We compare the distant sample with a large sample of nearby ($0.04<z<0.08$) early-type galaxies extracted from the Sloan Digital Sky Survey for which we determine sizes, masses, and densities in a consistent manner, using simulations to quantify systematic differences between the size measurements of nearby and distant galaxies.  We find a highly significant difference between the $\\sigma$-$\\reff$ distributions of the nearby and distant samples, regardless of sample selection effects.  The implied evolution in $\\reff$ at fixed mass between $z=1$ and the present is a factor of $1.97 \\pm 0.15$. This is in qualitative agreement with semianalytic models; however, the observed evolution is much faster than the predicted evolution. Our results reinforce and are quantitatively consistent with previous, photometric studies that found size evolution of up to a factor of 5 since $z\\sim 2$.  A combination of structural evolution of individual galaxies through the accretion of companions and the continuous formation of early-type galaxies through increasingly gas-poor mergers is one plausible explanation of the observations. ",
        "introduction": "Hierarchical galaxy formation models embedded in a $\\Lambda$CDM cosmology predict strong size evolution for massive galaxies.  A higher gas fraction in high-redshift galaxies leads to more dissipation and hence compact galaxies \\citep[e.g.,][]{robertson06, khochfar06a}, and subsequent evolution such as dry merging or accretion of smaller systems can increase the size of a galaxy \\citep[e.g.][]{loeb03, naab07}.  The models predict the strongest sample-averaged size evolution for the most massive galaxies \\citep{khochfar06b} because of large differences in the gas fraction at different redshifts and because the assembly of massive galaxies continues until very late epochs in a hierarchical framework\\citep[e.g.,][]{delucia06}.  Evidence for significant size evolution between $z\\sim 2$ and the present has been building up quickly over the past few years \\citep[e.g.][]{trujillo04, trujillo06a, franx08}. In particular, galaxies with low star formation rates and high stellar masses ($\\gtrsim 10^{11}~\\msol$) appear to be extremely compact from $z\\sim 1.5$ \\citep{daddi05, trujillo06b, trujillo07, longhetti07, cimatti08, rettura08} to $z\\sim 2.5$ \\citep{zirm07, toft07, vandokkum08b, buitrago08}.  Given the similarity between many of their observed properties there is likely to be an evolutionary connection between these distant compact galaxies and the present-day early-type galaxies despite the measured large difference of 2 orders of magnitude in surface mass density \\citep[e.g.,][]{zirm07}.  The measurements of sizes and densities of high-redshift galaxies are hampered by many systematic uncertainties (e.g., morphological \\textit{K}-corrections, surface brightness dimming, errors in photometric redshifts, and mass measurements). Most of these errors, however, are unlikely to fully account for the observed strong size evolution.  The uncertainty in the mass estimates may be the exception.  For the work in the literature, these mass estimates are always based on the photometric properties of the galaxies.  For a reasonable set of assumptions the photometric stellar mass estimates are not uncertain by more than a factor of 2 or 3 and would not change the inferred evolution significantly.  However, since we infer that $z\\sim 2$ galaxies must have physical central densities that are 3 orders of magnitude higher than those of local galaxies \\citep{vandokkum08b}, further verification of those apparently reasonable assumptions is warranted.  For example, a stellar initial mass function (IMF) that is radically different \\citep[e.g.,][]{larson05, fardal07, vandokkum08a, dave08} from a Salpeter-like IMF \\citep{salpeter55, scalo86, kroupa01, chabrier03, hoversten08} could reduce the stellar mass estimates by an order of magnitude, producing perfectly normal galaxies by today's standards.  The spectacular nature of these compact galaxies at $z\\sim 2$ could be confirmed by direct, kinematical mass measurements.  However, the quiescent nature of these objects and their consequent lack of emission lines \\citep{kriek06b} require absorption-line measurements of their stellar velocity dispersions, which should be as high as $400-500~\\rm{km~s}^{-1}$ \\citep{toft07, vandokkum08b}.  Unfortunately, with the currently available instrumentation this is not feasible. These $z \\sim 2$ galaxies are prohibitively faint at observed optical wavelengths \\citep[see, e.g.,][]{cimatti08}, and near-infrared spectroscopy is still maturing as a technique. Continuum detections in the observed NIR have only recently become possible \\citep{kriek06b} for the brightest sources, and no detection of absorption lines has been made.  At lower redshifts ($z\\sim 1$) absorption-line spectroscopy has for years been a powerful tool to study the evolution of distant early-type galaxies \\citep{vandokkum98, vandokkumstanford03, vandokkumellis03, wuyts04, vanderwel04, treu05a, holden05, vanderwel05, treu05b, diserego05, jorgensen05}. Size evolution is a gradual process \\citep[see, e.g.,][]{trujillo06a}; therefore, intermediate changes in sizes and densities should be observable at these redshifts.  In this paper we compile a sample of galaxies at redshifts $0.8<z<1.2$ with measured absorption line velocity dispersions from the literature and that are visually classified as early-type galaxies with the aid of \\textit{Hubble Space Telescope} ({\\it HST}) imaging from the Advanced Camera for Surveys \\citep[ACS;][]{ford98}. We measure the galaxies' sizes from these ACS data.  We then compare this distant sample with nearby early-type galaxies extracted from the Sloan Digital Sky Survey \\citep[SDSS;][]{york00}.  This comparison, with careful control of systematic uncertainties, allows us to verify that distant early-type galaxies are indeed significantly more compact than their local counterparts.  The advantage of this approach is that the size and density measurements are independent of the photometric properties of the galaxies apart from the surface brightness profile.  The absence of luminosity and other photometric properties from our analysis assures us that our study does not suffer from the strong possible biases in previous photometric work.  Moreover, deep, high-resolution ACS imaging allows us to determine sizes of $z\\sim 1$ galaxies to a precision comparable to that nearby galaxies and in a consistent manner.  Most previous studies verify for biases in the size determinations within their distant samples \\citep[e.g.,][]{trujillo04, trujillo06a, cimatti08} but do not extend this analysis to verify the consistency with size measurements of nearby galaxies.  In \\S~\\ref{nearby} we describe the sample of nearby early-type galaxies and derive the dynamical mass-size relation.  In \\S~\\ref{distant} we construct the sample of $z\\sim 1$ early-type galaxies, determining their masses and sizes in a manner that is consistent with the nearby sample.  In \\S~\\ref{sim} we quantify systematic effects in our size measurements through simulations.  In \\S~\\ref{results} we derive the evolution in the dynamical mass-size relation.  In \\S~\\ref{dis} we compare our results with previous measurements based on photometric mass estimates and semianalytic model predictions, and we discuss size evolution in the broader context of the evolving early-type galaxy population.  Finally, in \\S~\\ref{sum} we summarize our results and conclusions.  We adopt the following cosmological parameters: $(\\Omega_M,~\\Omega_{\\Lambda},~h) = (0.3,~0.7,~0.7)$. ",
        "conclusions": "\\label{dis}  \\subsection{Comparison with Photometric Size-Evolution Measurements}\\label{dis:1}  The main goal of this paper is to use dynamical measurements to investigate whether early-type galaxies were smaller and denser in the past.  Previous work has shown that the stellar mass surface density is higher, but there are a number of issues with such studies as they rely on stellar population models and they ignore possible changes in the underlying dark matter profile.  In Fig.~\\ref{lit_z} we compare the size-evolution results presented in the previous section with size-evolution results for early-type galaxies based on photometric mass estimates.  For all the literature samples we take the mean redshift and the mean stellar mass (normalized to the Kroupa IMF), and compute the mean offset from the local mass-size relation from \\citet{shen03}.  We include four intermediate redshift cluster galaxy samples with photometrically measured masses and sizes from WFPC2 or ACS imaging.  The data are described by \\citet{holden07} and the sizes are measured as described in this paper (\\S~\\ref{distant}).  These four clusters are \\clll at $z=0.33$, \\msl at $z=0.59$, and \\ms and \\cll, both at $z=0.83$.  Note that we also include \\ms in the present study with dynamical mass measurements.  The agreement between the independent measurements confirms that at least out to $z\\sim 1$ dynamical and photometric mass estimates based on optical colors and spectral energy distributions agree within the statistical errors as was previously shown by \\citet{vanderwel06b}, \\citet{rettura06}, and \\citet{holden06}.  The literature samples have all been selected in different ways, and so a direct comparison with our work may not be straightforward. Not all samples are morphologically selected; many are selected by their spectral or photometric properties. In the local universe there is substantial overlap between samples of early-type galaxies that are selected by different criteria; therefore, it is a reasonable assumption to suppose this to also be the case at high redshift, where different indicators (low star formation rates, red colors, smooth visual appearance) also reflect a common nature. Recently, several studies have shown hints that this is indeed the case \\citep[e.g.,][]{vandokkum08b, kriek08b}, but these issues need to be further addressed in the future.  Even with this cautionary proviso, the broad agreement between the results presented in this paper and the photometric results at higher redshifts is striking. All studies included in Fig.~\\ref{lit_z} are consistent with significant size evolution of several factors between $z\\sim 1-2$ and the present for galaxies with a given mass.  A linear fit in log-log space to our two data points at $z\\sim 0.06$ and $z\\sim 1$ gives $\\reff(z)\\propto (1+z)^{-0.98\\pm0.11}$.  With a linear fit to the photometric data the inferred rate of evolution is $\\reff(z)\\propto (1+z)^{-1.20\\pm0.12}$, where the error is obtained via a bootstrap/Monte Carlo simulation.  The broad agreement of our measurement of the size evolution of early-type galaxies with the photometric studies is encouraging and alleviates concerns about serious systematic effects that potentially could have compromised previous work.  Most notably, uncertainties in the photometric mass estimates used in all other previous work appear to have a limited impact, at least compared to factors of $\\gtrsim$5 which would mimic the strong, observed size and density evolution. Uncertainties in photometric mass estimates on the level of a factor of $\\sim$2 due to differences among the various stellar population models \\citep[e.g.,][]{bruzual03, maraston05} remain an issue, but to invoke, for example, an unconventional stellar IMF as an alternative to radically different structural properties of high-redshift early-type galaxies is no longer necessary.  Other systematic uncertainties cannot explain the observed evolution either.  In our size measurements, systematic effects have been taken into account (see Secs.~\\ref{distant:profile} and \\ref{sim}).  We are confident, for example, that we would detect low-surface brightness envelopes around distant galaxies.  Furthermore, we know that only a minority of morphologically selected $z\\sim 1$ early-type galaxies ($\\sim$10\\%) show signs of nuclear activity \\citep[e.g.,][]{rodighiero07, vanderwel07a} such that it is unlikely that central point sources affect our size measurements.  This is also clear from the fact that the residuals of our $R^{1/4}$ profile fits generally do not show central point sources and that none of the deep spectra used to measure dispersions show evidence for nuclear activity.  Furthermore, the good correspondence between the rest-frame wavelength of the imaging data sets used at different redshifts assures us that morphological $K$-corrections do not play a significant role.  Despite the broad consistency between our results and those previously published, the agreement is not perfect.  There is a marginal inconsistency at the $1.5~\\sigma$ level between the size-evolution measurement from kinematic data and the size-evolution measurement from photometric data shown in Fig.~\\ref{lit_z}.  This could point to the presence of some systematic effects within the $z>1.5$ results. Alternatively, the different studies sample galaxies with a wide range in masses, and therefore mass-dependent size evolution could lead to apparent discrepancies among the samples.  This is explored in the following section.  Our robust results strengthen the results from previous studies.  We conclude that early-type galaxies at $z=1$ are $\\sim 2$ times smaller than local early types with the same mass, and that at $z=2-2.5$ this size difference is likely increased to a factor of $\\sim$4, as previously observed by \\citet{zirm07}, \\citet{toft07}, \\citet{vandokkum08b}, and \\citet{buitrago08}.   \\subsection{Comparison with Model Predictions}\\label{dis:2}  The fact that we see considerable evolution in galaxy size with redshift is not surprising from a theoretical perspective.  Most semianalytic models of galaxy formation in a $\\Lambda$CDM universe predict substantial size evolution over the past several billion years.  A comparison between the observed and model-predicted amount of size evolution will help to identify the mechanism(s) that are responsible.  In Fig.~\\ref{M_dens} we directly compare the observed evolution in surface density with the predictions from the semianalytic work by \\citet{khochfar06b}.  For galaxies with a given mass the model significantly under-predicts the evolution in size and density, except, perhaps, for the most massive galaxies.  In our data set we see no indication that the magnitude of size and density evolution increases with galaxy mass, as predicted by the models.  In fact, the most massive galaxies in our sample are precisely the only ones that are not different from local massive galaxies.  Note, however, that statistically speaking the evidence for mass-dependent evolution is weak (see \\S~\\ref{results:mrevol}).  Moreover, the most massive galaxies in our distant sample are a special subset, brightest cluster galaxies.  Such galaxies have been shown to have properties that deviate from those of other massive galaxies \\citep[see, e.g.,][]{vonderlinden07, bernardi07b}.  By including the $z=1.5-2.5$ photometric samples discussed in \\S~\\ref{dis:1} we can place further constraints on the models.  In Fig.~\\ref{lit_M} we compare the observed size evolution of the available samples, normalized to $z=1$, with the model predictions from \\citet{khochfar06b}.  Representing the model predictions by a single line is justified by the fact that the predicted evolution of $\\log(\\reff)$ with $\\log(1+z)$ is very close to linear.  Again, the observed size evolution is stronger than that predicted by the model.  It is interesting to note that, in qualitative agreement with the model prediction, we see a hint that size evolution depends on mass in the compilation presented in Fig.~\\ref{lit_M}.  The samples containing on average the lowest-mass galaxies display marginally less evolution. It has to be kept in mind, however, that small sample sizes and systematic effects are more important for determining second-order effects such as mass dependence (see \\S~\\ref{results:mrevol}).  A clue that systematic uncertainties may play a role is the remaining difference between the kinematic and photometric samples. Alternatively, it may signify non-linear evolution of $\\log(\\reff)$ with $\\log(1+z)$.  \\begin{figure}[t] \\epsscale{1.2} \\plotone{lit_z.eps} \\caption{Size evolution with redshift as derived in this paper with dynamically determined masses (\\textit{large filled circles}) compared with previous results based on photometric masses (\\textit{small filled circles}).  The solid line connects our samples at $z\\sim 0.06$ and $z\\sim 1$, the dashed line is a linear least-squares fit to the small filled data points.  The open circles are samples of cluster galaxies with photometrically measured masses and serve as an illustration that size evolution shows a continuous trend between $z=2.5$ and the present.  The broad agreement in size evolution as derived from galaxies with dynamically and photometrically determines masses reinforces the conclusions of previous, photometric studies whose results were potentially mitigated by considerable systematic effects that do not affect our analysis. } \\label{lit_z} \\end{figure}  \\begin{figure}[t] \\epsscale{1.2} \\plotone{lit_M.eps} \\caption{Size evolution per unit redshift vs. mean galaxy mass of our sample (\\textit{large circle}) and samples taken from the literature (small squares; see Fig.~\\ref{lit_z} for references).  Samples consisting of high-mass galaxies show somewhat stronger size evolution than samples consisting of low-mass galaxies, which is qualitatively, but not quantitatively, consistent with the predictions from the semianalytic model from \\citet{khochfar06b} (\\textit{solid line}).  This conclusion should be considered highly tentative, however, as this interpretation is hampered by systematic uncertainties and small sample sizes. } \\label{lit_M} \\end{figure}  \\subsection{Size Evolution of Individual Galaxies}\\label{dis:3}  It appears that the observed size evolution of a factor of $\\sim$2 between $z=1$ and the present for early-type galaxies with masses $\\sim10^{11}~\\msol$ is similar to the predicted evolution for early-type galaxies that are an order of magnitude more massive (see Figs.~\\ref{M_dens} and \\ref{lit_M}).  This suggests that the mechanism responsible for increasing the average size of early-type galaxies with time may be well understood, but that it is not implemented correctly in the current semianalytic model from \\citet{khochfar06b}. The process of size evolution may occur at different times and under different circumstances than is now assumed.  This may be related to the late assembly of very massive galaxies in models of this kind \\citep[see also, e.g.,][]{delucia06}, a prediction that is challenged by various observations \\citep[e.g.,][]{cimatti06, scarlata07, cool08}.  It is beyond the scope of this paper to fully discuss these possible discrepancies.  Instead we will explore the question whether the proposed physical processes responsible for size evolution are consistent with the observed trends.  In the semi-analytic models it is assumed (and this is confirmed by numerical simulations) that mergers drive size evolution.  The gas content of merging galaxies largely determines the relative size of the merger remnant compared to its ancestors.  Because gas fractions were higher in the past, galaxies that form early will be smaller than galaxies that form late. In the framework of cosmological simulations this means that galaxies at high redshift will be smaller because they were formed through gas-rich mergers and that those merger remnants can grow over time through subsequent mergers with other galaxies that are progressively more devoid of gas.  The question is whether the observed size evolution is dominated by size evolution of individual galaxies or simply by the addition of larger galaxies over time.  At $z=1$ only about 30--50\\% of the present-day early-type galaxy evolution had formed \\citep{bell04b, brown07, scarlata07, faber07}.  If we assume that these galaxies will make up the 30--50\\% most dense early-type galaxies in the present-day universe, then the scatter in the local $\\mvir$-$\\reff$ relation implies that the sample-averaged size increases by a factor of 1.3--1.4 between $z=1$ and the present.  Such evolution is thus expected in the absence of size evolution of individual galaxies and this is less than the observed evolution of a factor of 2.  To explain the observed evolution by growth of the early-type galaxy population without changes in the sizes of individual galaxies, the number density of early-type galaxies is required to increase by an order of magnitude between $z=1$ and the present.  Such strong evolution is clearly ruled out by the above-mentioned determinations of the number density of red galaxies at $z\\sim 1$.  Similarly, at $z=2$ only $\\sim$10\\% of the galaxies with masses $\\gtrsim~10^{11}~\\msol$ had been assembled \\citep{kriek08b}; if those galaxies, evolve into the 10\\% most dense present-day early-type galaxies then an increase in average size by a factor of $\\sim$2 can be accounted for, less than the observed amount of evolution.  These arguments are in agreement with the conclusions from \\citet{cimatti08}, who show that local galaxies with the same sizes and masses as galaxies in the $z=1-2$ samples are so rare in the local universe that it can be confidently ruled out that their structure remains unchanged up until the present day.  Note, however, that these arguments may be affected by the aforementioned biases in the SDSS (\\S~\\ref{nearby:sigrad}).   We conclude that size evolution due to the addition of larger galaxies over time contributes at most half of the observed evolution in the $\\mvir$-$\\reff$ relation.  The remainder must be due to size evolution of individual galaxies.  Numerical simulations have demonstrated that when early-type galaxies accrete neighbors without significant dissipational processes $\\sige$ does not change by much and that, to first order, $\\reff$ increases linearly with mass.  This does not depend strongly on the mass of the accreted object, i.e., the mass ratio of the merger \\citep{boylan05, robertson06, boylan06}.  Simulations in a cosmological context show that an increase in size by a factor of 2 between $z\\sim 1$ and the present is certainly possible \\citep{naab07}.  The strong observed size evolution thus argues in favor of a scenario in which significant mass from low-mass companions is accreted onto existing early-type galaxies over the past $\\sim$7 Gyr, which also explains the broad tidal features that are frequently observed around early-type galaxies \\citep{vandokkum05}.  As shown by \\citet{feldmann08} such features are not necessarily, and are even quite unlikely to be, the result of major merger events and are most likely due to the accretion of low-mass, gas-poor satellites.  We note that the size evolution of individual galaxies and the evolution of the sample average are inseparable because galaxies evolve in mass as well as in size.  Nonetheless, it is important to distinguish this complex scenario from the simple picture in which early-type galaxies that form at different redshifts have different sizes but do not structurally evolve at later times.  The strong observed size evolution clearly rules out the latter, indicating that the build-up of the early-type galaxy population is a complex and ongoing process.  Finally, it is remarkable that the change in the sizes of early-type galaxies is consistent with and differs by less than 15\\% from the change in the scale factor of the universe, $1+z$.  Within the standard cold dark matter scenario this is likely a coincidence since dissipational, strongly non-linear processes that are decoupled from cosmic expansion dominate at the kiloparsec scale of forming galaxies. Nonetheless, we cannot exclude the possibility that there is an underlying, fundamental reason that galaxies are scale-invariant with respect to a co-moving coordinate system.  In an alternative description of dark matter, i.e., Bose-Einstein condensed, ultra-light particles with a $\\sim$10 kpc-sized wave function \\citep[fuzzy dark matter or FDM,][]{sin94,hu00}, sizes of halos and their occupying galaxies possibly follow the cosmic expansion rate \\citep{lee08}."
    },
    "0808/0808.4132_arXiv.txt": {
        "abstract": "We propose an explanation to the puzzling appearance of a wide blue absorption wing in the He~I~$\\lambda10830 \\rm{\\AA}$ P Cygni profile of the massive binary star $\\eta$ Car several months before periastron passage. Our basic assumption is that the colliding winds region is responsible for the blue wing absorption. By fitting observations, we find that the maximum outflow velocity of this absorbing material is $\\sim 2300 \\km \\s^{-1}$. We also assume that the secondary star is toward the observer at periastron passage. With a toy-model we achieve two significant results. (1) We show that the semimajor axis orientation we use can account for the appearance and evolution of the wide blue wing under our basic assumption. (2) We predict that the Doppler shift (the edge of the absorption profile) will reach a maximum $0-3~{\\rm weeks}$ before periastron passage, and not necessarily exactly at periastron passage or after periastron passage. ",
        "introduction": "\\label{sec:intro}  There is no boring feature when it comes to the massive binary system \\astrobj{$\\eta$ Car}. Its $P=5.54 \\yr$ ($P =2022.7 \\pm 1.3~$d; Damineli et al. 2008a) periodicity has been observed in the radio (Duncan \\& White 2003), IR (Whitelock et al. 2004), visible (e.g., van Genderen et al. 2006), X-ray (Corcoran 2005), and in many emission and absorption lines (e.g., Damineli et al. 2008a,b). Every orbital cycle the system experiences a spectroscopic event, defined by the fading, or even disappearance, of high-ionization emission lines (e.g., Damineli 1996; Damineli et al. 1998, 2000, 2008a,b; Zanella et al. 1984). The rapid changes in the continuum, lines, and in the X-ray properties (e.g., Martin et al. 2006,a,b; Davidson et al. 2005; Nielsen et al. 2007; van Genderen et al. 2006; Damineli et al. 2008b; Corcoran 2005) are assumed to occur near the periastron passages of the highly eccentric, $e \\simeq 0.9$, binary orbit (e.g., Hillier et al. 2006).  The He~I~$\\lambda10830 \\rm{\\AA}$ high excitation line has a complex P Cygni profile, composed of three blue-shifted peaks with significant variations over the cycle (Damineli et al. 1998; 2008b). The emission profile has significant variations over the cycle. The Doppler shifts of the peaks are of relatively low velocities, $\\vert v_{\\rm peaks} \\vert < 300 \\km \\s^{-1}$ (Damineli et al. 2008b). The location of the minimum of the profile (the deepest point in absorption) does not change with orbital phase, and stays at $v_{obs-m}=-570 \\km \\s^{-1}$. However, the absorption profile does change. In particular, just before periastron a wide blue wing appear in absorption, reaching $\\sim -1000 \\km \\s^{-1}$ a month before periastron, and $\\sim -1800 \\km \\s^{-1}$ at periastron, and the maximum equivalent width of absorption occurs $10 \\days$ after periastron passage (Damineli et al. 2008b). Damineli et al. (2008b) approximated the average radial velocity of the absorption profile at half intensity, and found it to change from $-640 \\km \\s^{-1}$ before phase zero to $-450 \\km \\s^{-1}$ shortly after phase zero.  An additional He~I~$\\lambda10830 \\rm{\\AA}$ blue absorption feature of up to $-1000 \\km \\s^{-1}$, is observed at several arcseconds from the center in the lobes (Smith 2002). In this paper we refer only to the He~I~$\\lambda10830 \\rm{\\AA}$ ground measurements of Damineli et al. (2008b). We note that other absorption and emission lines of He~I can be formed in different regions in the binary system (see also Kashi \\& Soker 2007b). In particular, some visible He~I lines can be formed in the hot winds of the two stars close to their origin, as compared with the He~I~$\\lambda10830 \\rm{\\AA}$ that is formed in cooler regions. We show that this cooler region can be the post-shock primary wind. Therefore, different He~I lines need not have the same behavior along the orbit.  Because of the winds' very complicated flow structure, when starting this project we limited ourself to build a toy-model in order to achieve two goals: (1) To show that the orientation where the secondary is toward us at periastron ($\\omega = 90^\\circ$), can be accounted for the development of a wide blue absorption wing starting several weeks before periastron passage. (2) To encourage a nightly observation of the He~I~$\\lambda10830 \\rm{\\AA}$ close to periastron passage. ",
        "conclusions": "\\label{sec:es}  We study the blue absorption wing of the He~I~$\\lambda10830 \\rm{\\AA}$ P Cygni profile of $\\eta$ Car. The two winds of the two stars collide, and the post shocked gas of the two winds flow on both sides of a surface (the discontinuity surface) that has a pseudo-conical structure: the conical shell. The shocked secondary's wind accelerates part of the shocked primary's wind to high velocity (Figure \\ref{fig:colliding}). This gas, and the segments of the shocked secondary's wind which cool to a low temperature near the contact discontinuity (Figure \\ref{fig:colliding}), are assumed to be responsible for the blue absorption wing. This is the fundamental assumption of our model. We use our previous results and assume an orientation where the secondary is toward the observer at periastron ($\\omega=90^\\circ$; see Figure \\ref{fig:omega}). For the conical shell we built a toy model (Figure \\ref{fig:cone}).  The absorption profile, up to a scaling factor, is calculated according to equations (\\ref{eq:sum}) and (\\ref{eq:I}), and the results are presented in Figure \\ref{fig:abs}. We are interested in the bluest part of the absorption profile, where intensity is lower; in our scaled units these are the contours in the range $I \\simeq 0.8-1$. Using our fundamental assumption, and a toy model for the conical shell of the colliding winds, we showed that with our orbital orientation we can account for the appearance of the wide blue wing several months before periastron. Other semimajor orientations $\\omega$, cannot reproduce the results under our fundamental assumption (figure \\ref{fig:otheromega}) This is our main result, namely, that if the absorber responsible for the blue wing of the He~I~$\\lambda10830 \\rm{\\AA}$ line reside in the winds collision region, then only the $\\omega \\simeq 90^\\circ$ can account for the blue wing of this line.  Our results also predict that the Doppler shift $v_D$ (the edge of the profile) will reach a maximum $0-3~\\rm{weeks}$ before periastron passage. Since close to periastron the conical shell starts to collapse onto the primary, and even before it experiences the effect of wrapping, it is hard to pinpoint the exact time of this maximum. Nevertheless this maximum should be observed before the event."
    },
    "0808/0808.1521_arXiv.txt": {
        "abstract": "This is the first of a series of papers devoted to the investigation of a large sample of brightest cluster galaxies (BCGs), their kinematic and stellar population properties, and the relationships between those and the properties of the cluster. We have obtained high signal-to-noise ratio, long-slit spectra of these galaxies with Gemini and WHT with the primary purpose of investigating their stellar population properties. This paper describes the selection methods and criteria used to compile a new sample of galaxies, concentrating on BCGs previously classified as containing a halo (cD galaxies), together with the observations and data reduction. Here, we present the full sample of galaxies, and the measurement and interpretation of the radial velocity and velocity dispersion profiles of 41 BCGs. We find clear rotation curves for a number of these giant galaxies. In particular, we find rapid rotation ($>$ 100 km s$^{-1}$) for two BCGs, NGC6034 and NGC7768, indicating that it is unlikely that they formed through dissipationless mergers. Velocity substructure in the form of kinematically decoupled cores is detected in 12 galaxies, and we find five galaxies with velocity dispersion increasing with radius. The amount of rotation, the velocity substructure and the position of BCGs on the anisotropy-luminosity diagram are very similar to those of ``ordinary'' giant ellipticals in high density environments. ",
        "introduction": "The galaxies in the centres of clusters are unique. They are usually the dominant, brightest and most massive, galaxies in their clusters. A wealth of imaging data has been accumulated for these intriguing objects (Malumuth $\\&$ Kirshner 1985; Schombert 1986, 1987, 1988; Postman $\\&$ Lauer 1995; Collins $\\&$ Mann 1998; Brough et al.\\ 2002; Laine et al.\\ 2003), however the spectroscopic data is limited to very small samples or narrow wavelength coverage.  Tonry (1984), Tonry (1985), Gorgas, Efstathiou $\\&$ Aragon-Salamanca (1990), Fisher, Illingworth $\\&$ Franx (1995a), Fisher, Franx $\\&$ Illingworth (1995b), Cardiel, Gorgas $\\&$ Aragon-Salamanca (1998) and Carter, Bridges $\\&$ Hau (1999) each investigated 18 or less galaxies, and moreover measured only three or less indices. Brough et al.\\ (2007) used a wide wavelength range but included only three brightest cluster galaxies (BCG). Von der Linden et al.\\ (2007) studied BCGs from the Sloan Digital Sky Survey (SDSS), however the spatial information is lacking since fibers were used.  Among BCGs there exists a special class of galaxies catalogued as cD galaxies, although there exists a great deal of confusion over the exact meaning of the classifications gE, D and cD. Matthews, Morgan $\\&$ Schmidt (1964) outlined the following definition: ``D galaxies have an elliptical-like nucleus surrounded by an extensive envelope. The supergiant D galaxies observed near the centre of a number of Abell's rich clusters have diameters 3 -- 4 times as great as the ordinary lenticulars in the same clusters. These very large D galaxies observed in clusters are given the prefix ``c'', in a manner similar to the notation for supergiant stars in stellar spectroscopy''. Later, Schombert (1987) defined gE galaxies as distinct from other early-type galaxies by their large size; D galaxies as being gE galaxies with a shallow surface brightness profile slope; and cD galaxies as D galaxies with large extended stellar haloes. The last are also more diffuse than normal ellipticals (Schombert 1986). Because of the confusion surrounding the definition of D galaxies, and the fact that they are rarely regarded as a separate type of object in the modern literature, this study will only refer to cD and non-cD BCGs (i.e BCGs containing a halo or not). For BCGs, we adopt the definition to comply with recent literature (for example Von der Linden et al.\\ 2007), where BCG refers to the central, dominant galaxy in a cluster. For a small fraction of clusters, the BCG might not strictly be the brightest galaxy in the cluster.  Approximately 20 per cent of rich clusters of galaxies contain a dominant central cD galaxy (Dressler 1984; Oegerle $\\&$ Hill 2001), although they can be found in poor clusters as well (Giacintucci et al.\\ 2007). Some clusters have more than one cD galaxy, but a cD galaxy is always the dominant member of a local subcluster. The surface-brightness profiles of cD galaxies are displaced above the de Vaucouleurs law (de Vaucouleurs 1948) at large radii. The break in the cD galaxy surface-brightness profile typically occurs between 24 and 26 mag/arcsec$^{2}$ in the V-band (Sarazin 1988). The interpretation of this deviation is that the galaxy is embedded in an extensive luminous stellar halo. Three main theories have been proposed over the last four decades to explain the properties of BCG, and in particular cD, galaxies.  \\paragraph*{Theory 1.} In the first theory, BCG formation is caused by the presence of \\textit{cooling flows} in clusters of galaxies (Cowie $\\&$ Binney 1977). Cooling flow clusters are common in the local universe and BCGs are most often found at the centres of these systems (Edwards et al.\\ 2007). If the central cluster density is high enough, intracluster gas can condense and form stars at the bottom of the potential well. Observations that support this idea are blue- and UV-colour excesses observed in the central galaxy of Abell 1795 (indicative of star formation) by McNamara et al.\\ (1996) and molecular gas detected in ten out of 32 central cluster galaxies by Salom\\'{e} $\\&$ Combes (2003). Cardiel et al.\\ (1998) obtained radial gradients for the D$_{4000}$ and Mg$_{2}$ spectral features in 11 central cluster galaxies. Their observations were consistent with an evolutionary sequence in which radio-triggered star formation bursts take place several times during the lifetime of the cooling flow in the centre of the cluster. However, McNamara $\\&$ O'Connell (1992) find only small colour anomalies with small amplitudes, implying star formation rates that account for at most a few percent of the material that is cooling and accreting onto the central galaxy.  More recently, \\textit{XMM-Newton} observations showed that the X-ray gas in cluster centres does not cool significantly below a threshold temperature of $kT\\sim1-2$ keV (Jord\\'an et al.\\ 2004, and references therein). The central cluster galaxies often host radio-loud AGN which may account for the necessary heating to counteract radiative cooling (von der Linden et al.\\ 2007). Although BCGs are probably not completely formed in cooling flows, the flows play an important role in regulating the rate at which gas cools at the centres of groups and clusters.  \\paragraph*{Theory 2.} The second theory was proposed by Merritt (1983) and suggests that the essential properties of BCGs, and in particular those with haloes, are determined when the clusters collapse (\\textit{primordial origin}). Thereafter, frequent mergers of galaxies would be inhibited by the relatively high velocities between galaxies. Merritt (1983) argued that all galaxies had large haloes early in the life of the cluster. These haloes were then removed by the mean cluster tidal field during the initial collapse and returned to the cluster potential, except for the central member which remained unaffected because of its special position with respect to the cluster potential.  \\paragraph*{Theory 3.} The third, and most widely accepted, theory is in the context of the $\\Lambda$CDM cosmology and relates the formation of the central galaxy to mergings with or captures of less massive galaxies, and is known as \\textit{``galactic cannibalism\"}. It was first proposed by Ostriker $\\&$ Tremaine (1975) and developed by Ostriker $\\&$ Hausman (1977).  The most complete quantitative prediction of the formation of BCGs in the now standard CDM model of structure formation is by De Lucia $\\&$ Blaizot (2007). They used N-body and semi-analytic techniques to study the formation and evolution of BCGs and found that, 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 $z \\sim 5$ and 80 per cent at $z \\sim 3$) and in many small galaxies. They also found that BCGs assemble late: half of their final mass is typically locked up in a single galaxy after $z \\sim 0.5$ (illustrated in their figure 9). A very similar conclusion was reached by Romeo et al.\\ (2008), who performed N-body and hydrodynamical simulations of the formation and evolution of galaxy groups and clusters in a $\\Lambda$CDM cosmology to follow the build-up of two clusters and 12 groups. Observationally, Aragon-Salamanca et al. (1998) examined the K-band Hubble diagram for BCGs up to a redshift of $z$ = 1. They found that the BCGs had grown by a factor of two to four since $z$ = 1. Brough et al. (2002) found a similar result but discovered that the mass growth depended on the X-ray luminosity of the host cluster. They found that BCGs in high X-ray luminosity clusters showed no mass accretion since $z$ = 1 as opposed to BCGs in low X-ray luminosity clusters which grew by a factor of four. However, the recent near-infrared photometric study of 42 BCGs by Whiley et al.\\ (2008) in the 0.2 $<$ $z$ $<$ 1 range contradicts this. They studied the colour and rest-frame $K$-band luminosity evolution of BCGs and found it to be in good agreement with population synthesis models of stellar populations which formed at $z \\sim $ 2 and evolved passively thereafter.  Using the Millennium Simulation, De Lucia et al.\\ (2006) studied how formation histories, ages and metallicities of elliptical galaxies depend on environment and on stellar mass. Their Figure 9 shows the effective number of progenitors of early-type galaxies as a function of galaxy stellar mass. The number of effective progenitors is less than two for galaxies up to stellar masses of $\\simeq10^{11}$ M$_{\\odot}$. This function suddenly increases up to the value of approximately five effective progenitors for the mass of a typical BCG.  A related theory, called \\textit{tidal stripping}, was first proposed by Gallagher $\\&$ Ostriker (1972). Cluster galaxies that pass near the gravitational centre of the cluster may be stripped of some of their material by the tidal forces from the cluster potential or the central galaxy potential. The stripped material falls to the centre of the potential well, and could contribute to the observed haloes of cD galaxies. The most massive galaxies surrounding the central galaxy would be preferentially depleted as they are most strongly affected by dynamical friction (Jord\\'an et al.\\ 2004). The difference between stripping and primordial origin is that stripping (and cD halo formation) begins after cluster collapse whereas primordial origin assumes that the tidal events occur before collapse, and that the cD halo is not a consequence of tidal stripping (Schombert 1988).  The observation of multiple nuclei in central galaxies favours the cannibalism theory (Postman $\\&$ Lauer 1995). Tonry (1984) observed the velocity and velocity dispersion profiles of NGC6166 and NGC7720 and their multiple nuclei. He found that the stellar velocity dispersion of the central galaxy demonstrated that the multiple nuclei are not following circular orbits. This was followed by a bigger sample of 14 multiple nuclei BCGs for which the redshifts and stellar velocity dispersions are presented in Tonry (1985). Yamada et al.\\ (2002) showed that the BCG in a cluster at $z=1.26$ is composed of two distinct sub-units that are likely to fully merge on a time-scale of $10^{8}$ years. Jord\\'an et al.\\ (2004) studied the globular cluster systems in BCGs with confirmed haloes from Hubble Space Telescope observations. They concluded that the observed globular cluster metallicity distributions are consistent with those expected if BCGs galaxies form through cannibalism of numerous galaxies and protogalactic fragments that formed their stars and globular clusters before capture and disruption, although they state that the cannibalism scenario is not the only possible mechanism to explain these observations. The Jord\\'an et al.\\ (2004) globular cluster data also suggests that BCGs experienced their mergers prior to cluster virialisation, yet the presence of tidal streams suggest otherwise (Seigar, Graham $\\&$ Jerjen 2007).  Carter $\\&$ Metcalfe (1980) and West (1989) showed that the major axis of BCGs tends to be aligned with the major axis of the cluster galaxy distribution. Recent studies of the Coma cluster (Torlina, De Propris $\\&$ West 2007) show strong evidence that there are no other large-scale galaxy alignments other than for the BCGs. This also confirms that BCGs form via mechanism related to collimated infall of galaxies along the filaments and the growth of the cluster from the surrounding large scale structure (Boylan-Kolchin, Ma $\\&$ Quataert 2006).  In the CDM cosmology it is now understood that local massive cluster galaxies assemble late through the merging of smaller systems. In this picture, cooling flows are the main fuel for galaxy mass-growth at high redshift. This source is removed only at low redshifts in group or cluster environments, due to AGN feedback (De Lucia $\\&$ Blaizot 2007).  Some of the outstanding issues are: whether BCGs have a different formation mechanism than elliptical galaxies; and if the formation of BCGs are controlled by environment. To address these issues the present project was initiated, and long-slit spectra of a large and statistically significant sample of BCGs were obtained. These data provide a set of spectral indices, covering a wide wavelength range. The luminous central galaxies will be contrasted with other early-type galaxies to look for relative differences in their evolution, using features sensitive to stellar population age and abundances of various elements.  The first part of this project entails the determination of the stellar kinematics through derivation of velocity and velocity dispersion profiles. A forthcoming paper will be devoted to the measurement and analysis of line strengths for this sample of BCGs.  This paper is organised as follows. Section 2 details the sample and the selection criteria, followed by a description of the observations and data reductions in Section 3. The kinematical measurements are described in Section 4 and the individual galaxy kinematic profiles and notes are presented in Section 5. Section 6 discusses the kinematic properties of BCGs as a class, compared to other Hubble types. Conclusions and future work are given in Section 7. ",
        "conclusions": "\\label{kinematicanalysis}  (i) \\textit{Five out of 41 BCGs (ESO349-010, ESO444-046, ESO552-020, NGC3311, PGC026269) were found to have a positive velocity dispersion gradient.}  The radial kinematic studies done so far on early-type galaxies are mostly limited to normal ellipticals, for which flat or decreasing velocity dispersion profiles are found. The majority of the results previously obtained for very small samples of BCGs are similar to those of normal ellipticals. Fisher et al.\\ (1995a) found one galaxy (IC1101) in their sample of 13 BCGs with a positive velocity dispersion gradient. Carter et al.\\ (1999) found one (NGC6166) of their sample of three BCGs to have a positive velocity dispersion gradient, although Fisher et al.\\ (1995a) did not find it for this galaxy. Brough et al.\\ (2007) found negative velocity dispersion gradients in five out of their sample of six brightest cluster and group galaxies (the other one had a zero velocity dispersion gradient).  Both IC1101 and NGC6166 form part of the sample of BCGs studied here, but in both cases the measured velocity dispersion profiles do not reach the radius achieved in the much smaller samples in the above mentioned studies. Thus, the positive velocity dispersion gradient could not be confirmed for IC1101 or NGC6166 (see Figures \\ref{fig:NGC6166kin} and \\ref{fig:IC1101kin}). However, five other BCGs were found to have a positive velocity dispersion gradient, though admittedly the slope is marginally positive in some cases. If these positive velocity dispersion gradients are not caused by systematic errors in the various data reduction or velocity dispersion measurements by different authors, then they imply a rising mass-to-light ratio.  \\medskip (ii) \\textit{At least 12 BCGs (NGC3842, NGC4889, NGC7647, ESO488-027, IC5358, MCG-02-12-039, NGC1713, NGC2832, NGC4839, NGC6269, NGC7649 and UGC05515) show clear velocity substructure in their profiles.}  From studies of elliptical galaxies in high density environments (e.g. Koprolin $\\&$ Zeilinger 2000), the incidence of KDCs is observed to be about 33 per cent, rising to 50 per cent when projection effects are considered. Hau $\\&$ Forbes (2006) finds KDCs in 40 per cent of their isolated elliptical galaxies. For BCGs, we have found at least 12 of the 41 BCGs show clear velocity substructure, amounting to 29 per cent of the sample (intermediate and minor axis data included).  KDCs can be the result of a merger event (Koprolin $\\&$ Zeilinger 2000), but can also occur when the galaxy is triaxial and supports different orbital types in the core and main body (Statler 1991). The fact that the incidence of KDCs in BCGs compares with that found for normal elliptical galaxies in high density environments suggests that the two classes share the same fraction of galaxies with triaxial shapes.  \\medskip (iii) \\textit{NGC6034 and NGC7768 possess significant rotation (134 and 114 km s$^{-1}$ respectively) along the major axis. Several other BCGs show rotation that is $>$ 40 km s$^{-1}$ and more than three times the standard error: ESO346-003, GSC555700266 and NGC4839 (major axis spectra); ESO488-027, IC5358 and UGC05515 (intermediate axis spectra), as do the two elliptical galaxies NGC4946 (major axis spectra) and NGC6047 (intermediate axis spectra).}  Carter et al.\\ (1999) found small rotation along the major axis at large radii (30 -- 40 arcsec) for their sample of three BCGs, which is consistent with the nearly complete lack of rotation found near the centres of the sample of 13 BCGs by Fisher et al.\\ (1995a). According to Fisher at al.\\ (1995a), the lack of rotation found in samples of BCGs are in agreement with the expectation of declining importance of rotation with increasing luminosity for elliptical galaxies. The lack of rotation is also compatible with the idea that these objects formed through dissipationless mergers (Boylan-Kolchin et al.\\ 2006). The remnants left by mergers with or without dissipation are expected to differ in their kinematical structure. In a merger which involves gas-rich galaxies, the gas will form a disk. After the gas has been removed from the system at the end of the merger (through ejection and converted into stars), the remnant will show rotation (Bournaud et al. 2005). Whereas in a merger where dissipationless processes dominate, the remnant will show little or no rotation (Naab $\\&$ Burkert 2003; Cox et al. 2006). In this study, clear rotation above 100 km s$^{-1}$ was found for NGC6034 and NGC7768, while most BCGs showed little or no rotation. This kinematical differentiation (the existence of slow and fast rotators) in early-type galaxies is also clearly visible in the SAURON data presented by Emsellem et al.\\ (2006).  The amount of flattening that is expected due to rotation in a galaxy depends on the balance between ordered and random motions, and this can be quantified using the anisotropy parameter. The rotation of elliptical galaxies is conventionally expressed as the anisotropy parameter, defined as $(V_{\\rm max}/\\sigma_{0})^{\\ast} =$ ($V_{\\rm max}/\\sigma_{0}$)/$\\sqrt{\\epsilon/1-\\epsilon}$\\ \\ (Kormendy 1982), where the rotational velocity $V_{\\rm max}$ is measured as in Section \\ref{Sec:kinematics}, and the central velocity dispersion $\\sigma_{0}$ is taken as the measurement for the central velocity dispersion in Table \\ref{table:GalVel}. A value of $(V_{\\rm max}/\\sigma_{0})^{\\ast} \\approx 1$ would be expected if a galaxy is flattened by rotation. The anisotropy parameter $(V_{\\rm max}/\\sigma_{0})^{\\ast}$ can be used to separate galaxies that are rotationally supported from those that are supported by $\\sigma$ anisotropy where the value is substantially less than unity. The division occurs at about $(V_{\\rm max}/\\sigma_{0})^{\\ast} = 0.7$ (Bender, Burstein $\\&$ Faber 1992).  Figure \\ref{fig:Anisotropy} shows the anisotropy parameter as a function of galaxy $B$-band luminosity, where the above mentioned division is indicated by the horizontal dashed line. Note that the errors indicated are only propagated from the errors on the velocity measurements taken to be the extreme radial velocity points, and the errors on the central velocity dispersion. They do not take into account the general uncertainties involved in determining the most extreme velocity measurements. Therefore, care has to be taken when interpreting individual points on the diagram. The only notable BCG data point that lies significantly above the dashed line is for the galaxy PGC026269, which possesses moderate rotation (51 km s$^{-1}$) and surprisingly low central velocity dispersion (222 km s$^{-1}$). However, the ellipticity of this galaxy is zero and it is not rotationally supported.  Another factor that complicates the dynamical interpretation of individual points is that the observed ellipticity is a global property of the galaxy, whereas the kinematic measurements taken here only reflect the kinematics along the axis where the slit was placed, and only close to the centre of the galaxy. For example, a disk component may dominate the measured kinematics but will have little effect on the ellipticity, making the galaxy appear to rotate faster than its global ellipticity would suggest (Merrifield 2004).  For comparison, the sample of isolated ellipticals from Hau $\\&$ Forbes (2006) are also plotted in Figure \\ref{fig:Anisotropy}. Of this sample, 11 galaxies were observed along the major axis, and the central velocity dispersions were derived from the bins closest to the galaxy cores. The BCGs show less rotational support than the isolated elliptical galaxies, as a class. Spiral bulge (typically rotationally supported) and giant elliptical data from Bender et al.\\ (1992) are also plotted. Their central velocity dispersions were derived over the whole half-light radii of the galaxies. Even though large rotation velocities were found for a few individual cases, the BCGs are consistent with the general trend for very massive galaxies.  \\begin{figure} \\centering \\includegraphics[scale=0.32]{Anisotropy.eps} \\caption[The Anisotropy-luminosity Diagram]{The anisotropy-luminosity diagram. The horizontal dashed line separates the rotationally supported galaxies from the anisotropic galaxies as described in the text. Only the BCGs for which major axis spectra were taken (within 10 degrees) are plotted. The box plotted in the figure outlines the region containing data on giant ellipticals from Bender et al.\\ (1992).} \\label{fig:Anisotropy} \\end{figure}"
    },
    "0808/0808.0461_arXiv.txt": {
        "abstract": "The 9-month SWIFT Burst Alert Telescope (BAT) catalog provides the first unbiased ($N_H < 10^{24}$\\,cm$^{-2}$) look at local ($<z> = 0.03$) AGN.  In this paper, we present the collected X-ray properties (0.3 -- 12\\,keV) for the 153 AGN detected.  In addition, we examine the X-ray properties for a complete sample of non-beamed sources, above the Galactic plane (b$\\ge 15^{\\circ}$).  Of these, 45\\% are best fit by simple power law models while 55\\% require the more complex partial covering model.  One of our goals was to determine the fraction of ``hidden'' AGN, which we define as sources with scattering fractions $\\le 0.03$ and ratios of soft to hard X-ray flux $\\le 0.04$.  We found that ``hidden'' AGN constitute a high percentage of the sample (24\\%), proving that they are a very significant portion of local AGN.  Further, we find that the fraction of absorbed sources does increase at lower unabsorbed 2--10\\,keV luminosities, as well as accretion rates.  This suggests that the unified model requires modification to include luminosity dependence, as suggested by models such as the 'receding torus' model (Lawrence 1991).  Some of the most interesting results for the BAT AGN sample involve the host galaxy properties.  We found that 33\\% are hosted in peculiar/irregular galaxies and only 5/74 hosted in ellipticals.  Further, 54\\% are hosted in interacting/merger galaxies.  Finally, we present both the average X-ray spectrum (0.1--10\\,keV) and $\\log N$-$\\log S$ in the 2-10\\,keV band.  With our average spectrum, we have the remarkable result of reproducing the measured CXB X-ray power law slope of $\\Gamma \\approx 1.4$ (Marshall {\\it et al.} 1980).  From the $\\log N$-$\\log S$ relationship, we show that we are complete to $\\log S \\ge -11$ in the 2--10\\,keV band.  Below this value, we are missing as many as 3000 sources at $\\log S = -12$.  Both the collected X-ray properties of our uniform sample and the $\\log N$-$\\log S$ relationship will now provide valuable input to X-ray background models for $z \\approx 0$. ",
        "introduction": "Active galactic nuclei (AGN) are among the most powerful sources of energy in the Universe, where the brightest quasars can outshine all of the stars in their host galaxy by 100 times.  Optical studies of AGN reveal strong narrow and broad emission lines indicative of AGN activity in the nearby Universe (z$ < 1$).  However, X-ray and optical surveys fail to select the same AGN samples \\citep{2004ASSL..308...53M}.  X-ray surveys identify more sources whose 2 -- 10\\,keV emission is obscured by high column density absorbing material in the line of sight.  Still, even the X-ray surveys are affected by heavy obscuration, making it difficult to detect sources with column densities $> 10^{24}$\\,cm$^{-2}$ (Compton thick sources).  This provides a major question that AGN surveys need to address: how many heavily obscured or Compton thick sources exist?  \\citet{2000MNRAS.318..173M} estimated that the number could be as high as an order of magnitude more than the unobscured sources, which are easily detected in optical and soft X-ray surveys.  In order to account for these additional heavily obscured sources and determine their contribution to the cosmic X-ray background, very hard X-ray ($> 10$\\,keV) surveys are needed.  Only at these wavelengths is the AGN emission penetrating enough to pass through much of the dust and gas in the line of sight.  With the Burst Alert Telescope (BAT) on board SWIFT, we now have the first sensitive unbiased AGN survey, towards all but the most heavily obscured sources ($N_H > 10^{24}$\\,cm$^{-2}$).  This is due to BAT's sensitivity at very high X-ray energies (14 -- 195\\,keV).  With a much larger sample than previous (such as HEAO-1) and contemporary (Integral is most sensitive along the Galactic plane) very hard X-ray missions, BAT is the first sensitive all-sky very hard X-ray survey in 28 years.  From the 9-month catalog of BAT AGN \\citep{2007arXiv0711.4130T}, a large enough sample has been obtained (153 sources) to determine the statistical properties.  In the catalog paper, only the X-ray derived column density and a complexity flag were indicated.  However, in this paper we provide a more detailed look at the X-ray properties, including archival data, analyses from the literature, and previously unanalyzed SWIFT X-ray Telescope (XRT) observations.  This work follows upon an XMM-Newton follow-up study of 22 BAT AGN, in which we reported that the ``hidden''/buried AGN described in \\citet{2007ApJ...664L..79U} may be a significant fraction of the BAT AGN \\citep{2008ApJ...674..686W}.  Our goals are two-fold.  First, we present the X-ray properties for the entire 9-month catalog.  Second, we examine in more detail the collective properties of a uniform sample.  This sample consists of non-beamed sources with Galactic latitudes $\\ge 15^{\\circ}$. In Section~\\ref{bat-data}, we describe the observations and the spectral fits, including data from ASCA, XMM-Newton, Chandra, Suzaku, and SWIFT XRT.  In Section~\\ref{bat-spectra}, we describe the general properties of the entire BAT 9-month AGN sample.  In Section~\\ref{bat-unbiased}, we describe in more depth the properties of our uniform sample.  In Section~\\ref{bat-cxb}, we present the average X-ray spectrum as well as the 2--10\\,keV $\\log N$-$\\log S$ relation. These X-ray properties can now be used as input to X-ray background models for $z \\approx 0$. We then discuss the properties of the host galaxies in Section~\\ref{bat-host}.  Finally, we summarize our results in Section~\\ref{bat-summary}. ",
        "conclusions": "\\label{bat-summary} In this paper, we present the X-ray properties of a uniform sample of very hard X-ray (14 -- 195\\,keV) selected AGN.  We present a number of interesting results that highlight the many uses of a uniform very hard X-ray survey.  This paper is complimentary to the 9-month AGN survey paper \\citep{2007arXiv0711.4130T}, which presents the 14--195\\,keV properties of the sources.  Additionally, this paper confirms the results of our earlier study on {\\it XMM-Newton} observations of a representative sample of the BAT AGN  \\citep{2008ApJ...674..686W}.  Among these, we show that: (1) the X-ray and optical classifications agree, i.e. Sy 1s have low X-ray column densities while Sy 2s are more obscured, (2) the average power law index, $\\Gamma \\approx 1.8$, agrees with the results from HEAO-1 \\citep{1982ApJ...256...92M}, (3) ``hidden'' AGN are a significant fraction of local AGN, where we can now quantify this value as $\\approx 24$\\%, and (4) nearly half  (45\\%) of local AGN are well-fit by a simple model (all with $\\log N_H < 23$) while the remaining sources (55\\%) require a more complex model.  In addition, this paper presents a number of additional, important results.  From examining the host galaxy properties, we found that the majority of the X-ray obscuration is not simply from the host galaxy (by comparing host inclination ($b/a$) to X-ray column density).  The most surprising results, however, were that many of the host galaxies are peculiar/irregular galaxies (33\\%).  Further, an even larger fraction (54\\%) have either a close companion galaxy or are known mergers. This is observational proof that galaxy interactions may be driving activity in local supermassive black holes.  However, we also find that the distribution of AGN 2--10\\,keV luminosities and accretion rates, as well as morphologies, are the same between interacting and non-interacting hosts.  While, it is unclear what these results mean, however, there appears to be more than one trigger besides mergers for local AGN activity.  Our team is currently compiling higher quality images and photometry to better quantify these results.  From our uniform sample (102 sources with $|b| \\ge 15^{\\circ}$), we found that the distributions of both unabsorbed 2 -- 10\\,keV luminosity and accretion rate are significantly lower for Sy 2s than Sy 1s.  While earlier studies found this connection in 2 -- 10\\,keV luminosity \\citep{2003ApJ...596L..23S, 2003ApJ...598..886U}, this is the first time it has been reported in accretion rate.    We also showed that the fraction of obscured AGN is indeed larger for lower luminosities (absorption corrected $L_{2-10 keV}$) and accretion rates.  However, we note that the most heavily obscured sources ($\\log N_H \\ge 23$) do not dominate this relationship. Since the unified model predicts differences between absorbed and unabsorbed sources are a product of viewing angle alone, our results provide a challenge, arguing in favor of a luminosity-dependent AGN model.  Another result involves the correlation between accretion rate and $\\Gamma$.  In \\citet{2008ApJ...674..686W}, we had found indication of a connection between Eddington ratio (or 2 -- 10\\,keV luminosity) and $\\Gamma$ using the spectral fits of multiple observations for individual sources.  The fact that we did not observe a correlation in our larger sample seems to be a result of our sources having a larger range of Eddington ratios (or 2 -- 10\\,keV luminosities).  We suggest that previous studies, for instance by Shemmer et al. 2006, see this correlation because their samples have a narrower range of properties (being mid- to high luminosity AGNs).  The primary correlation appears to be with accretion rate and not hard band luminosity. Such a correlation should appear when comparing $\\Gamma$ to  $L^{corr}_{2-10 keV}/L_{Edd}$ for multiple observations of individual sources or for a sample of sources with a narrow range of accretion rates.  In a similar manner, we found that while our sample did not immediately confirm the X-ray Baldwin effect, binning the sources by luminosity, we were able to reproduce the anti-correlation between unabsorbed 2 -- 10\\,keV luminosity and Fe K EW.  The primary  anti-correlation, however, again appears to be with Eddington rate.  When we binned the values by our Eddington ratio proxy, we found that $EW \\propto {L^{corr}_{2-10 keV}/L_{Edd}}^{0.26 \\pm 0.03}$ (agreeing with the results of \\citet{2007AA...467L..19B}).  Since both $\\Gamma$ and Fe K EW are dependent on accretion rate, this suggests that the $\\Gamma$-EW correlation found by  \\citet{2007ApJ...664..101M} is a result of the accretion rate dependences.  Having classified the X-ray spectra of our sample into simple and complex categories, we were able to examine the properties of the two sub-samples in more detail.  For the simple model sources, we found that 41\\% of the sources exhibited a soft excess.  Having modeled this parameter with a simple blackbody model, we found the average temperature to be $kT = 0.10$\\,keV.  We also found that there was a significant amount of scatter in this value ($\\sigma = 0.07$\\,keV), contrasting with the \\citet{2004MNRAS.349L...7G} results for PG quasars.  We found no correlation between the blackbody temperature and Eddington ratio, black hole mass, or photon index.  However, we did find a correlation between the luminosity of the blackbody component and the luminosity in the power law.  This relationship is linear ($L_{pow} \\propto L_{kT}$) and may provide a challenge to the current soft excess models.  Examining the complex model sources, we found that the majority of these sources included absorbed AGN.  Of the 4 Sy 1s in this category, all have complex absorption features in their X-ray spectra.   For these sources, we showed that the nature of the soft emission ($L_{0.5-2 keV}$) for these sources is unclear.  Over half have soft band luminosities low enough to be the result of galactic emission from star formation/X-ray binaries.  However, of the sources with higher soft luminosities, 3 are ``hidden''/buried AGN.  This argues that the soft emission may be scattered AGN emission ($\\le 0.03$).  An important result we found is that the ``hidden''/buried AGN, sources with a high covering fraction, are a significant fraction of local AGN.  Among the complex sources, 45\\% are ``hidden''.  For these sources, we found that the FIR luminosity is not consistent with an increased star formation rate, as suggested by \\citet{2007ApJ...664L..79U}.  However, without higher quality X-ray spectra and multi-wavelength observations, we are unable to further explore the nature of these sources.  While BAT is quite good at finding heavily obscured sources, we found that none of the 9\\,month sources in our uniform sample have spectra consistent with heavily obscured Compton-thick objects ($N_H > 1.4 \\times 10^{24}$\\,cm$^{-2}$).   However, we do detect sources classified as Compton-thick in other studies based on a reflection dominated spectrum or strong Fe K EW (for instance 3C 452 and NGC 4945).  Additionally, we did not include an analysis of the very complex source NGC 6240, which may also be classified as Compton thick.  Since the Compton hump lies above 10\\,keV, spectral fits with and without reflection can be degenerate in the 0.1--10\\,keV band.  Therefore, a full analysis of the Compton thick nature of the BAT sources must be deferred to future studies.  One remarkable result we found came from the average spectrum we constructed in the 0.1--10\\,keV band with the measured spectral properties of our uniform sample.  Here, our data reproduce the measured slope of the CXB ($\\approx 1.4$).  This highlights the importance of the BAT survey in selecting heavily absorbed sources.  More importantly, this is observational proof that the combination of BAT-detected absorbed and unabsorbed local AGN replicate the shape of the CXB.  If the distribution of source properties at $z \\approx 1$, where much of the CXB originates, is the same as that of the BAT-detected AGN, the spectral paradox is resolved.  To test our completeness in the 2--10\\,keV band, we plotted the distribution of $\\log N$-$\\log S$ for the entire uniform sample.  This showed that while the sample is complete in the 14--195\\,keV band \\citep{2007arXiv0711.4130T}, we are only complete above $\\log S = -11$ in the 2--10\\,keV band.  Further, this distribution suggests that we are missing as many as 3000 sources at $\\log S = -12$, requiring that these sources have 14--195\\,keV fluxes below the current flux limit of the BAT survey.  Possibly these sources are ``hidden'' AGN with even higher X-ray columns ($\\log N_H \\ge 24$).  Such sources must have a high ratio of $F_{2-10 keV}/F_{14-195 keV}$, like NGC 1068.  Also, they may or may not contain Compton-thick sources, an answer to which our data can not supply.  These results, in addition to the X-ray properties (including column densities and spectral indices) will provide important input for CXB models at low redshift ($z \\approx 0$).  Overall, our analysis of the X-ray properties (and some host galaxy properties) show the interesting nature of very hard X-ray selected AGN.  In order to understand the properties further, we are continuing to collect and analyze the properties in the optical through spectra (Winter et al. in prep) and imaging (Koss et al. in prep), the IR through {\\it Spitzer} observations (Weaver et al. in prep), and radio (Sambruna et al. in prep).  Additionally, BAT is continuing to discover more sources at fainter fluxes, with sensitivity increasing as $\\sqrt t$.  With the additional sources in future BAT catalogs, we will obtain an even better understanding of local AGN."
    },
    "0808/0808.1927_arXiv.txt": {
        "abstract": "We discuss the dynamics of microquasar jets in the interstellar medium, with specific focus on the effects of the X-ray binaries' space velocity with respect to the local Galactic standard of rest. We argue that, during late stages in the evolution of large scale radio nebulae around microquasars, the ram pressure of the interstellar medium due to the microquasar's space velocity becomes important and that microquasars with high velocities form the Galactic equivalent of extragalactic head--tail sources, i.e., that they leave behind trails of stripped radio plasma.  Because of their higher space velocities, low--mass X-ray binaries are more likely to leave trails than high--mass X--ray binaries.  We show that the volume of radio plasma released by microquasars over the history of the Galaxy is comparable to the disk volume and argue that a fraction of a few percent of the radio plasma left behind by the X-ray binary is likely mixed with the neutral phases of the ISM before the plasma is removed from the disk by buoyancy.  Because the formation of microquasars is an unavoidable by-product of star formation, and because they can travel far from their birth places, their activity likely has important consequences for the evolution of magnetic fields in forming galaxies.  We show that radio emission from the plasma inside the trail should be detectable at low frequencies.  We suggest that LMXBs with high detected proper motions like XTE J1118+480 will be the best candidates for such a search. ",
        "introduction": "\\label{sec:introduction} The interaction of AGN jets with their environments has been under investigation for several decades, partly because it is easily observable through the morphology of radio lobes \\citep[e.g.][]{miley:80} and X-ray cavities \\citep[e.g.][]{mcnamara:07}.  Because the most powerful AGNs are essentially stationary in the centers of galaxy clusters, models of this interaction typically only consider stationary atmospheres that jets run into (see, e.g., \\citealt{heinz:06b} for recent work on jets in dynamic atmospheres).  The evolution in this case can be separated into three distinct phases: (1) the early momentum driven phase, where the ram pressure of the jet is significant for the dynamical evolution and the source evolves into an elongated structure with narrow cocoons, (2) the energy driven phase, where the slowed-down jet plasma inflates supersonically expanding lobes, excavating a quasi-spherical cavity (this phase is well described by the self-similar solution by \\citealt{castor:75,kaiser:97,heinz:98}), and (3) the late, sub-sonic evolution, when the radio lobes (the reservoirs of relativistic gas released by jets) are in pressure equilibrium with the environment. In the case of AGNs, this radio plasma is buoyant in the host galaxy/galaxy cluster atmosphere \\citep[see][for a more detailed review]{reynolds:02}.  It is becoming increasingly clear that X-ray binaries (XRBs) are also producing relativistic jets over a wide range in accretion rate \\citep{fender:00,fender:01,gallo:03,fender:04}.  The inner regions of these XRB jets, where their dynamics is governed by the atmosphere of the compact object powering them, are very similar to the inner regions of AGN jets \\citep{heinz:03a}.  However, \\cite{heinz:02c} showed that some critical differences exist between XRB jets and AGN jets in how they interact with the larger scale environment, well outside of the sphere of influence of the central compact object.  One of the differences is that, in comparison with AGN jets, the ISM poses a much weaker barrier to microquasar jets. This is because the XRB jet thrust (in other words, the ram pressure delivered by the jet) is much larger in comparison to the inertial density of the ISM than it is in the case of AGN jets.  The second fundamental difference between the two cases is, of course, that XRBs do not reside in the centers of dark matter halos with stratified gaseous atmospheres.  Instead, they travel through regular Galactic ISM with some space velocity $v_{\\rm XRB}$, set by supernova kicks and orbital dynamics.  The velocity dispersion of XRBs implies that a significant fraction of these sources are moving with large (supersonic) velocities through the ISM, which will have important consequences for their dynamics.  VLBI parallax measurements have shown that a microquasars can be moving with velocities in excess of $100\\,{\\rm km\\,s^{-1}}$ with respect to the local standard of rest \\citep{mirabel:01}.  As a result, the ram pressure of the ISM will act on the radio plasma released by the source, sweeping back the outer layers of the radio lobe.  In this paper, we argue that this aspect has important consequences for the interaction of some XRB jets with their environment: Instead of inflating stationary cocoons, in many cases they will produce trails of radio emitting plasma (made up of relativistic particles and magnetic fields) as they travel through the Galaxy.  A very similar situation is encountered in pulsar bow shock nebulae \\citep{cordes:96,frail:96}, where a relativistic wind inflates a channel in the ISM through which the pulsar is moving, and, of course, in exragalactic head-tail sources, which are the direct equivalent of jets propagating into a moving medium.  The paper is organized as follows: In \\S\\ref{sec:dynamics}, we will present a simple parametric model of the interaction of the jets with the ISM.  In \\S\\ref{sec:discussion} we discuss the consequences for particle and magnetic field input into the ISM, \\S\\ref{sec:observations} discussed the observational signatures of this interaction, and \\S\\ref{sec:summary} presents a brief summary of the paper. ",
        "conclusions": "\\label{sec:summary} We proposed that jets from X-ray binaries (microquasars) leave behind trails of non-thermal, synchrotron emitting plasma as they move through the interstellar medium.  LMXBs are more likely to leave such trails due to their longer life expectancy and, more importantly, due to the higher expected space velocities, while HMXBs likely produce more stationary radio lobes.  The total plasma volume deposited by XRBs over the history of the Galaxy is comparable to the total disk volume and constitutes a non-negligible fraction of the halo volume.  We argue that a fraction $f_{\\rm mix}$ of order a few percent of this plasma is mixed into the thermal ISM.  The magnetic field thus released can easily provide the seed field for Galactic dynamos to produce the fields observed in spiral galaxies even under the most conservative assumptions. Since HMXBs turn on rapidly after an episode of star formation, this mechanism should be important for galactic magnetic field production and maintenance within a time frame of about $10^{7}\\,{\\rm yrs}$ from the first supernovae.  We showed that radio emission from the plasma inside the trail should be visible primarily at low frequencies, since the radiative cooling time of the plasma at GHz frequencies limits emission to a roughly spherical region around the binary.  LOFAR, MWA, LWA and other future low frequency, wide field instruments are ideally suited for searches of this emission.  LMXBs with high detected proper motions like XTE J1118+480 will be the best candidates for such a search."
    },
    "0808/0808.1050_arXiv.txt": {
        "abstract": "Various laboratory-based experiments are underway attempting to detect dark matter directly. The event rates and detailed signals expected in these experiments depend on the dark matter phase space distribution on sub-milliparsec scales. These scales are many orders of magnitude smaller than those that can be resolved by conventional N-body simulations, so one cannot hope to use such tools to investigate the effect of mergers in the history of the Milky Way on the detailed phase-space structure probed by the current experiments.  In this paper we present an alternative approach to investigating the results of such mergers, by studying a simplified model for a merger of a sub-halo with a larger parent halo.  With an appropriate choice of parent halo potential, the evolution of material from the sub-halo can be expressed analytically in action-angle variables, so it is possible to obtain its entire orbit history very rapidly without numerical integration.  Furthermore by evolving backwards in time, we can obtain arbitrarily-high spatial resolution for the current velocity distribution at a fixed point. Although this model cannot provide a detailed quantitative comparison with the Milky Way, its properties are sufficiently generic that it offers qualitative insight into the expected structure arising from a merger at a resolution that cannot be approached with full numerical simulations. Preliminary results indicate that the velocity-space distribution of dark matter particles remains characterized by discrete and well-defined peaks over an extended period of time, both for single and multi-merging systems, in contrast to the simple smooth velocity distributions sometimes assumed in predicting laboratory experiment detection rates.  In principle, this structure contains a wealth of information about the formation history of the Milky Way's dark halo. ",
        "introduction": "Dark matter (hereafter DM) appears to be the dominant mass component of galaxies and large-scale structures in the Universe. The first evidence came in the 1930s (Zwicky 1933, Smith 1936), but it was only in the 1970s that observations of the rotation curves of galaxies demonstrated that DM dominates the masses of galaxies (Rubin \\& Ford 1970, Rubin, Thonnard \\& Ford Jr 1980). These observations showed that many rotation curves are approximately flat, or even rising, in the outer region of galaxies, where there is little luminous matter and so a Keplerian decline is expected.  Subsequently, work on the hierarchical structure formation paradigm showed that non-baryonic material known as ``cold dark matter'' (CDM) is required to match the observed large-scale structure of the Universe (Peebles 1982). The term ``cold'' derives from the fact that this material was non-relativistic at the epoch of matter-radiation equality. The density of this CDM has subsequently been indirectly measured by various experiments such as 2dFGRS (Percival et al., 2001), WMAP (Dunkley et al. 2008) and the Sloan Digital Sky Survey (Tegmark et al. 2006).  Particle physics provides us with various well-motivated candidates for the CDM, including weakly interacting massive particles (WIMPs). WIMPs can potentially be directly detected in the laboratory via their elastic scattering on target nuclei, and numerous experiments are currently underway to try to detect this phenomenon (e.g. Angle et al. 2007, Ahmed et al. 2008). The signals expected in these experiments, the number of recoil events per unit energy (and in some case its temporal and angular dependence), depend on the WIMP velocity distribution in the solar neighbourhood. A single stream of dark matter particles produces a step in the energy spectrum, detectable by a detector. The energy at which the step occurs is determined by the speed of the particles composing the stream (in the rest frame of the detector), while the height and position of the step vary annually, due to the Earth's orbit. Multiple streams would lead to a more complicated picture, and a superposition of enough such streams would ultimately be indistinguishable from a smooth distribution, depending on the detector resolution.  The question that we seek to address here is what form one might expect for this distribution in reality.  Predictions of the expected signals are often based on simplified models, which assume, for example, that the WIMP velocity distribution is Maxwellian (Freese, Frieman \\& Gould 1988) or a multivariate Gaussian (Evans, Carollo \\& de Zeeuw 2000, Helmi, White \\& Springel 2002). These models rely on the assumption that the Milky Way halo has reached a steady state so that the ultra-local DM phase-space distribution is smooth. However since structures form hierarchically, and the age of the Universe is not large compared with relevant dynamical timescales (such as the crossing time), this assumption is somewhat questionable.  Numerous N-body simulations have been performed studying the hierarchical formation and evolution of DM halos. Such simulations find that DM halos contain ubiquitous substructure, in particular in their outer regions (Moore et al.\\ 1999, Klypin et al.\\ 1999). However, the salient question here is whether or not the DM distribution is smooth on the scales probed by direct detection experiments.  Unfortunately, the resolution of even the best N-body simulations is many orders of magnitude larger than the relevant scales. The Sun's circular velocity around the centre of the Galaxy is $v_{\\odot} \\approx 200 \\; \\rm{km/sec}$, so that over a course of a year a terrestrial DM detector travels a distance \\begin{equation} \\rm{r_{det}} \\approx v_{\\odot} \\; \\tau_{\\rm{exp}} \\sim (200 \\; \\rm{km/sec}) \\; ({1 \\; \\rm{yr}}) \\sim 0.1 \\; \\rm{mpc} \\,, \\end{equation} whereas current N-body simulations cannot resolve scales smaller than ${\\sim \\; 100 \\, \\rm pc}$ (e.g. Diemand et al. 2007).  This represents an insurmountable problem for the conventional simulation techniques, indicating that a completely different, specialized approach is necessary to describe in detail the ultra-fine DM distribution probed by direct detection experiments. A first attempt at such a specialized simulation was carried out by Stiff and Widrow (2003; hereafter SW). Their method used a reverse simulation process to calculate the DM speed distribution, $f(v)$, at a single spatial point of the phase space, representing a detector. More specifically they ran a simulation of the formation of a DM halo and at the end of the simulation put down a uniform grid (in velocity space) of mass-less test particles at the point of interest. They then evolved the test and simulation particles back to the initial time, found where the phase-space sheet of the test particles intersects the initial DM phase-space distribution, and hence calculated the density of the test particles at the final detector position. In this way, they found that the DM distribution in the solar neighbourhood is characterized by a number of discrete peaks. This numerical approach allowed them to use initial conditions which reproduce a realistic hierarchical-formation model for the Milky Way. However, a drawback of this technique was that it proved numerically unstable, and, in order to stabilize the reverse integration, SW were obliged to introduce a softening length of some $20\\, {\\rm kpc}$ into the gravitational force law that they applied. This softening is worryingly large as it significantly exceeds the solar radius in the Milky Way, $R_0 \\approx 8.5\\, {\\rm kpc}$, so might be expected to affect the inferred DM phase space distribution impinging on a terrestrial detector.  More recently Vogelsberger et al. (2008) have formulated a technique for calculating the evolution of the fine-grained DM density in both static potentials and N-body simulations. They argue that the small-scale DM distribution that a terrestrial experiment would observe can be described by a multivariate Gaussian, in apparent contradiction to the SW results, perhaps indicating that the SW analysis had been compromised by the modification that they had to make to the gravitational force.  In this paper we develop a complementary approach to the analysis of the ultra-fine DM distribution. Following SW, we calculate the DM distribution in the solar neighbourhood via a backward evolution method, but using a simplified model for the potential that allows the system to be expressed in action-angle (hereafter AA) variables. Although this simplified potential provides a less realistic representation of the Milky Way, its qualitative properties are similar, and it has the great benefit of being analytically soluble. Thus, the gravitational force does not have to be artificially softened (allowing one to test whether this effect did compromise the SW results), and one can very rapidly explore parameter space without the computational overhead of numerical integration.  The paper is organized as follows. In Section~2 we present the simplified model that we use to describe the interaction between a galaxy like the Milky Way and a merging sub-halo. Section~3 contains our initial results, and we conclude in Section~4 with a discussion. ",
        "conclusions": ""
    },
    "0808/0808.1099_arXiv.txt": {
        "abstract": "We report the results of a search for pure rotational molecular hydrogen emission from the circumstellar environments of young stellar objects with disks using the Texas Echelon Cross Echelle Spectrograph (TEXES) on the NASA Infrared Telescope Facility and the Gemini North Observatory. We searched for mid-infrared H$_{2}$ emission in the \\sone, \\stwo, and \\sfour ~transitions. Keck/NIRSPEC observations of the H$_{2}$ \\snine ~transition were included for some sources as an additional constraint on the gas temperature. We detected H$_{2}$ emission from 6 of 29 sources observed: AB Aur, DoAr 21, Elias 29, GSS 30 IRS 1, GV Tau N, and HL Tau. Four of the six targets with detected emission are class I sources that show evidence for surrounding material in an envelope in addition to a circumstellar disk. In these cases, we show that accretion shock heating is a plausible excitation mechanism. The detected emission lines are narrow ($\\sim$10 km s$^{-1}$), centered at the stellar velocity, and spatially unresolved at scales of 0.4\\arcsec, which is consistent with origin from a disk at radii 10-50 AU from the star. In cases where we detect multiple emission lines, we derive temperatures $\\gtrsim$ 500 K from $\\sim$1 M$_{\\oplus}$ of gas. Our upper limits for the non-detections place upper limits on the amount of H$_2$ gas with T $>$500~K of less than a few Earth masses. Such warm gas temperatures are significantly higher than the equilibrium dust temperatures at these radii, suggesting that the gas is decoupled from the dust in the regions we are studying and that processes such as UV, X-ray, and accretion heating may be important. ",
        "introduction": "Studying the structure and evolution of circumstellar disks is crucial to developing an understanding of the process of planet formation. Observations of dust emission and modeling of the spectral energy distributions (SED) of disks have revealed much about the dust component from a few stellar radii out to hundreds of AU \\citep{zuckerman01}. While circumstellar disks are composed of both dust and gas, the gas component dominates the mass of the disk, with molecular hydrogen (H$_{2}$) being the most abundant constituent. In order to develop a complete picture of the structure and evolution of protoplanetary disks, it is important to observe the gas.  Observations at different wavelengths probe different disk radii. Submillimeter observations sample gas at large radii ($> 50$ AU) \\citep{semenov05}, while near-infrared CO \\citep{najita03,blake04} and H$_{2}$O \\citep{carr04, thi05} observations allow for study of the inner few AU. Spectral lines in the mid-infrared (5-25 $\\mu$m) provide a means to investigate gas in the giant planet region of the disk and beyond (10-50 AU) \\citep{najita07a}.  Several mid-infrared spectral diagnostics have been shown to be useful probes of gas in disks.  These include [NeII] at 12.8 $\\mu$m \\citep{pascucci07,lahuis07,herczeg07}, H$_{2}$O rotational transitions \\citep{carr08,salyk08}, [FeI] at 24 $\\mu$m \\citep{lahuis07}, and, based on a theoretical analysis of debris disks, [SI] at 25.2 $\\mu$m, and [FeII] at 26 $\\mu$m \\citep{gorti04}.  Molecular hydrogen should make up the bulk of the mass in disks, but is a challenge to detect. Bright far-ultraviolet (FUV) H$_{2}$ emission from classical T Tauri stars may be produced in the irradiated disk surface \\citep{herczeg02,bergin04}. At longer wavelengths, rovibrational and pure rotational transitions are generally weak because H$_2$ lacks a permanent dipole moment. Near-infrared emission in the \\textit{v}=1-0 \\sone ~rovibrational transition of H$_{2}$ has been detected from T Tauri stars \\citep{bary03,ramsay07,carmona08b} and may be the result of excitation by UV and X-ray irradiation \\citep{nomura07,gorti08}. Near-infrared adaptive optics fed, integral field spectroscopy of six classical T Tauri stars that drive powerful outflows has revealed that most of the H$_{2}$ emission is spatially extended from the continuum \\citep{beck08}. The properties of the emission are consistent with shock excitation from outflows or winds rather than UV or X-ray excitation from the central star. The FUV H$_{2}$ emission probes gas between 2000 and 3000 K \\citep{herczeg04,nomura05} and the 2.12 $\\mu$m H$_{2}$ line traces gas at T$>$1000 K \\citep{bary03}. The mid-infrared H$_{2}$ lines considered in this paper are most sensitive to gas at lower temperatures.  Owing to their small Einstein A-values, the pure rotational mid-infrared H$_{2}$ lines remain optically thin to large column densities (N$_{H_{2}} > 10^{23}$ cm$^{-2}$) and will be in LTE at the densities found in disks. For a disk with a strong mid-infrared continuum, the dust becomes optically thick well before the H$_{2}$ lines. Observable line emission is present only when there is a temperature inversion in the atmosphere of the disk or if there is a layer of dust-depleted gas separate from the optically thick dust. The process of dust coagulating into larger grains or settling out of the disk atmosphere can allow a larger column of gas to be observed. A disk that has an optically thin mid-infrared continuum, implying a very small amount of dust in the disk or dust grains that have grown large compared to mid-infrared wavelengths, would allow the entire disk to be observable. However, it is not known whether such disks have large quantities of gas and, in the absence of gas heating through collisions with dust grains, another heating mechanism is necessary in dust-free environments, such as UV or X-ray heating \\citep{glassgold04,gorti04,nomura07}.  A number of groups have searched for H$_{2}$ emission from protoplanetary disks in recent years. \\citet{thi01} reported the detection of several Jupiter masses of warm gas in a sample of disk sources based on \\textit{Infrared Space Observatory} (ISO) observations of the H$_{2}$ \\szero ~and \\sone ~lines. However, follow-up observations from the ground with improved spatial resolution did not confirm these results \\citep{richter02,sheret03,sako05}. \\citet{richter02} used the Texas Echelon Cross Echelle Spectrograph (TEXES) on the NASA 3m Infrared Telescope Facility (IRTF) to set upper limits on the warm gas mass within the disks of six young stars. \\citet{sheret03} searched for H$_{2}$ emission using MICHELLE on the United Kingdom Infrared Telescope and set upper limits on the emission from disks around two stars. The first group to use an 8-meter class telescope in the search for molecular hydrogen was \\citet{sako05} using the Cooled Mid-Infrared Camera and Spectrometer on the 8.2 m Subaru telescope to set upper limits for emission in the \\sone ~line around four young stars. Using the Infrared Spectrograph (IRS) aboard the \\textit{Spitzer Space Telescope}, \\citet{pascucci06} reported the non-detection of H$_{2}$ lines in their sample of 15 young Sun-like stars, while \\citet{lahuis07} detected the \\stwo ~and \\sthree ~lines in $\\sim$8\\% of the 76 circumstellar disks in their sample. Recently, ground-based observations with high resolution spectrometers on 8-meter class telescopes have produced both detections of the mid-infrared H$_{2}$ lines in the Herbig Ae stars AB Aur and HD 97048 \\citep{bitner07,martin07} and stringent upper limits in a sample of six Herbig Ae/Be stars and one T Tauri star \\citep{carmona08a}.  Three mid-infrared pure rotational H$_{2}$ lines are observable from the ground: \\sone ~($\\lambda = 17.035 ~\\mu$m), \\stwo ~($\\lambda = 12.279 ~\\mu$m), and \\sfour ~($\\lambda = 8.025 ~\\mu$m). When multiple optically thin lines are observed, line ratios permit the determination of the temperature and mass of the emitting gas. Ratios of these three lines are most sensitive to temperatures of 200-800 K. Two additional pure rotational H$_{2}$ lines are observable near 5 $\\mu$m: \\seight ~($\\lambda = 5.053 ~\\mu$m) and \\snine ~($\\lambda = 4.695 ~\\mu$m), extending our temperature sensitivity to hotter gas. The high spectral resolution possible with an instrument like TEXES (Lacy et al. 2002) increases our sensitivity to narrow line emission and helps determine the location of the emission. By making observations at high spectral resolution, we maximize the line to continuum contrast while minimizing atmospheric effects by separating the lines from nearby telluric features. A further benefit of high spectral resolution is that we are able to estimate the location of the emitting gas if coming from a disk under the assumption of Keplerian rotation.  In this paper, we present the results of a search for molecular hydrogen emission in disk sources using TEXES on both the IRTF and Gemini North telescopes. We observed 29 sources spanning a range of mass, age, and accretion rate in order to constrain the amount of warm gas in the circumstellar disks of these stars. ",
        "conclusions": "We have carried out a survey for pure rotational H$_{2}$ emission from the circumstellar environments surrounding a sample of 29 stars with disks and detected emission from 6. In the case of non-detections, our upper limits constrain the amount of T $>$ 500 K gas in the surface layers of the circumstellar disks to be less than a few Earth masses. Several objects in our survey have transition object SEDs implying the presence of an optically thin, dust-depleted inner disk: DoAr 21, GM Aur, HD 141569, and LkCa 15. Among these sources, only DoAr 21 shows H$_{2}$ emission and it appears to be far from the star. One possible explanation for transitional SEDs is grain growth \\citep{strom89,dullemond05} whereby the inner disk becomes optically thin yet remains gas rich. This gas could be heated through accretion as well as X-ray and UV heating. GM Aur is of particular interest since it has an accretion rate ($\\sim$10$^{-8}$ M$_\\odot$ yr$^{-1}$) similar to typical CTTS accretion rates so should have as much gas in the inner disk as a typical CTTS. If grain growth is responsible for the transitional SEDs of these sources, the available heating mechanisms are insufficient to produce detectable H$_{2}$ line emission. Alternatively, grain growth may not be a good explanation for a transition object SED, as suggested by other demographic data \\citep{najita07b}.  In all cases, the detected emission lines are narrow and centered at the stellar velocity. The narrow range of line widths, FWHM between 7 and 15 km s$^{-1}$, along with the fact that the line fluxes are all similar, suggests that the mechanism for exciting the emission may be the same in each case. Four of the six targets with detected emission are class I sources that show evidence for surrounding material in an envelope in addition to a circumstellar disk. It is possible, and likely in the case of HL Tau, that the H$_{2}$ emission we observe is a result of gas in the circumstellar envelope being shock heated by an outflow. However, the fact that all of the H$_{2}$ line centroids in our sample are within a few km s$^{-1}$ of the stellar velocity argues against this being the case for all of our detections.  Under the assumption of emission from a disk in Keplerian rotation, the narrow line widths imply that the emission arises at disk radii from 10-50 AU. At such large disk radii, additional heating of the gas besides heating due to collisions with dust grains is required to explain the temperatures derived from our H$_{2}$ observations. Both X-ray/UV irradiation of the disk surface layer and accretion shocks resulting from matter infall onto the disk are plausible candidates. With the exception of DoAr 21, all of the sources where we detect H$_{2}$ emission possess both a circumstellar disk and a surrounding envelope of material. This lends support to the possibility that the H$_{2}$ emission we observed may be the result of shocks in the disk due to infalling material.  Models of molecular hydrogen emission from disks that assume sufficient levels of stellar X-ray and UV irradiation \\citep{gorti08} predict line fluxes that are consistent with our observations. In contrast, models which assume smaller values of stellar UV and X-ray irradiation \\citep{nomura07} produce weaker H$_{2}$ emission than observed in our sample. We looked for evidence of a correlation between X-ray/UV luminosity and the presence of H$_{2}$ emission but found none. We note that the X-ray and UV luminosities used for the purpose of searching for a correlation with H$_{2}$ emission were not measured at the same time. To definitively test for a correlation between X-ray/UV luminosity and the presence of H$_{2}$ emission will require a series of coordinated observations."
    },
    "0808/0808.1785_arXiv.txt": {
        "abstract": "In this paper, using 2MASS photometry, we study the structural and dynamical properties of four young star clusters viz. King~16, NGC~1931, NGC~637 and NGC~189. For the clusters King~16, NGC~1931, NGC~637 and NGC~189, we obtain the limiting radii of $7'$, $12'$, $6'$ and $5'$ which correspond to linear radii of 3.6~pc, 8.85~pc, 3.96~pc and 2.8~pc respectively. The reddening values $E(B-V)$ obtained for the clusters are 0.85, 0.65--0.85, 0.6 and 0.53 and their true distances are 1786~pc, 3062~pc, 2270~pc and 912~pc respectively. Ages of the clusters are 6~Myr, 4~Myr, 4~Myr and 10~Myr respectively. We compare their structures, luminosity functions and mass functions ($\\phi(M) = dN/dM \\propto M^{-(1+\\chi)}$) to the parameter $\\tau = t_{age}/t_{relax}$ to study the star formation process and the dynamical evolution of these clusters. We find that, for our sample, mass seggregation is observed in clusters or their cores only when the ages of the clusters are comparable to their relaxation times ($\\tau \\geq 1$). These results suggest mass seggregation due to dynamical effects. The values of $\\chi$,  which characterise the overall mass functions for the clusters are 0.96 $\\pm$ 0.11, 1.16 $\\pm$ 0.18, 0.55 $\\pm$ 0.14 and 0.66 $\\pm$ 0.31 respectively. The change in $\\chi$ as a function of radius is a good indicator of the dynamical state of clusters.   \\bf{star clusters: young -- near-infrared photometry -- color--magnitude diagrams -- pre-mainsequence stars -- initial mass function--relaxation time-- 2MASS} ",
        "introduction": "Star clusters are the most fundamental stellar sytems in understanding the star formation process, which is still full of mysteries. A good study of these stellar systems is the basis of understanding galaxies which are the larger building blocks of the universe \\citep{lynga82,janes81,friel95,bonatto05,piskunov06}.  \\end{abstract} Homogeneous samples of photometric data, coupled with uniform methods of data analysis are essential to make statistical inferences based on the fundamental parameters of clusters. These studies can contribute to understanding the galactic disk, formation and evolution of clusters, molecular cloud fragmentation, star formation and evolution. In this work, we study a sample of young clusters viz.  King~16, NGC~1931, NGC~637 and NGC~189  using photometric data from the Two~Micron~All~Sky~Survey (2MASS) \\citep{Skrutskie06}.  The 2MASS covers 99.99\\% of the sky in the near-infrared $J$~(1.25~$\\mu$m), $H$~(1.65~$\\mu$m) and $K_{s}$~(2.16~$\\mu$m) bands (henceforth $K_{s}$ shall be refered to as $K$). Hence the 2MASS database has the advantages of being homogeneous, all sky (enabling the study of the outer regions of clusters where the low mass stars dominate) and covering near infrared wavelengths where young clusters can be well observed in their dusty environments. \\cite{dutra01} discovered 42 objects at infra-red wavelengths using the 2MASS survey. Many papers devoted to the study of clusters using the 2MASS have been presented in the past few years \\citep{bica03,bica06,tadross07,kim06} showing the potential of this database.  We use the 2MASS database to study the sample of four young clusters, to investigate the structure and dynamical state of these clusters close to their time of formation. The sample is selected on the basis that all the four clusters have an age of $\\approx$10~Myr reported in literature (see Table 1) and have been formed in different environments. We study the structures and dynamical states of our sample of clusters and determine their mass functions (MFs) and degree of mass seggregation in various regions of the clusters. To study the sample, we construct radial density profiles (RDPs), color--magnitude diagrams (CMDs),  color--color diagrams, luminosity functions (LFs) and MFs. Such studies are not possible using heterogenous datasets where unknown biases may be present.  The initial mass function (IMF) is the distribution of stars of varying masses from the original parent cloud. The universality of the IMF and the influence of environment on star formation is still a matter of debate. As these clusters are very young (age $\\leq$ 10 Myr), their MF may be approximated as the IMF. The sample of clusters have been made from differing initial conditions and subject to varied influences of external interactions with the galactic field, thus leading to observable differences, which we explore.  However, from a recent study of \\cite{kroupa07}, even in the case of very young clusters, there is a change in the MF due to the dynamics of young clusters which loose a significant fraction of their stars at an early age.  Mass segregation is the redistribution of stars according to their masses, thus leading to the concentration of high mass stars near the centre and the low mass ones away from the centre. This has been observed in a variety of clusters, both young and old. The variation of the MF of these clusters is determined in different regions of the clusters and their values are compared.  Further, we estimate from the value of $\\tau = t_{age}/t_{relax}$, the degree of mass segregation expected due to dynamical effects and compare it with our observations. The relaxation time $t_{relax}$ is a characteristic time in which there is an equipartition of energy and the high mass stars with lesser kinetic energy sink to the core and the low mass stars move to the outer regions of the cluster \\citep{binney08}.  The value of $\\tau$  indicates whether an excess of high mass stars in the cores of clusters is a result of dynamical evolution or the imprint of the star formation process itself. The parameter $\\tau$ has been described as an evolutionary parameter \\citep{bonatto05} which indicates the extent to which the cluster has relaxed. It relates to the core and overall MF flattening.  For large values of $\\tau$,  the high mass stars sink to the centre and the low mass stars with high velocities move towards the outskirts and hence the MFs of clusters show large-scale mass segregation and low-mass stars evaporation. We report the presence of gaps in the main sequence associated to physical processes in stars \\citep{kjeldsen91}.    The plan of the paper is as follows: Section 2 describes the clusters in our sample and shows the corresponding RDPs and the values obtained for the limiting radii for these clusters. Section 3 describes the method of selecting cluster members  and the corresponding values of fundamental parameters obtained.  LFs and MFs are described in Section 4 and a comparative study of these clusters is in the concluding Section 5. ",
        "conclusions": "In this paper we have studied four young clusters of comparable ages to understand their structure and dynamics. The RDPs of the clusters have been plotted and the parameters for the clusters such as reddening, distance and age have been determined using isochrone fits.  We have also plotted the LFs in the $J$, $H$ and $K$ bands and used the derived mass--luminosity relation to find the MFs using all three bands independently. The $\\chi$ values have been determined for different regions and the overall clusters as a function of the parameter $\\tau$. We use the difference in $\\chi$ values to estimate the level of mass segregation of the clusters and their cores.  In the case of King~16, where $\\tau$ = 3  for the core, the core is clearly relaxed as in indicated by its flat $\\chi=-0.44$, while for the outer regions where  $\\tau$ = 0.57, the cluster has begun to relax. In the case of NGC~1931, the core is relaxed and the outer region  seems to have begun relaxation. In the case of NGC~637, the core is relaxed ($\\tau=4/0.039$), and there  appears to be a redistribution of stars in the cluster, indicated by a progressive increase in $\\chi$ from 0.39 to 1.15. The outer halo has a larger value of $\\chi$ compared to the overall cluster (0.55) as, it appears that the inner halo has thrown away a large number of low mass stars to the outer halo, thus causing an excess in low mass stars and a steeper $\\chi$. In the case of NGC~189, where $\\tau$ =2.79, the cluster has already relaxed, as is indicative of the flat $\\chi$ of the core. The outer core has a flatter $\\chi$, probably because low mass stars which were thrown out of the inner halo during relaxation, have been lost and hence the outer halo has a deficit of low mass stars.  As, seen from our analysis, mass seggregation is observable in the cores of the clusters King~16, NGC~1931 and   NGC~637  where the cluster ages are comparable to the relaxation times. In the case of NGC 189, where the relaxation time is lesser than the age of the cluster, we notice a flatter mass function. This  implies that the observed mass seggregation in these clusters is a dynamical effect. However, larger samples will improve the statistics and give us a better insight in the physical processes leading to the structure and dynamical evolution of clusters."
    },
    "0808/0808.2053_arXiv.txt": {
        "abstract": "Infrared Dark Clouds (IRDCs) represent the earliest observed stages of clustered star formation, characterized by large column densities of cold and dense molecular material observed in silhouette against a bright background of mid-IR emission. Up to now, IRDCs were predominantly known toward the inner Galaxy where background infrared emission levels are high. We present {\\it Spitzer} observations with the Infrared Camera Array toward object G111.80+0.58 (G111) in the outer Galactic Plane, located at a distance of $\\sim$3\\,kpc from us and $\\sim$10\\,kpc from the Galactic center. Earlier results show that G111 is a massive, cold molecular clump very similar to IRDCs. The mid-IR {\\it Spitzer} observations unambiguously detect object G111 in absorption. We have identified for the first time an IRDC in the outer Galaxy, which confirms the suggestion that cluster-forming clumps % are present throughout the Galactic Plane. However, against a low mid-IR background such as the outer Galaxy it takes some effort to find them. ",
        "introduction": "Massive stars are believed to form almost exclusively in stellar clusters \\citep{Blaauw1964,deWit2005}, where supposedly the majority of all stars in the Galaxy form. The precursors to these stellar clusters are cold, dense and, most importantly, massive molecular clumps. A major, unexpected and exciting result from the {\\it Midcourse Space Experiment} ({\\it MSX}) and {\\it Infrared Space Observatory} ({\\it ISO}) is the discovery of dark clouds seen in silhouette against the bright mid-IR background toward the inner Galactic Plane \\citep{MSXIRDC1998,ISOCAM1996}. In contrast to the well-known dark clouds seen in visual extinction against the stellar distribution, these so-called Infrared Dark Clouds (IRDCs) have much higher column densities ($\\gtrsim$\\,10$^{23}$\\,cm$^{-2}$), are % at large distances ($\\gtrsim$\\,1\\,kpc) and hence are typically much more massive. Follow-up studies \\citep[e.g.][]{Carey1998,Teyssier2002,Simon2006co} show that they are cold ($<$\\,20\\,K), dense ($\\sim$\\,10$^5$\\,cm$^{-3}$) and, indeed, among the most massive molecular clumps yet found in our Galaxy ($M$$\\sim$100--10$^5$\\,M$_{\\odot}$). Many IRDCs contain compact (sub-) millimeter cores \\citep[$M$$\\sim$10--10$^3$\\,\\msun, $n$$\\sim$10$^3$--10$^7$\\,cm$^{-3}$, $T$$\\sim$15--30\\,K, e.g.,][]{Carey2000,Teyssier2002,Garay2004,Rathborne2006}. While these cores were first considered to represent the pre-stellar core phase of clustered star formation, recent studies show that at least some cores in IRDCs contain % proto-stellar objects \\citep[e.g.,][]{Redman2003,Ormel2005,Rathborne2005}. Nevertheless, IRDCs are likely to represent the earliest stages of clustered star formation \\citep[e.g.,][]{Menten2005,Rathborne2006} and they are generally referred to as cluster-forming clumps. Their ambient physical conditions can illuminate the role of IRDCs in the star forming process and can provide insight into the differences between low- and high-mass star formation, the nature of the initial mass function and the impact of environment on star formation. Moreover, their % core forming properties may reveal the important mechanisms involved in forming high-mass stars, e.g., competitive accretion, merging of lower-mass stars or massive disk accretion \\citep[e.g.,][and references therein]{Larson2007}.  The identification of IRDCs is by necessity biased toward cold molecular clouds having sufficient column density against a bright mid-IR background, in order to be seen as absorption features. Hence, identifying IRDCs is considerably easier toward the mid-IR bright spiral arms and molecular ring in the inner Galactic Plane \\citep{Simon2006msx}, i.e., within 90\\degree\\ from the Galactic Center. An unbiased, more complete understanding of clustered star formation and its dependence on environment and other external properties, e.g., interstellar radiation field, metallicity, external pressure and dynamics, requires a study of IRDC-like objects in more quiescent regions, such as the outer Galaxy, i.e., between $l$=90\\degree and $l$=270\\degree.  Frieswijk et al. (2008) produce a catalog of extended red objects in the outer Galactic Plane. The objects are identified in the {\\it Two Micron All Sky Survey Point Source Catalog} ({\\it 2MASS,} \\citealt{Skrutskie2006}) as clusters of stars that show a statistically redder color distribution compared to their local surroundings. A sample, consisting of 414 objects covered by the Canadian Galactic Plane Survey \\citep[CGPS,][]{Heyer1998}, was investigated in detail using the CO data of the CGPS. Over 90\\% of the objects correlate morphologically with CO emission with a typical FWHM of the order of 5\\,km\\,s$^{-1}$. This suggests that the observed reddening is due to the presence of foreground molecular clouds. The kinematic distance of the clouds, derived from the CO radial velocity, indicates that $\\sim$15\\% are located beyond $\\sim$\\,3\\,kpc from us. Some of these distant clouds ($\\sim$10) do not show significant mid- and far-IR emission in, e.g., MSX and IRAS, and are potentially cold, cluster-forming clumps in the outer Galaxy, or `IRDC-like' objects. \\citet[][F07]{Frieswijk2007AA} studied the physical characteristics of a candidate IRDC-like object (G111) by using molecular line observations. At a radial velocity of -52\\,km\\,s$^{-1}$, G111 is located at a kinematic distance of $\\sim$\\,5\\,kpc, assuming IAU standard constants. However, within 15\\arcmin\\ on the sky and at similar radial velocity are the (massive) star forming regions NGC 7538 and S 159. These are part of the Cas OB2 complex in the Perseus spiral arm, located at $\\sim$3\\,kpc from us and $\\sim$10\\,kpc from the Galactic center. Hence, a distance of 3\\,kpc is assumed for G111. A detailed investigation of the molecular lines conducted by F07 revealed multiple cold and massive cores (P1--P10, see Fig.\\ref{fig1} and Sec.\\ref{conclusions}) along a filamentary cloud structure. The main conclusion is that G111 has global properties in agreement with values found for IRDCs, e.g., $M$\\,$\\gtrsim$\\,3000\\,\\msun, $N$\\,$\\gtrsim$\\,10$^{23}$\\,cm$^{-2}$, $T$\\,$\\lesssim$\\,20\\,K. The mass and density of individual cores are low when compared to sub-millimeter cores in IRDCs. However, with peak column densities in excess of 10$^{23}$\\,cm$^{-2}$, the most massive cores in G111 should, provided that sufficient mid-IR background is present, be observable as absorption features in a similar way to inner Galaxy IRDC cores.  We subsequently requested {\\it Spitzer} observations toward G111 to test the idea that the cores can be observed in absorption. {\\it MSX} observations were not able to resolve this, presumably because the low sensitivity did not allow a detection of contrast between the molecular cloud and the background. With the sensitivity of the Infrared Camera Array (IRAC) on {\\it Spitzer} we argued that it should be possible to observe this contrast. In this paper, we present the first outcomes of the data. Combined with the results in F07, we conclude that we have, beyond doubt, a detection of an IRDC in the outer Galaxy. The paper is organized as follows: Section \\ref{observations} gives % an overview of the observations. Section \\ref{results} presents the results. In Section \\ref{conclusions} we discuss the observations and give our interpretation of the data % compared to the results published in F07. We conclude the section with some future prospects. ",
        "conclusions": "\\label{conclusions} The mid-IR properties of G111 compare well to other IRDCs. The {\\it MSX} survey data have been used by \\citet{Simon2006msx} to identify a large number of IRDCs in the inner Galaxy by their mid-IR contrast relative to the background. A typical 8\\,$\\mu$m extinction observed toward IRDCs, determined from the ratio between the on- and off-source emission after foreground and zodiacal light correction, is 1--2\\,mag \\citep[e.g.,][]{MSXIRDC1998,Carey2000}. We find extinction values toward G111 to be in agreement with this, at least toward the darkest cores. More so, the extinction measurement gives only a strict lower limit to the column of material. The cloud does not only absorb background emission, but re-radiates low-level emission as well. For a cold dark cloud ($<$25\\,K) this will be negligible. However, once embedded objects start heating parts of the cloud it may become important. The contrast between the cloud and the background, and thus the measured extinction, decreases as the cloud emits, but of course the column of material does not. The molecular column densities of positions P1--P10 are $\\sim$10$^{22}$\\,cm$^{-2}$, peaking at $\\approx$\\,7$\\times$10$^{22}$\\,cm$^{-2}$ for P8 (F07). The latter corresponds to 2.8\\,mag extinction at 8\\,$\\mu$m (70\\,mag in $A_V$)\\footnote{This is a factor of 2 above the value in F07, because of a too low conversion factor used by F07.}. These estimates demonstrate that parts of this cloud have a mid-IR extinction comparable to values found for inner Galaxy IRDCs \\citep[e.g.,][]{Carey1998,Hennebelle2001}.  Recall in this that the selection of IRDCs from {\\it MSX} data is based on high contrast \\citep{Simon2006msx}. In the outer Galaxy, a high contrast is not guaranteed because of the low-level background emission. For example, compare the background of the IRDCs toward the edge of W51, i.e., $\\sim$\\,60\\,MJy\\,sr$^{-1}$ (IRAC/MIPS cycle 1 observations, K. Kraemer et al., priv. comm.) with the background that we observe ($\\sim$\\,20\\,MJy\\,sr$^{-1}$). Note that one bright 8\\,$\\mu$m emission feature clearly is in the background of G111, since a dark lane is visible in front of IRAS23136+6111 (see Figure \\ref{fig1}, lower blow-up).  The four-color images show a clustering of star-like objects near the dark filaments. Though these stars could coincidently be in the foreground, this positional consistency suggests that some cores might have started to form stars. The excess in steller surface density was also pointed out by F07 toward P5 and P8, based on the 2MASS data. Note that IRDCs are often associated with active star formation \\citep[e.g.,][]{Redman2003,Rathborne2005}. Star forming activity may further reveal itself by the presence of `green fuzzy emission', i.e., as weak extended 4.5\\,$\\mu$m features. This emission is often attributed to shocked gas arising from outflow-activity, e.g., the pure rotational H$_2$ lines S(11) at 4.18\\,$\\mu$m through S(9) at 4.69\\,$\\mu$m and the ro-vibrational lines of CO at 4.45--4.95\\,$\\mu$m  \\citep[e.g.,][]{Noriega2004,Marston2004}. Shocked gas features are, besides for nearby star formation, also frequently observed toward IRDCs \\citep[e.g.,][]{Rathborne2005,Beuther2007}. A first impression of the emission at 4.5\\,$\\mu$m suggests that such features are present toward some cores in G111. However, further investigation is required to confirm this. \\\\ \\ \\\\ \\textbf{Notes on individual cores}\\\\ The positions P1 to P10 in Figure \\ref{fig1} were selected by F07 based on their high column density. A detailed investigation showed that they represent cold (10--20\\,K), dense ($>$10$^3$\\,cm$^{-3}$) and massive cores ($\\sim$100\\,\\msun, P8$\\sim$\\,1000\\,\\msun) where stars might form. Star forming activity was investigated by means of 2MASS star-colors and -counts and the presence of warm gas (NH$_{3}$, $^{13}$CO). The following discussion summarizes some of the results in F07 and includes the mid-IR characteristics of the cores presented in this paper.  P1, P2, P3 and P7 do not show indications of active star formation. P1 shows enhanced 8\\,$\\mu$m emission, which may be a chance encounter along the line of sight. The other cores show a decrease in 8\\,$\\mu$m emission, corresponding to an $A_V$ of $>$10\\,mag. The peak extinction appears somewhat offset from the C$^{18}$O peak, which is true for most cores. This may be explained by C$^{18}$O freeze-out in the densest regions. P4 is not a core but part of the C$^{18}$O filament extending to the east. Several extinction peaks at 8\\,$\\mu$m ($A_V$$\\gtrsim$10\\,mag) are present. P9 and P10 are not evident as cores in C$^{18}$O and show no specific features at 8\\,$\\mu$m. P6 shows a decrease of 8\\,$\\mu$m emission ($A_V$$\\gtrsim$7\\,mag), but the contrast with the background may be reduced due to the presence of the bright source near the C$^{18}$O peak. Whether this source is physically associated cannot be determined at present. P5 and P8 show signs of active star formation. The NH$_{3}$ lines indicate the presence of warm gas. Both cores have associated 2MASS sources with typical near-IR colors of YSOs. Furthermore, the stellar surface density of 2MASS peaks at a value almost twice that of the surrounding field for P8. These results are supported by the four-color images that reveal signs of active star formation in the form of clustering of objects and 4.5\\,$\\mu$m `green fuzzes'. The 8\\,$\\mu$m contrast indicates peak extinctions of at least 12.5\\,mag in $A_V$ for both cores. IRAS23136+6111 was not a target position in F07, but is nonetheless interesting to include here. C$^{18}$O emission reveals a narrow filamentary structure. The four-color image depicts a dark lane and a clustering of point sources matching the location of the filament. We thus conclude that it must be in the foreground of the bright IRAS object. Further investigation is required to see if the dark lane and the IRAS source are associated, which may indicate a triggered star forming event.  To summarize, the {\\it Spitzer} observations support the results in F07. Except P5 and P8, where star formation may have started, the cores appear to be in a cold, pre-stellar phase. \\\\ \\ \\\\ \\textbf{Future prospects}\\\\ Supplemental IRAC and MIPS 24\\,$\\mu$m data from {\\it Spitzer} are expected. Combining the IRAC bands with the 2MASS data enables a discrimination between stars and YSOs for many of the clustered point sources in the cloud vicinity. Adding 24\\,$\\mu$m data will allow a proper SED modeling of the dust emission. The characteristic features of envelopes, disks and photospheres in the SED will enable a determination of the evolutionary state, i.e., Class 0, I, II or III, of the apparent discrete sources \\citep[e.g.,][]{Robitaille2007}. Star forming activity can be further characterized through an analysis of the spatial distribution and strength of polycyclic aromatic hydrocarbon emission (e.g., 3.3, 6.2 and 7.7\\,$\\mu$m covered by the 3.6, 5.8 and 8\\,$\\mu$m bands, respectively) and shock activity, exemplified by the 4.5\\,$\\mu$m features in the blow-up images in Figure \\ref{fig1} \\citep[e.g., bow shocks;][]{Neufeld2006,Velusamy2007}. Considering the distance to the cloud, high-resolution continuum and spectroscopic observations (near-IR to millimeter) are essential to improve on the spatial information."
    },
    "0808/0808.2323_arXiv.txt": {
        "abstract": "\\noindent We discuss the 21cm power spectrum (PS) following the completion of reionization. In contrast to the reionization era, this PS is proportional to the PS of mass density fluctuations, with only a small modulation due to fluctuations in the ionization field on scales larger than the mean-free-path of ionizing photons. We derive the form of this modulation, and demonstrate that its effect on the 21cm PS will be smaller than 1\\% for physically plausible models of damped Ly$\\alpha$ systems. In contrast to the 21cm PS observed prior to reionization, in which HII regions dominate the ionization structure, the simplicity of the 21cm PS after reionization will enhance its utility as a cosmological probe by removing the need to separate the PS into physical and astrophysical components. As a demonstration, we consider the Alcock-Paczynski test and show that the next generation of low-frequency arrays could measure the angular distortion of the PS at the percent level for $z\\sim3-5$. ",
        "introduction": "Recently, there has been much interest in the feasibility of mapping the three-dimensional distribution of cosmic hydrogen through its resonant spin-flip transition at a rest-frame wavelength of 21cm~(Furlanetto, Oh \\& Briggs~2007; Barkana \\& Loeb~2007). Several experiments are currently being constructed (including MWA~\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/}, LOFAR~\\footnote{http://www.lofar.org/}, PAPER ~\\footnote{http://astro.berkeley.edu/~dbacker/EoR/}, 21CMA~\\footnote{http://web.phys.cmu.edu/~past/}, GMRT~\\footnote{Pen et al.~(2008)}) and more ambitious designs are being planned (SKA~\\footnote{http://www.skatelescope.org/}).  One driver for mapping the 21cm emission is the possibility of measuring cosmological parameters from the shape of the underlying power spectrum (PS; see Loeb \\& Wyithe 2008). During the epoch of reionization, the PS of 21cm brightness fluctuations is shaped mainly by the topology of ionized regions, rather than by the PS of matter density fluctuations which is the quantity of cosmological interest~(McQuinn et al.~2006; Santos et al.~2007; Iliev et al.~2007). As a result, the line-of-sight anisotropy of the 21cm PS due to peculiar velocities must be used to separate measurements of the density PS from the unknown details of the astrophysics (Barkana \\& Loeb~2005; McQuinn et al.~2006). The situation is expected to be simpler both prior to the formation of the first galaxies (at redshifts $z\\ga 20$, Loeb \\& Zaldarriaga~2004; Lewis \\& Challinor~2007; Pritchard \\& Loeb~2008), and following reionization of the intergalactic medium (IGM; $1\\la z\\la6$) -- when only dense pockets of self-shielded hydrogen, such as damped Ly$\\alpha$ absorbers (DLA) and Lyman-limit systems (LLS) survive ~(Wyithe \\& Loeb~2008; Chang et al.~2007; Pritchard \\& Loeb~2008).   In this paper we focus our discussion on the post-reionization epoch (Wyithe \\& Loeb~2007; Chang et al.~2007).  The DLAs which contain most of the neutral hydrogen mass in the Universe at $z\\la 6$ are expected to be hosted by galactic mass dark matter halos~(Wolfe, Gawiser \\& Prochaska~2005). A survey of 21cm intensity fluctuations after reionization would measure the modulation of the cumulative 21cm emission from a large number of galaxies~(Wyithe \\& Loeb~2008; Wyithe, Loeb \\& Geil~2008; Chang et al.~2007). Regarding the measurement of the 21cm PS, this lack of identification of individual galaxies is an advantage, since by not imposing a minimum threshold for detection, such a survey collects all the available signal. This point is discussed in Pen et al.~(2008), where the technique is also demonstrated via measurement of the cross-correlation of galaxies with unresolved 21cm emission in the local Universe.  Studying the 21cm PS after (rather than during) reionization offers two advantages.  First, it is less contaminated by the Galactic synchrotron foreground, whose brightness temperature scales with redshift as $(1+z)^{2.6}$ ~(Furlanetto, Oh \\& Briggs~2006). Second, because the UV radiation field is nearly uniform after reionization, it should not imprint any large-scale features on the 21cm PS that would mimic the cosmological signatures.  In addition, on large spatial scales the 21cm sources are expected to have a linear bias analogous to that inferred from galaxy redshift surveys.   Most previous studies of post-reionization 21cm fluctuations have assumed that the 21cm emission traces perturbations in the matter density, and have considered peculiar motions only after a spherical average~(Wyithe \\& Loeb~2008; Pritchard \\& Loeb~2008, Loeb \\& Wyithe~2008).  The exception is a recent paper discussing measurements of the dark energy equation of state (Chang et al.~2007). Since the neutral gas resides within collapsed dark matter halos, galaxy bias plays an important role in setting the 21cm fluctuation amplitude.  In this paper we derive the 21cm PS after reionization in the context of the formalism that has been developed to calculate it during reionization~(Barkana \\& Loeb~2005; McQuinn et al.~2006; Mao et al.~2008). By framing the derivation this way, the relative merits of cosmological constraints from 21cm surveys at redshifts before and after reionization can be more easily understood. In addition, this formalism provides a framework to describe the possible effect of fluctuations in the ionizing background. We then compute the Alcock-Paczynski effect~(Alcock \\& Paczynski~1979) as an example for the cosmological utility of the post-reionization 21cm PS. In our numerical examples, we adopt the standard set of cosmological parameters ~(Komatsu et al.~2008), with values of $\\Omega_{\\rm b}=0.24$, $\\Omega_{\\rm m}=0.04$ and $\\Omega_Q=0.76$ for the matter, baryon, and dark energy fractional density respectively, and $h=0.73$, for the dimensionless Hubble constant. ",
        "conclusions": "In this paper we have derived the 21cm power spectrum (PS) following the completion of reionization. Our approach is to derive the PS within the formalism that has been developed to calculate the 21cm PS during reionization~(Barkana \\& Loeb~2005; McQuinn et al.~2006; Mao et al.~2008). By framing the derivation this way we are able to directly compare the relative merits of cosmological constraints from 21cm surveys at redshifts prior to and post reionization. We have derived expressions for the components of the post-reionization 21cm PS that are due to the density field, the ionization field, and their cross-correlation. As is the case prior to reionization, we find that these components contribute to the observed PS in proportions that depend on the angle relative to the line-of-sight along which the power is measured. However, in difference from the situation prior to reionization, we have shown that all components of the 21cm PS are directly proportional to the PS of the underlying matter fluctuations, with a small but predictable modulation on scales below the mean-free-path of ionizing photons. We have derived the form of this modulation, and have shown that its effect on the observed PS will be at less than the 1\\% level for physically plausible astrophysical models of DLA systems. The simplicity of the 21cm PS after reionization stands in contrast to the astrophysical uncertainty during the reionization epoch, where HII regions dominate the 21cm signal. This simplicity will enhance the utility of the 21cm PS after reionization as a cosmological probe (Loeb \\& Wyithe 2008) by removing the need to separate the PS into physical and astrophysical components (Barkana \\& Loeb 2005).  To illustrate the utility of the 21cm PS after reionization as a cosmological probe, we have examined the Alcock-Paczynski~(1979) test. Our calculations show that the next generation of low-frequency arrays could measure the angular distortion of the PS to around $\\sim 1\\%$ at $z\\sim3.5$. It has previously been shown that the scale of Baryonic Acoustic Oscillations, which constitutes a standard ruler ~(Blake \\& Glazebrook~2003; Seo \\& Eisenstein~2005; Eisenstein et al.~2005; Padmanabhan et al.~2007), can be also be used to independently probe $H$ and $D_{\\rm A}$ in 21cm PS, and hence to measure the equation of state of the dark energy ~(Wyithe, Loeb \\& Geil~2008; Chang et al.~2007). With the caveat that non-linear evolution of the 21cm PS must be quantitatively understood, our simple analysis indicates that the precision achievable via the Alcock-Paczynski~(1979) test using the 21cm PS after reionization could be better than those available via a 21cm fluctuation measurement of the acoustic scale of baryonic oscillations for a given observing strategy (Shoji, Jeong \\& Komatsu 2009). More generally, our analysis shows that the 21cm PS after reionization shares the same favorable features as a galaxy redshift survey. The advantage of using 21cm fluctuations lies in the fact that individual sources need not be resolved. This would allow PS measurements using 21cm fluctuations to be extended to higher redshifts.     \\bigskip  \\noindent {\\bf Acknowledgments.} The research was supported by the Australian Research Council (JSBW), by NASA grants NNX08AL43G and LA, by FQXi, and by Harvard University funds (AL). We thank an anonymous referee for correcting an earlier error."
    },
    "0808/0808.0921_arXiv.txt": {
        "abstract": "We present a high signal-to-noise spectrum of a bright galaxy at $z$ = 4.9 in 14 h of integration on VLT FORS2.  This galaxy is extremely bright, $i_{850} = 23.10 \\pm 0.01$, and is strongly-lensed by the foreground massive galaxy cluster Abell 1689 ($z=0.18$).  Stellar continuum is seen longward of the Ly$\\alpha$ emission line at  $\\sim7100$ \\AA, while intergalactic H~I produces strong absorption shortward of Ly$\\alpha$. Two transmission spikes at $\\sim$6800 \\AA \\ and $\\sim$7040 \\AA \\ are also visible, along with other structures at shorter wavelengths. Although fainter than a QSO, the absence of a strong central ultraviolet flux source in this star forming galaxy enables a measurement of the H~I flux transmission in the intergalactic medium (IGM) in the vicinity of a high redshift object.  We find that the effective H I optical depth of the IGM is remarkably high within a large 14 Mpc (physical) region surrounding the galaxy  compared to that seen towards QSOs at similar redshifts.  Evidently, this high-redshift galaxy is located in a region of space where the amount of H~I is much larger than that seen at similar epochs in the diffuse IGM.  We argue that observations of high-redshift galaxies like this one provide unique insights on the nascent stages of baryonic large-scale structures that evolve into the filamentary cosmic web of galaxies and clusters of galaxies observed in the present universe. ",
        "introduction": "Hydrodynamic simulations tell us that dark matter near the epoch of galaxy formation collapses into an ordered filamentary pattern, the so-called ``cosmic web.\" In turn, this cosmic web is thought to cradle high-redshift galaxies in dense nodes that are opaque to the extragalactic ultraviolet background radiation. Observations of H~I surrounding these high-redshift objects provide information on the likely reservoirs from which galaxies assemble their gas.  Hitherto, it has only been possible to measure H~I opacities towards bright Quasi-stellar Objects (QSOs); unfortunately, the high ultraviolet flux from the QSOs  ionizes the hydrogen clouds in their vicinity, thereby making  the determination of H~I cloud physical conditions unreliable close to the QSO.  New techniques and facilities in the past decade have enabled detailed observational studies of the IGM in the early universe ($z > 5$).  For example, the spectra of high-redshift QSOs show a plethora of H~I Ly$\\alpha$ absorption lines, the ``Ly$\\alpha$ forest\", as well as lines from a wide variety of heavier elements, all of which provide detailed information about the high-redshift IGM along the line-of-sight towards the QSO. Around the QSOs themselves, however, the ultraviolet flux is so high that it typically ionizes the surrounding gas.   Thus QSOs are not ideal for studying the pervasive IGM close to the QSO. Instead, we suggest that bright galaxies may provide better probes for the study of IGM conditions in the proximity of high-redshift objects, as these objects emit fewer ultraviolet photons than QSOs do.  The subject of this paper is the Ly$\\alpha$ forest in the spectrum of an unusually bright high-redshift galaxy at $z = 4.866$ (Frye et al. 2002, 2007). The galaxy is situated behind the massive galaxy cluster Abell 1689 ($z = 0.183$), which magnifies the starlight of this background object by a factor of 10.3 by the effect of gravitational lensing (Broadhurst et al. 2005). Multicolor Hubble Space Telescope (HST) imaging of one of the faint lensed images, hereafter designated $A1689\\_7.1$, is shown in Figure 1. We assume a cosmology for this paper of $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m,0} = 0.3$, and $\\Omega_{\\Lambda,0} = 0.7$.  \\begin{figure} \\includegraphics[viewport=60 -10 250 350,scale=0.7]{f1_small.pdf} \\caption{HST Advanced Camera for Surveys (ACS) $gri$ optical photo of  $A1689\\_7.1$ and surrounding region near the center of the galaxy cluster Abell 1689 ($z=0.18$). A single galaxy at $z=4.9$ has been gravitationally-lensed by this foreground galaxy cluster (with member galaxies visible here as yellow spheroid-shaped objects) into a system of three arclets; $A1689\\_7.1$ is the brightest of the three and is shown above and in the inset.  This arclet is extremely bright, $i_{775}=23.10 \\pm 0.01$ magnitudes, after being magnified by a factor of $10.3$.  $A1689\\_7.1$  is also spatially resolved, with an estimated unlensed angular size of 0.094\\arcsec,  corresponding to an intrinsic linear size of only 600 pc.  $A1689\\_7.1$ was discovered as a part of a ground-based galaxy redshift survey (Frye et al. 2002, 2007). \\label{fig_image}} \\end{figure}  \\begin{figure} \\includegraphics[viewport = 0 0 700 520, scale=0.35]{f2.pdf} \\caption{VLT FORS2 spectrum of  $A1689\\_7.1$ at $z=4.866$.  The stellar continuum of the galaxy is seen longward of the Ly$\\alpha$ emission line at $\\sim7100$ \\AA, while intergalactic H~I produces strong absorption shortward of Ly$\\alpha$. Two transmission spikes at $\\sim$6800 \\AA \\ and $\\sim$7040 \\AA \\ are also visible, along with other structures at shorter wavelengths.  The data enable a measurement of the flux transmission along the line of sight to this star-forming galaxy. The best-fit \\citet{Bruzual:03} model spectrum without corrections for dust extinction and attenuation by the intervening Ly$\\alpha$ forest is overlaid, including a gray region marking our 1$\\sigma$ continuum-placement uncertainties. We select a Chabrier inital mass function (Padova 1994), stellar evolution tracks, and solar metallicity for the model, and give the details regarding the fitting procedure in \\S4. The inset figure shows a theoretical  H~I Ly$\\alpha$ profile absorption fit to the data, with horizontal bars delimiting the physical extents of the proximity zone for a QSO (red solid line), and the extent of the first of the five equally-spaced redshift bins in our study over which we measure the H~I opacity (blue-dashed line), respectively.  Owing to the lack of a significant proximity zone for our galaxy, a new physical scale is probed close to a high redshift object within a $\\sim14$ physical Mpc radius (blue dashed line).\\label{fig_spec }} \\end{figure} ",
        "conclusions": "Despite the strong magnification, the intrinsic luminosity of the galaxy is not high; we measure an unabsorbed luminosity of $L_{1400} = 7.7 \\times 10^{28}$ ergs s$^{-1}$ Hz$^{-1}$, and an unlensed $K = 24.5$, two magnitudes fainter than $K^*$ at $z=3$ \\citep{Shapley:01}. Galaxies at $z \\apg 5$ are faint, and hence only a few high quality spectra exist (Price et al. 2007; Kawai et al. 2006; Dow-Hygelund et al. 2007; Franx et al. 1997). Some of these are at $z > 6$ where $\\tau_{GP}^{eff}$ is expected to be high (from the QSO measurements), but the spectrum of the gamma-ray burst GRB060510B at $z = 4.94$ (Price et al. 2007). shows some similar characteristics to the spectrum of $A1689\\_7.1$.   The IGM transmission measured over a broad redshift bin of the former is $T = 0.18$, similar to the values observed toward QSOs. However, this broad average smears out the structure in $\\tau_{GP}^{eff}$ as a function of redshift.  The GRB060510B spectrum shows evidence of higher than expected GP optical depth extending beyond the region affected by the damped Ly$\\alpha$ absorption at the redshift of the GRB host. It would be interesting to reanalyze this GRB spectrum using smaller redshift bins.    H~I opacities were also measured  inside the proximity zones of QSOs; although significantly ionized, the proximity zones are found to have neutral hydrogen fractions that exceed theoretical expectations.  This indicates that at least some QSOs are also found in regions of gaseous overdensity with large sizes of $15h^{-1}$Mpc \\citep{Guimaraes:07}.  Might $A1689\\_7.1$ still be in the process of accreting much of its mass from its overdense surrounding cosmic structure? Theoretical models predict that H~I gas is funneled into young galaxies via large-scale filamentary structures \\citep{Keres:05, Birnboim:03}. We have presented here observational evidence that H~I column densities are higher than expected near one high-redshift galaxy.  Based on the large physical size implied by the H~I excess, it is unlikely that this gas will accrete onto a single galaxy. Alternatively, post-starburst galaxies are known to drive high-velocity outflows with velocities of $\\apll 2000$ km s$^{-1}$ in low-ionization stages such as Mg~II \\citep{Tremonti:07}.  Given sufficient time, the observed excess of H~I optical depth could be explained by such an outflow; however $A1689\\_7.1$ is too young by a significant factor given the best fit stellar age of 100$\\pm 14$ Myr. On the other hand, QSOs are known to drive outflows with terminal velocities in excess of $10^4$ km s$^{-1}$.  $A1689\\_7.1$ does not show obvious signatures of being an AGN,  but even if there was a low-metallicity outflow undetected in our spectrum, the velocities would still fall short by more than a factor of ten of that required to explain the H~I excess.  As more data become available to study the neutral H~I absorption towards the several recently discovered strongly-lensed LBGs, it may be found to be typical for high-redshift objects, both galaxies {\\it and} QSOs, to be located in regions of space with neutral hydrogen gas fractions significantly larger than that of the pervasive IGM. This re-evaluation of the H~I structure in the densest regions of the universe at these epochs will provide the necessary calibration for modeling the formation and evolution of galaxies in the early universe."
    },
    "0808/0808.2115_arXiv.txt": {
        "abstract": "Models of galactic chemical evolution (CEMs) show that the shape of the stellar initial mass function (IMF) and other assumptions regarding star formation affect the resultant abundance gradients in models of late-type galaxies. Furthermore, intermediate mass (IM) stars undeniably play an important role in the buildup of nitrogen abundances in galaxies. Here I specifically discuss the nitrogen contribution from IM/AGB stars and how it affects the N/O-gradient. For this purpose I have modelled the chemical evolution of a few nearby disc galaxies using different IMFs and star formation prescriptions. It is demonstrated that N/O-gradients may be used to constrain the nitrogen contribution from IM/AGB-stars. ",
        "introduction": "Much progress has been made in modelling nucleosynthesis in stars in general, but the origin of nitrogen still remains somewhat uncertain. In the classical picture, the apparent increase in nitrogen production with metallicity, as inferred from spectroscopy of galactic and extragalactic HII-regions, is a manifestation of nitrogen being either primary or secondary (produced from primordial elements or with heavier elements as \"seeds\"). However, reality often turns out to be more complex than our simplest models of it. The origin of nitrogen is probably no exception.  During the last few years, reasonably accurate stellar nitrogen abundances have become available also for Population II stars (see, e.g., \\cite{Cayrel04, Ecuvillon04, Israelian04, Spite05}). The rather large nitrogen abundances relative to iron seen in halo stars in the Milky Way suggest that a relatively large fraction of the nitrogen should come from Type II supernovae, especially at low metallicity. However, recent stellar evolution models, including nucleosynthesis, all appear to predict significant nitrogen production in intermediate-mass (IM) stars (\\cite{Marigo01, Izzard04, Gavilan05, Karakas07}). The typical production factor would be of the order $\\sim 10^{-3}$ of the initial stellar mass. If IM-stars of all metallicities produce that much nitrogen, it would require that metal-poor galaxies (or regions within a galaxy) with low N/O-ratios have very young stellar populations, in fact, they have to be so young that IM-stars have not yet made any significant contribution to the interstellar nitrogen abundance of these objects, unless the production of oxygen is much greater at low metallicity.  In this work I have considered radial gradients of O/H and N/O in the Milky Way and five other nearby late-type galaxies in order to find additional clues to the origin of nitrogen. The advantages of considering the abundance gradients within galaxies, rather than over-all trends obtained from the global properties of galaxies, are that the same stellar nucleosynthesis prescriptions (yields) must be able to simultaneously reproduce the radial abundance gradients in galaxies with quite different properties (sizes, total masses, gas distributions, etc.) and the very same yields must also be consistent with the fact that galaxies with similar O/H-gradients may have quite different N/O-gradients. Models of chemical evolution (CEMs) may thus be able to provide insight and constraints on stellar yields when several galaxies are modelled and analysed in parallel. ",
        "conclusions": "Employing a standard IMF (\\cite{Scalo86}) with a high-mass cut-off at $m_{\\rm up}=60/70 M_\\odot$ the O/H-gradients are well-reproduced by numerical CEMs. However, the N/O-gradients cannot, in general, be reproduced simultaneously. This inconsistency can be explained in several ways, e.g., \\begin{enumerate} \\item the production of nitrogen in intermediate-mass stars is strongly correlated with metallicity, \\item the IMF varies over the galactic disc (or, effectively, from galaxy to galaxy) such that the N/O-ratio is affected, \\item the infalling gas that forms the disc is not primordial, but significantly enriched with heavier elements (from Pop. III or Pop. II stars?), \\item or the P-method produces systematic errors. \\end{enumerate} I believe that the explanations (a) and (b) are more likely then the other two. The main reasons are: (1) many ingredients of stellar evolution models that affects the nucleosynthetic yields (e.g., mass loss rates and the efficiency of hot-bottom burning) are still highly uncertain, and (2) the turn-over in the IMF at low stellar masses appears correlated with the Jeans mass (\\cite{vanDokkum08}), (3) significant effects from chemically enriched infall is only expected in the outer parts of the discs, where the abundances are low, and (4) the P-method is in very good agreement with the standard electron-temperature method (\\cite{Pilyugin01a, Pilyugin01b}). Furthermore, the stellar yields and detailed properties of the IMF are probably the greatest uncertainties in CEMs in general.  \\begin{figure} \\begin{center} \\includegraphics[width=132mm]{mattsson1_fig4} \\end{center} \\caption{\\label{imf} Variations of the IMF. Changing the slope of a simple power-law IMF ($m_{\\rm up} = 50 M_\\odot, $ $m_{\\rm low}= 0.08 M_\\odot$ are left unchanged) will shift the O/H-gradient up or down (see left panel), but leave $d({\\rm O/H})/d{\\rm R}$ almost unaffected. A Salpeter (1955) IMF ($x = 1.35$) which provides an acceptable fit to the observed N/O-ratios (right panel).} \\end{figure}"
    },
    "0808/0808.4113_arXiv.txt": {
        "abstract": "The red spectral shape of the visible to near infrared reflectance spectrum of the sharply-edged ring-like disk around the young main sequence star HR\\,4796A was recently interpreted as the presence of tholin-like complex organic materials which are seen in the atmosphere and surface of Titan and the surfaces of icy bodies in the solar system. However, we show in this {\\it Letter} that porous grains comprised of common cosmic dust species (amorphous silicate, amorphous carbon, and water ice) also closely reproduce the observed reflectance spectrum, suggesting that the presence of complex organic materials in the $\\hra$ disk is still not definitive. ",
        "introduction": "$\\hra$ is a nearby (distance to the Earth $d\\approx 67\\pm 3\\pc$) young main-sequence (MS) star (age $\\approx 8\\pm 3\\myr$; Stauffer et al.\\ 1995) of spectral type A0V (effective temperature $\\Teff\\approx 9500\\K$) with a large infrared (IR) excess which has recently aroused considerable interest. Imaging observations describe dust scattering in the optical (Debes et al.\\ 2008) and near-IR (Augereau et al.\\ 1999; Schneider et al.\\ 1999) and thermal emission at mid-IR (Jayawardhana et al.\\ 1998; Koerner et al.\\ 1998; Telesco et al.\\ 2000; Wahhaj et al.\\ 2005). They reveal a ring-like disk with maximum at $\\simali$70$\\AU$ distance from the central star and $\\simali$17$\\AU$ width that is sharply truncated at the inner and the outer edge. The structure of the HR\\,4796A disk has important implications for planetesimal evolution (Kenyon et al.\\ 1999). Furthermore, possibly existing planets may generate disk asymmetries through gravitational confinement or perturbation (Wyatt et al.\\ 1999) or form rings with sharp edges (Augereau et al.\\ 1999, Klahr \\& Lin 2000, Th\\'ebault \\& Wu 2008).  Very recently, Debes et al.\\ (2008) measured a visible to near-IR photometric reflectance spectrum of the dust ring around $\\hra$. To fit the observed spectrum (which is characterized by a steep red slope increasing from $\\lambda \\approx 0.5\\mum$ to 1.6$\\mum$ followed by a flattening of the spectrum at $\\lambda >1.6\\mum$), Debes et al.\\ (2008) argued for the presence of tholin-like organic material in the disk around $\\hra$. Tholin, a complex organic material, was detected as a major constituent of the atmosphere and surface of Titan and the surfaces of icy bodies in the solar system. The detection of tholin in the $\\hra$ disk -- if confirmed -- would imply that these potential basic building blocks of life may be common in extra-solar planetary systems as well (see van Dishoeck 2008 for an overview of organic matter in space). Its young age places HR\\,4796A at a transitional stage between massive gaseous protostellar disks around young pre-MS T-Tauri and Herbig Ae/Be stars ($\\simali$1$\\myr$) and evolved and tenuous debris disks around MS ``Vega-type'' stars ($\\simali$100$\\myr$; see Jura et al.\\ 1993, Chen \\& Kamp 2004). Detecting organic matter in this system will provide valuable information about the formation and evolution of planetary systems.  The observed thermal emission of the disk, however, was previously reproduced with a model population of porous grains consisting of the common cosmic dust species amorphous silicate, amorphous carbon, and water ice (Li \\& Lunine 2003a). In this {\\it Letter} we question the existence of tholin as a major dust component in the HR\\,4796A disk and study alternative dust models to reproduce the observed reflectance spectrum. ",
        "conclusions": "\\label{discussion} Debes et al.\\ (2008) calculated the reflectance spectra of ``astronomical silicates'', water ice, hematite Fe$_2$O$_3$ (which is found on the surface of Mars and responsible for its redness), and UV laser ablated olivine (which has been used to explain the spectral reddening of silicate-rich asteroids due to space weathering; Brunetto et al.\\ 2007). But the scattering spectra of these minerals and ice are too neutral at $\\lambda$\\,$\\sim$\\,0.5--1.6$\\mum$ to match the steep red spectral slope of the reflectance spectrum of the $\\hra$ disk. Debes et al.\\ (2008) therefore resorted to organic materials. They found that tholin or its mixture with other dust species (e.g. water ice or olivine) are able to reproduce the observed red spectral slope. This led them to suggest that ``{\\it the presence of organic material is the most plausible explanation for the observations}'', with a cautionary note that ``... {\\it longer wavelength scattered light observations will further constrain the (tholin-based) grain models, particularly around 3.8--4$\\mum$, where a large absorption feature is seen for different grain sizes of tholins. This would help to directly confirm whether Titan tholins are an adequate proxy for the material in orbit around HR\\,4796A.}''  However, as shown in Figure \\ref{fig:kml_refl}, simple porous dust models consisting of dust species (amorphous silicate, amorphous carbon, water ice) which are commonly considered to dominate in the interstellar medium (ISM), envelopes around evolved stars, and dust disks around young stars closely reproduce the observed reflectance spectrum of the $\\hra$ disk. Our model provides at least a viable alternative to the tholin-based models of Debes et al.\\ (2008). While the tholin organic dust model predicts a strong feature around 3.8--4$\\mum$ characteristic of tholin (Debes et al.\\ 2008), the porous dust model presented here predicts a strong band at $\\simali$3.1$\\mum$, attributed to the O--H stretching mode of water ice.  The tholin organics of which the optical constants were adopted by Debes et al.\\ (2008) were made from DC discharge of 90\\% N$_2$ and 10\\% CH$_4$ gas mixture (Khare et al.\\ 1984). They are extremely N-rich and optically very different from amorphous carbon. The dust in the HR\\,4796A disk should be continuously replenished. This is indicated by the low size cutoff $\\amin$ (of a few micrometers) of the dust required to reproduce the reflectance spectrum (see \\S2 and Table \\ref{tab:para}) combined with considerations of dust lifetimes based on the radiation pressure and Poynting-Robertson effects (see Fig.\\,9 of Li \\& Lunine 2003a). The replenishing source would likely arise from collisional cascades of larger bodies like planetesimals, asteroid-like and comet-like bodies. We suggest, with interstellar dust as the building blocks of the parent bodies, it is more reasonable to assume that the dust in the HR\\,4796A disk is composed of amorphous silicate, amorphous carbon, and water ice.\\footnote{% The tholin model may still be viable if the dust in the $\\hra$ disk originates from the surface layers of planetesimals, possibly through excavating collisions (J.H. Debes, private coomunication). In this scenario, the entire planetesimal need not be composed of tholins; they might reside only on the surface where they are created and then be released. Due to significant processing and possible alteration through planetesimal formation, the surface compositions of planetesimals in the $\\hra$ disk may not resemble the ISM composition. Based on what we know from our own Solar System, methane ice and water ice may reside on the surfaces of large planetesimals. These ices are exposed to the stellar UV flux and may ultimately produce tholin-like organic residues. }   The model which best fits both the observed reflectance spectrum and the IR emission requires highly porous dust (with $P=0.73$ which corresponds to $P=0.90$ if ice is sublimated; see \\S2). While it is natural to recognize that cold conglomeration of dust grains in molecular clouds can lead to highly porous dust structures,\\footnote{% A porosity in the range of $0.80\\simlt P\\simlt 0.97$ is expected for dust aggregates formed through coagulation as shown both theoretically (Cameron \\& Schneck 1965; Wada et al.\\ 2008) and experimentally (Blum et al.\\ 2006). } at a first glance, it is harder to accept that comparatively more violent collisions between larger bodies would result in such a morphology in the resulting debris.  To address this concern, we take the interplanetary dust particles (IDPs) as an analog for the dust in debris disks. The anhydrous chondritic IDPs collected in the stratosphere possibly of cometary origin show a highly porous structure (Brownlee 1987). Love et al.\\ (1994) have measured the densities of $\\simali$150 unmelted chondritic IDPs with diameters of $\\simali$5--15$\\mum$, using grain masses determined from an absolute X-ray analysis technique with a transmission electron microscope and grain volumes determined from scanning electron microscope imaging. They found that these particles have an average density of $\\simali$2.0$\\g\\cm^{-3}$, corresponding to a moderate porosity of $\\simali$0.4. More recently, Joswiak  et al.\\ (2007) identified 12 porous cometary IDPs (based on their atmospheric entry velocities) with an average density of $\\simali$1.0$\\g\\cm^{-3}$, corresponding to $P\\approx 0.7$. Much higher porosities ($>$0.9) have been reported for some very fluffy IDPs (e.g. MacKinnon et al.\\ 1987, Rietmeijer 1993), despite that highly porous IDPs are probably too fragile to survive atmospheric entry heating. Low densities are also derived for different groups of meteoroids\\footnote{% By definition, meteoroids are small bodies in the mass range of $\\simali$$10^{-4}$--$10^8\\g$, which orbit the Sun in interplanetary space. The atmospheric trajectories of meteors, i.e. the brightness generated by meteoroids passing through atmosphere, contain information about meteoroid orbits and densities. Meteoroids are accordingly classified into groups with different orbits, structure and composition. The densities given above are valid for between $\\simali$50\\% and 73\\% of cometary meteor observations, depending on the observation method. The same studies show densities of $\\simali$2$\\g\\cm^{-3}$ for the remaining cometary meteoroids, as well as for a significant part of the meteoroids ascribed to asteroids, which is also moderately porous (see Mann 2008). } ascribed to comets: the densities are $\\simali$1.0, 0.75, and 0.27$\\g\\cm^{-3}$, respectively (Ceplecha et al.\\ 1998), corresponding to a porosity of $\\simali$0.7, 0.8, and $>$0.9.    It is therefore reasonable to assume that the dust in the $\\hra$ disk has a high porosity (at least for those originated from cometary bodies). We are not sure about the relative contributions to the $\\hra$ disk from asteroid collisions and cometary activity. Note that this is a long-standing problem even for our own solar system (e.g. see Lisse 2002).  To conclude, we argue that the presence of tholin-like complex organic materials in the $\\hra$ disk is still not conclusive since the observed red spectral shape of the disk can be closely reproduced by models of porous dust comprised of common cosmic dust species. A more thorough study of the scattered light over a range of scattering angles would further constrain the optical properties of the dust."
    },
    "0808/0808.4055_arXiv.txt": {
        "abstract": "Based on 2MASS J and Ks photometry for the open star clusters  NGC 2383, NGC 2384, Pismis 6, Pismis 8 and using color magnitude diagrams with isochrones fit, we found an age of $\\log (\\mbox{age})$ = 8.3 (200 $\\pm$ 6 Myr) for NGC 2383 and $\\log (\\mbox{age})$ = 6.9 (8 $\\pm$ 6 Myr) for NGC 2384. For Pismis 6 and Pismis 8 we adopted a range of $\\log (\\mbox{age})$ = 6 - 7 (1 - 10 Myr). Because they similar ages, Pismis 6 and Pismis 8 may have been formed in the same Giant Molecular Cloud, and we concluded they are a good candidate for a binary system. In the case of NGC 2383 and NGC 2384, because the big age difference found we conclude that  most probably they are born in different environments and as well are not physically connected. ",
        "introduction": "Open clusters are very important objects in the study of stellar evolution because their members are all of very similar ages and chemical composition. This way the effects of other more subtle variables on the properties of stars are much more easily studied than they are for isolated stars. The total number of open clusters known in our Galaxy is over 1600, (see \"New catalog of optically visible open clusters and candidates\"  Dias et al. \\cite{dias1}) of these the only well established double or binary cluster is NGC 869 and NGC 884 (known also as $h + \\chi$ Persei), located at a distance of more than 2 kpc from the Sun. The existence of other possible double clusters has been proposed earlier from Pavloskaya et al. \\cite{pav2}, but not been seriously looked into. Subramaniam et al. \\cite{sub3} examined existing catalogues of open clusters and suggested 18 probable binary open star clusters. The aim of this study is to determine the ages of two probable couples NGC 2383 - NGC 2384 and Pismis 6 - Pismis 8. Our research is based on J and Ks photometry from Two Micron All Sky Survey (2MASS project used two highly-automated 1.3-m telescopes, and provide all-sky photometry in J (1.25 microns), H (1.65 microns), and Ks (2.17 microns) bands). ",
        "conclusions": "Using 2MASS  J and Ks photometry for the open star clusters  NGC 2383, NGC 2384, Pismis 6 and Pismis 8, and fitting CMDs with isochrones based on the Geneva models, we found $\\log (\\mbox{age})$ = 8.3 (200 $\\pm$ 6 Myr) for NGC 2383, $\\log (\\mbox{age})$ = 6.9 (8 $\\pm$ 6 Myr) for NGC 2384, and range of $\\log (\\mbox{age})$ = 6 - 7 (1 - 10 Myr) for Pismis 6 and Pismis 8. Pismis 6 and Pismis 8 have similar age and may have been formed in the same GMC, and  we conclude that they are a good candidate for binary cluster. In contrast NGC 2383 and NGC 2384 have a big age range between, and may be not formed in the same GMC."
    },
    "0808/0808.3443_arXiv.txt": {
        "abstract": "We present the results from arcsecond resolution observations of various line transitions at 1.3 mm toward hypercompact HII region G28.20-0.04N. With the SMA data, we have detected and mapped the transitions in the CH$_{3}$CN, CO, $^{13}$CO, SO$_{2}$, OCS, and CH$_{3}$OH molecular lines as well as the radio recombination line H30$\\alpha$. The observations and analysis indicate a hot core associated with G28.20-0.04N. The outflow and possible rotation are detected in this region. ",
        "introduction": "The scenario of massive star formation remains unclear and has been observationally challenging because of large distances, clustered formation environments, and shorter evolutionary timescales of massive stars. Ultracompact HII (UCHII) regions are considered signposts of massive star formation, but do not represent the earliest stage of massive star forming process (e.g., Churchwell 2002). Hot molecular cores are defined as compact ($\\leq$ 0.1 pc), dense ($\\geq 10^7$ cm$^{-3}$) and warm ($\\geq$ 100 K) molecular cloud cores (Kurtz et al. 2000). The observations from various wavelengths have suggested that hot cores are the sites for massive star formation, and that they represent the early phase of the evolution prior to UCHII regions (Kurtz et al. 2000; Gibb et al. 2000, 2004; Churchwell 2002). So far only twenty or so hot cores from high-resolution observations have been reported in the literature. The observations of hot cores are important for understanding the evolutionary sequence and physical conditions of massive star formation. Hypercompact HII (HCHII) regions have smaller sizes and higher densities compared with UCHII regions, and they probably correspond to a transition phase from hot cores to UCHII regions.  The HCHII region G28.20-0.04N at a distance of 5.7 kpc has been observed at radio and millimeter wavelengths (Fish et al. 2003; Sollins et al. 2005; Sewilo et al. 2004, 2008; Keto et al. 2008). Masers, rotation, inflow, and outflow in G28.20-0.04N have been revealed from observations of molecular and radio recombination lines at centimeter wavelengths (Argon et al. 2000; Menten 1991; Sollins et al. 2005; Sewilo et al. 2008). However, the star formation environment and physical conditions of the hot core in this region are still poorly understood.  Due to the compact and dense nature of hot cores, high angular resolution observations using molecular lines with high critical densities and excitation temperatures at (sub)millimeter wavelengths are crucial to uncovering the physical conditions and kinematics in G28.24-0.04N. In this letter, we present the Submillimeter Array (SMA) \\footnote {The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} observations of molecular lines at 1.3 mm toward the hot core in G28.24-0.04N. ",
        "conclusions": ""
    },
    "0808/0808.2359_arXiv.txt": {
        "abstract": "To probe further the possible nature of the unidentified source IGR J17098-3628, we have carried out a detailed analysis of its long-term time variability as monitored by RXTE/ASM, and of its hard X-ray properties as observed by INTEGRAL.  INTEGRAL has monitored this sky region over years and significantly detected  IGR J17098-3628 only when the source was in this dubbed active state. In particular, at $\\ge$ 20 keV, IBIS/ISGRI caught an outburst in March 2005, lasting for $\\sim$5 days with detection significance  of 73$\\sigma$ (20-40 keV) and with the emission at $< $200 keV. The ASM observations reveal that the soft X-ray lightcurve shows a similar outburst to that detected by INTEGRAL, however the peak of the soft X-ray lightcurve either lags, or is preceded by, the hard X-ray ($>$20 keV) outburst by $\\sim$2 days. This resembles the behavior of X-ray novae like XN 1124-683, hence it further suggests  a LMXB nature for IGR J17098-3628. While the quality of the ASM data prevents us from drawing any definite conclusions, these discoveries are important clues that, coupled with future observations, will help to resolve the as yet unknown nature of IGR J17098-3628. ",
        "introduction": "INTEGRAL has detected roughly 500  sources at energies $\\ge$ 20 keV \\citep{boda}. Among them, X-ray binary systems were identified for 32\\% of the times, and another 26\\% remained unidentified. As one of the unidentified sources, IGR J17098-3628 was discovered \\citep{source} at the end of March 2005 by INTEGRAL/IBIS during the private Open Program observation dedicated to the deep view to the Galactic center. The source was located at R.A.(2000) $=17^{h}09^{m}48^{s} $, Dec. $=-36\\degree 28\\arcmin 12\\arcsec $, with an error box of $2\\arcmin$ at 90$\\%$ confidence. The source got a peak flux at $\\sim$ 60 and 95 mCrab for 18-45 and 45-80 keV, respectively. The   20-60 keV flux evolved from $\\sim50$ mCrab at 2005 March 26 00:00, to $\\sim9$ mCrabs at April 3 23:30 (UTC) \\citep{453}. The preliminary spectral analysis \\citep{spe} showed the source spectrum became quite soft along this evolution. Following INTEGRAL, IGR J17098-3628 was observed by RXTE on March 29 of 2005, and was detected at 80 mCrab in the 3--20 keV band, with a hard power-law tail of a spectral index $\\sim$ 2.5. Then the source was assumed a BHC and X-ray Nova (Grebenev et al. 2007). The column density was found to be less than 1$\\times$10$^{22}$ atoms cm$^2$.  Later on, Swift observed this source on May 1 of 2005, and refined its location to R.A.\\ (2000) $=17^{\\rm h}09^{\\rm m}45.9^{\\rm s} $, Dec.\\ $=-36\\degree 27\\arcmin 57\\arcsec $, with 90$\\%$ confidence, and an uncertain radius of 5 arcseconds \\citep{kenn}. This made the search possible for the potential counterparts in optical and  infrared bands. From the 2MASS ALL-SKY Catalog a possible counterpart was found, J17094612-3627573 \\citep{kong}. However, the most recent optical and infrared observations, made with the 6.5m Magellan-Baade telescope, revealed that 2MASS J17094612-3627573 is in fact composed off several sources \\citep{stee}. Furthermore, Very Large Array (VLA) observations of IGR J17098-3628 were made on March 31, April 5, and May 4, 2005 at 4.86 GHz \\citep{radio}, which  led to the discovery of a radio transient at 0.8 arcseconds from the Swift XRT position, and hence regarded as the possible radio counterpart. Finally on July 9, Swift/XRT observed IGR J17098-3628 again and found that the spectrum was  fitted by an absorbed disk blackbody model. The column density obtained was (0.89$\\pm$0.02$)\\times$10$^{22}$ atoms cm$^{-2}$ \\citep{kenn2}.   In this paper, we report the analysis of all available observations on IGR J17098-3628  carried out by IBIS/ISGRI and JEMX onboard INTEGRAL, and the All Sky Monitor (ASM) onboard RXTE. This allows us to trace the source behavior in X-rays back to 1996 with ASM and back to 2002 with INTEGRAL. We put the outburst in 2005 in this context. ",
        "conclusions": "IGR J17098-3628 may have stepped into an active phase in the beginning of 2005. INTEGRAL detected a days-long outburst extending up to energies of $\\sim$200 keV. A similar outburst is detected in the soft X-ray (1.5-12 keV) band, however the onset of the outburst lags the hard X-ray outburst by $\\sim$2 days, in a fashion reminiscent of X-ray novae like XN 1124-683 \\citep{nova}. The spectral analysis of the initial stages of the outburst detected by INTEGRAL supports the association with X-ray novae (Grebenev et al. 2007), the cool disk temperature and small inner radius of the accretion disk both suggest that IGR J17098-3628 is a newly discovered black hole system.  Current theoretical models of the formation of X-ray novae include the mass transfer instability (MTI) model \\citep{hkl86}, the disk thermal instability  (DTI) model (Cannizzo et al. 1982; Faulkner et al. 1983; Meyer \\& Meyer-Hofmeister 1984; Huang \\& Wheeler 1989; Mineshige \\& Wheeler 1989; Ichikawa et al. 1994) Unfortunately, the only way to discriminate between these models relies on detailed differences in the observed soft energy spectrum of the source and so the lack of any detailed X-ray spectra of the IGR J17098-3628 outburst in this energy band  prevents us from drawing any conclusions on the X-ray nova mechanism for IGR J17098-3628. That the three episodes of the hard X-ray outburst observed by INTEGRAL have to be described by different models might suggest rather strong time evolution. i.e.,   the spectral index changed from $\\sim$ 1.7 when rising to roughly 2.6 when decaying suggesting the cooling of the source.    Although the time variability of X-ray novae is not well understood theoretically, observations of a variety of X-ray novae not only show that the hard X-ray flux often rises and peaks earlier than the outburst in soft X-rays, but can also show precursor events at hard X-rays several days prior to the main outburst (Chen et al. 1997). While it is not possible to associate the outburst in IGR J17098-3628 with any particular X-ray nova mechanism spectrally, it is clear that the typical accretion scenario for the low mass X-ray binaries holding a black hole is not capable of reproducing the observed lag between the hard and soft X-ray outbursts. In this scenario, the soft X-rays are produced by the accretion disk, while the hard X-rays are produced by comptonization of soft seed photons in the inner part of the accretion disk \\citep{disalvo,stella}. Clearly, large changes in the hard X-ray flux can, at best, be coeval with large changes in the soft X-ray seed photons and cannot precede variations of the soft X-ray photons that the process of comptonization relies on to produce the hard X-ray tail. One potential mechanism that may reproduce the observed lag is for an instability in the disk to trigger the formation of a temporary jet, which produces hard X-rays via inverse comptonization, in advance of any change in the flux of soft X-rays from the disk.  A more straightforward scenario in understanding this different time behaviour  between the hard and the soft X-rays may lie in the tie-up with  the structure of the accretion flow. The precursor event can be explained as the transition from the \"hard\" state to \"soft\" state in the BHC system. The accretion time in the inner parts of the disk is much smaller than the duration of X-ray nova's outburst therefore changes in the spectral state of these sources are connected with general changes in structure of the accretion flow rather than in changes in number of soft seed photons for Comptonization. Observations show that all BH systems have the hard Comptonized spectrum at low accretion rates and the soft DBB(disc blackbody) spectrum (with a possible weak hard tail) at high accretion rates. This may be connected with approaching of the inner radius of a cold accretion disk to the radius of marginally stable orbit when the accretion rate increases. The observed lag between outbursts in hard X-rays and soft X-ray may be directly connected with this observed dependence: the accretion rate is small at the initial stage of the outburst thus the spectrum is hard and the outburst is observed in hard X-rays, but later the accretion rate rises and the spectrum becomes soft thus the outburst is observed in soft X-rays.    Although no clear statement can be made regarding their nature, the discovery  the lag/precursor event revealed by the concurrent RXTE ASM and INTEGRAL observations in 2005, is important piece of the IGR J17098-3628 puzzle that will help to resolve the as yet unknown nature of this perplexing source."
    },
    "0808/0808.0809_arXiv.txt": {
        "abstract": "People usually smile when astrophysicists assert that we are {\\it sons of the stars}, but human life confirms this sentence: about 65\\% of the mass of our body is made up of oxygen, carbon occurs in all organic life and is the basis of organic chemistry, nitrogen is an essential part of amino acids and nucleic acids, calcium is a major component of our bones. Moreover, phosphorus plays a major role in biological molecules such as DNA and RNA (where the chemical codes of life is written) and our blood carries oxygen to tissues by means of the hemoglobin (an iron pigment of red blood cells). All these elements have been created in stars. I just list some examples related to human body, but also common element such as aluminum, nickel, gold, silver and lead come from a pristine generation of stars. The abundances in the Solar System are in fact due to the mixing of material ejected from stars that polluted the Universe in different epochs before the Sun formation, occurred about 5 billion years ago, after the gravitational contraction of the proto-solar cloud. 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. They are responsible for the nucleosynthesis of the main component of the cosmic s-elements. ",
        "introduction": "The majority of the isotopes heavier than iron (A$\\ge$56) are synthesized by neutron capture processes. The observed heavy elements distribution shows the presence of two main components, correlated to different nucleosynthetic processes: the s (slow) process  and the r (rapid) process (on the basis of the definitions given by \\cite{b2hf} in their pioneering work). The r process requires high neutron densities, and it is believed to occur during explosive phases of stellar evolution (Novae, SuperNovae and/or X-rays binaries). The s process is characterized by a slow neutron capture with respect to the corresponding $\\beta$ decay: stable isotopes capture neutrons, while the radioactive ones decay ($\\beta^-$ or $\\beta^+$) or capture a free electron. These isotopes are mainly created in the Thermally Pulsing AGB (TP-AGB) phase of low mass stars (1 $\\leq$ M/M$_{\\odot}$ $<$ 4), where freshly synthesized elements are carried out to the surface by means of a recurrent mechanism called Third Dredge Up (TDU) (see \\cite{stra06} and references therein). \\begin{figure}[tb] \\centering \\includegraphics[width=8.5cm]{mcore.ps} \\caption{Evolution of the positions in mass of the inner border of (top to bottom): convective envelope, H-burning shell and most energetic mesh of the He-burning shell, during the Thermally Pulsing AGB phase of a model with initial mass M=2M$_\\odot$ and solar metallicity..} \\label{mcore} \\end{figure} In this phase the stellar structure consists of a partially degenerate carbon-oxygen core, an He shell separated from an H shell by the He-intershell region and by a convective envelope (see Fig.\\ref{mcore}). The energy required to supply the surface irradiation is mainly provided by the H burning shell, located just below the inner border of the convective envelope. This situation is recurrently interrupted by the growing up of thermonuclear runaways, driven by violent He-burning ignitions. As a consequence of a Thermal Pulse (TP), the region between the two shells (He-intershell) becomes unstable against convection (for a short period), the external layers expand and, later on, the H shell burning temporarily dies down. In the He-intershell, He is partially converted into carbon. During the AGB phase, main neutron sources are the $^{13}$C($\\alpha$,n)$^{16}$O reaction, active in radiative layers during the interpulse period \\cite{stra95}, and the $^{22}$Ne($\\alpha$,n)$^{25}$Mg reaction, marginally activated within the convective shell originated by the TP. In order to obtain a sufficient amount of $^{13}$C for the activation of the s-process, a diffusion of protons from the H-rich envelope into the $^{12}$C-rich radiative zone is needed: the diffused proton are captured from the abundant carbon via the $^{12}$C(p,$\\gamma$)$^{13}$N($\\beta^-$)$^{13}$C nuclear chain, leading to the formation of a tiny $^{13}$C-pocket. ",
        "conclusions": "The present work demonstrates that, nowadays, the computational power allows the coupling between a stellar evolutionary code and a full nuclear network. A different treatment of the internal border of the convective envelope with respect to a bare Schwarzschild criterion allows the formation of the so-called $^{13}$C pocket. For the first time in the literature, it has been possible to directly compare theoretical models and observational data both from a physical (luminosities, surface temperatures) and a chemical (light elements and heavy elements abundances) point of view . Moreover, we furnish a uniform  set of yields (from hydrogen to lead) at different metallicities. The importance of the adopted mass loss rate and of molecules contribution to opacity has been pointed out. While the first problem is still under analysis, the second one has been solved by using opacity tables with enhanced carbon and nitrogen abundances. Finally, we have addressed the importance of very low metallicity AGB models, by describing first preliminary (and promising) results."
    },
    "0808/0808.0038_arXiv.txt": {
        "abstract": "We present a sample of 16 radio galaxies, each of which is characterized by a wide,   elongated emission gap with   fairly sharp and straight edges between the two radio lobes.    This particular subset of the ``superdisk'' radio galaxies is chosen because of a highly asymmetric location of the host elliptical galaxy relative to the gap's central axis. In addition to posing a considerable challenge to the existing models, such a morphology also means that the two jets traverse highly unequal distances through the superdisk material. One thus has a possibility to directly investigate if the marked asymmetry between the two jets' interaction with the (much denser) ambient medium, during their propagation, has a significant import for the brightness of the hotspot forming near each jet's extremity. We also propose a new explanation for the formation of superdisks through the merger of a smaller elliptical galaxy with the massive host, in which the gas attached to the infalling galaxy deposits its angular momentum into the host's circumgalactic gas, thereby causing it to flatten into a fat pancake, or superdisk.   The asymmetric location of the host galaxy can be assisted by the kick imparted to it during the merger. We also suggest a physical link between  these radio galaxies and those with {\\bf X}-shaped and {\\bf Z}-symmetric radio lobes, commonly believed to arise from mergers of two galactic nuclei, each harboring a supermassive black hole. ",
        "introduction": "A few years ago we noted  that a fraction of extragalactic double radio sources exhibit radio morphologies and other properties that led us to infer that they required gaseous ``superdisks\" extending tens of kiloparsecs in both diameter and thickness \\citep[hereafter GKW00]{gkw00}. The key morphological property shared by these radio galaxies is the sharp, quasi-linear edges of the radio lobes on the sides facing the central elliptical host galaxy \\citep{gkw96,gkna}, as seen in at least a dozen radio galaxies at low to moderate redshifts.  We showed that such superdisks (SDs) can provide consistent, alternative explanations for some key correlations found among the parameters of extragalactic radio sources, such as the Laing-Garrington effect \\citep[e.g.,][]{garr88,lain} and the correlated radio-optical asymmetries  \\citep[e.g.,][]{mcca}.  Furthermore, in high-$z$ RGs the strong tendency for redshifted component of the diffuse, quiescent Ly$\\alpha$ emission to be on the side of the brighter hot spot has been explained by postulating that the redshifted Ly$\\alpha$ emission originates from the region of the radio lobe on the near side of the nucleus, and is thus subject to much less dilution due to the intervening dust \\citep[e.g.,][]{hump,gkw05}. As discussed in GKW00, these Ly$\\alpha$ RGs provide optical evidence for high-$z$ SDs. Because the sharp-edged morphology would only be noticed when the radio jets are oriented close to the plane of the sky, SDs are subject to a strong negative selection effect, and GKW00 argued that even though the observed cases are few, the phenomenon may not be so uncommon. In this paper we present an enlarged sample of radio galaxies (RGs) exhibiting SDs and discuss, in particular, their properties related to structural asymmetry.  The sharp emission gaps seen in double radio sources are most commonly attributed to a blocking of the backflowing radio plasma in the lobes by a denser thermal plasma  associated with the parent elliptical galaxy \\citep[e.g.,][]{leahw,wiitg,blac}. Such backflows are to be expected and simulations indicate that they will indeed be diverted by galactic ISM \\citep[e.g.,][]{wiitn}. This approach was an alternative to the original proposal by \\citet{sche} according to which the central emission gap arose from the pinching of the inner parts of the lobes by the higher gas pressure of the host galaxy.  Both of these mechanisms, however, seem incompatible with cases where the edges of the gaps are long and straight; they seem particularly incapable of explaining the highly asymmetric SDs where the host elliptical is found almost at one edge of the radio emission gap.  A third possible explanation for large gaps has recently been put forward by \\citet{gerg}; they suggest that SDs could be carved out as two galaxy cores containing supermassive black holes (SMBHs) merge and the associated pair of jets undergoes a rapid precession during the later stages of the merger.  While this novel process might indeed create a gap, it too may have difficulty in explaining the   sharp inner edges of the lobes, particularly those of the highly asymmetric SDs we are exploring here.  Initially we proposed that the SD is primarily made of the interstellar medium bound to the RG itself, perhaps originating from the cool gas belonging to gas-rich disk galaxies previously captured by the giant elliptical host of the powerful RG \\citep[e.g.,][]{stat}. We argued that the tidal stretching and heating of that gas during the capture was sufficient to produce very large fat pancakes \\citep{gkna,gkw00}. Since in some extreme cases (e.g., 0114$-$476; Table 1) SDs were found to have widths running into several hundred kiloparsecs, an alternative possibility was also considered in GKW00, according to which at least some SDs trace the gaseous filaments of the ``cosmic web\".  In a recent paper, we have attempted to explain the occurrence of SDs in high-$z$ RGs which have yet to acquire a significant circumgalactic medium (CGM) or reside in an intracluster medium (ICM) \\citep[ICM;][]{gkwj}. We argued that the material forming the SD in such RGs likely arises from the nuclear wind which is believed to be associated with the AGN activity \\citep[e.g.,][]{soke}. We investigated the conditions under which this wind material would be squeezed between the two radio lobes into a pancake shaped SD. It was shown that for a wide range of reasonable wind and jet parameters those jets launched within a few tens of Myr subsequent to the initiation of wind production could quickly catch up to the wind-blown bubble and then the lobes growing outside the bubble could indeed squeeze the bubble in such a way as to produce SDs.  Such a squeezing, however, is unlikely in $z < 1$ RGs, which are usually surrounded by a significant amount of X-ray emitting CGM with a pressure at least that of the radio lobes \\citep[e.g.,][]{cros}. Therefore, another mechanism seems to be required at least for the SDs in low $z$ RGs.  In \\S 2 we tabulate a significant number of additional RGs with highly asymmetric SDs.  We also briefly discuss some interesting morphological properties of these asymmetric SDs. We then propose in \\S 3 a new mechanism based on a merger of an elliptical galaxy, also rich in hot gas, with the massive host elliptical. In this scenario the orbital angular momentum of the captured galaxy's gas will be gradually transferred to the CGM gas of the host galaxy, thus transforming its quasi-spherical CGM halo into a fat pancake. This scenario ties in well with our explanation for the {\\bf Z}-symmetries seen in some {\\bf X}-shaped RGs (XRGs) \\citep[hereafter GKBW03]{gkbw}. Conclusions are summarized in \\S 4. ",
        "conclusions": "We have highlighted a particularly interesting and intriguing class of RGs which not only shows an SD type morphology but also the host galaxy is seen close to the edge of the SD and is thus grossly offset from the midplane of the SD. This morphological asymmetry poses a serious difficulty for existing SD models, particularly when the superdisk is found in RGs at $z \\lesssim 1$, where a significant amount of hot gas surrounding the host is usually present. We have shown that the two radio lobes extend significantly more symmetrically about the SD midplane than they extend relative to  the host galaxy.  We have proposed a scenario whereby gaseous SDs acquire their form through the injection of angular momentum from the gas belonging to a galaxy whose eventual merger with the massive elliptical host triggers the jets responsible for the RG. The gradual transfer of the angular momentum into the hot CGM of the host elliptical flattens its CGM into a fat-pancake shaped superdisk, which can account for the strip-like emission gaps observed between the radio lobes of such RGs. In this picture, the grossly off-centered location of the host (relative to the central plane of the SD) is relatively straightforward to understand in terms of the motion of the host galaxy during the period of jet activity, partly assisted by a linear momentum kick imparted by the merged galaxy. A corollary to this picture is that the SD material (which appears to be ``docking the tails'' at least in these RGs \\citep[cf.,][]{jenk} is not gravitationally bound to the radio-loud elliptical. Interestingly, this scenario also provides a possible physical link between the SDs  and the XRGs that are widely believed to manifest mergers of two galaxies containing SMBHs.  The merger is likely to launch a pair of jets and if the host had already been recently active, a double-double radio structure can arise, as seen in three RGs in the present sample (0035$+$130, 0114$-$476, and 1155$+$256). It is interesting that  for between 10 and 13 of these 16 RGs, the brighter of the two outer hotspots is associated with the radio lobe adjoining the host galaxy.  This indicates a diminishing of a jet's radio emitting potential following a long passage through the denser thermal material associated with the SD, as compared to a jet that mainly traverses the (light) relativistic plasma within a radio lobe. Perhaps a striking manifestation of this is seen in the RGs 0106$+$729 and 1939$+$605  where the radio jet appears to turn dissipative after propagating through the SD.  No such dissipation is observed,  at least in the present sample, for the jets that propagate mainly through the radio lobe. This calls for a follow-up study once a significantly larger sample of highly asymmetric superdisks becomes available.  We intend to investigate the realm of applicability of this model through simulations of the merger of two  massive ellipticals with substantial circumgalactic gas."
    },
    "0808/0808.2384_arXiv.txt": {
        "abstract": "We report the discovery of ripple-like X-ray surface brightness oscillations in the core of the Centaurus cluster of galaxies, found with 200~ks of \\emph{Chandra} observations.  The features are between 3 to 5 per cent variations in surface brightness with a wavelength of around 9~kpc. If, as has been conjectured for the Perseus cluster, these are sound waves generated by the repetitive inflation of central radio bubbles, they represent around $5\\times 10^{42}\\ergps$ of spherical sound-wave power at a radius of 30~kpc. The period of the waves would be $10^7$~yr. If their power is dissipated in the core of the cluster, it would balance much of the radiative cooling by X-ray emission, which is around $1.3 \\times 10^{43} \\ergps$ within the inner 30~kpc. The power of the sound waves would be a factor of four smaller that the heating power of the central radio bubbles, which means that energy is converted into sound waves efficiently. ",
        "introduction": "Ripples in the X-ray surface brightness of galaxy clusters were first discovered in the Perseus cluster of galaxies \\citep{FabianPer03}. As the surface brightness is roughly proportional to the density of the intracluster medium (ICM) squared, these are density ripples.  \\cite{FabianPer03} explained these variations as sound waves generated by the inflation of the bubbles of relativistic plasma in the core of the cluster by the central active nucleus. Providing that these sound waves can be dissipated as they travel \\citep{FabianReynolds05}, they have the ability to transport energy from the central nucleus to the core of the cluster. If they carry significant quantities of energy, they may provide the distributed heating required \\citep{Voigt04} to prevent large quantities of the ICM from cooling \\citep{PetersonFabian06,McNamaraNulsen07}.  Subsequent theoretical work \\citep{Ruszkowski04,Sijacki06} have shown that sound waves may be able to help heat the cores of clusters in a gentle distributed fashion.  A deep 900~ks observation of Perseus showed the ripples in exquisite detail enabling their amplitude to be measured \\citep{FabianPer06}. Subsequently, we used the amplitude of the ripples to calculate the amount of energy propagated in them if they are sound waves \\citep{SandersPer07}. They would carry enough energy in Perseus to combat a significant fraction of the cooling rate, and appear to decline in power with radius.  The Perseus observations also reveal thick spherical shells around each inner bubble at a higher gas density and pressure than their surroundings. These high pressure shells have a sharp outer edge where the density abruptly jumps by about 30 per cent, and is interpreted as a weak shock. We have assumed that such high pressure regions propagate outward to become the ripples. A problem with the weak shock interpretation is that there is no temperature jump coincident with the density one. The observed limit on any jump in temperature is inconsistent with an adiabatic shock.  The region has been studied and discussed in detail by \\cite{Graham08Per}. It is noted there that the Northern optical filaments of cold gas stop at the shock front, so cold gas is likely being mixed with the hot gas there due to turbulence generated at that front. This could explain a reduction in temperature of the post shock gas.  Sound waves or ripples have not been detected in other systems. However, they may still be an important contribution to the heating processes in clusters as they could be difficult to detect. They were not found until 200~ks of \\emph{Chandra} time of the Perseus cluster was obtained. Perseus is the X-ray brightest galaxy cluster by a factor of 1.7 (in the 2-10~keV band; \\citealt{Edge90}). Clearly the detectability of ripples will vary with cluster surface brightness and observation time. It will also vary with ripple amplitude and wavelength and with cluster properties. These effects have been examined numerically by \\cite{Graham08}, finding that it is difficult to observe these features with current instruments in clusters other than Perseus. The detectability of ripples, however, depends on unknown factors, such as the wavelength and amplitude of the ripples. These factors can be estimated with a large degree of uncertainty from the properties of the cluster, assuming that they are responsible for heating the cluster core and have a wavelength similar to the bubble size.  Here we report on some ripples which we find in a 200~ks \\emph{Chandra} observation of the Centaurus cluster of galaxies. These data were examined previously by \\cite{Fabian05} and \\cite{SandersEnrich06}, examining the thermal distribution of gas in the cluster core and its metallicity.  The Centaurus cluster of galaxies is a nearby ($z=0.0104$; \\citealt{LuceyCurrieDickens86a}) X-ray bright galaxy cluster ($L_{X,2-10\\keV} = 2.9 \\times 10^{43} \\ergps$; \\citealt{Edge90}). It contains a low radiative efficiency nucleus \\citep{Taylor06} and distinct bubbles of relativistic plasma \\citep{Taylor02} displacing the intracluster medium, as seen by \\emph{Chandra} \\citep{SandersCent02}.  We assume $H_0 = 70 \\kmpspMpc$, which translates into scale of 213~pc per arcsec. ",
        "conclusions": "We detect surface brightness deviations in a \\emph{Chandra} image of the Centaurus cluster of galaxies. These features are between 3 and 5 per cent in size. Assuming that they are sound waves, the wave power is approximately $5\\times 10^{42}\\ergps$ in the inner 30~kpc, with a period of $10^7$ yr and a wavelength of 9~kpc. The power is close to the X-ray luminosity of this central region. It is therefore energetically possible that sound waves are the mechanism by which the power of the central black hole is transmitted quasi-isotropically through the cluster. Provided that the power in the waves is gradually dissipated as heat, they provide the means by which radiative cooling is balanced by heating in the inner cluster core."
    },
    "0808/0808.2667_arXiv.txt": {
        "abstract": "We present a systematic study of line widths in the [\\ion{O}{3}]$\\lambda$5007 and H$\\alpha$ lines for a sample of 86 planetary nebulae in the Milky Way bulge based upon spectroscopy obtained at the \\facility{Observatorio Astron\\'omico Nacional in the Sierra San Pedro M\\'artir (OAN-SPM)} using the Manchester Echelle Spectrograph.  The planetary nebulae were selected with the intention of simulating samples of bright extragalactic planetary nebulae.  We separate the planetary nebulae into two samples containing cooler and hotter central stars, defined by the absence or presence, respectively, of the \\ion{He}{2}\\,$\\lambda$6560 line in the H$\\alpha$ spectra.  This division separates samples of younger and more evolved planetary nebulae.  The sample of planetary nebulae with hotter central stars has systematically larger line widths, larger radii, lower electron densities, and lower H$\\beta$ luminosities.  The distributions of these parameters in the two samples all differ at significance levels exceeding 99\\%.  These differences are all in agreement with the expectations from hydrodynamical models, but for the first time confirmed for a homogeneous and statistically significant sample of galactic planetary nebulae.  We interpret these differences as evidence for the acceleration of the nebular shells during the early evolution of these intrinsically bright planetary nebulae.  As is the case for planetary nebulae in the Magellanic Clouds, the acceleration of the nebular shells appears to be the direct result of the evolution of the central stars. ",
        "introduction": "The interacting stellar winds model of \\citet{kwoketal1978} provides a general framework for understanding planetary nebulae.  Within this framework, a variety of theoretical and numerical studies of the hydrodynamics have been undertaken to understand and predict the kinematic properties of planetary nebulae.  Our modern view of the kinematics of planetary nebulae was established in the early 1990's \\citep[e.g.,][]{kahnwest1985, breitschwerdtkahn1990, kahnbreitschwerdt1990, mellema1994}.  These studies showed that the AGB envelope is accelerated in two phases.  The first phase is a result of the shock wave initiated by the ionization front that sweeps through the AGB envelope as the central star's temperature increases.  Then, as the central star's wind energy increases, it drives a pressure-driven central bubble that accelerates the AGB  envelope through ram pressure.  More recent numerical work that attempts to include more realistic AGB evolution confirms these basic results \\citep[e.g.,][]{villaveretal2002, perinottoetal2004}.  Eventually, the central star's wind ceases and the inner part of the AGB envelope backfills towards the central star while the outer part maintains its momentum-driven evolution \\citep[e.g.,][]{garciaseguraetal2006}.  These models have been used extensively to interpret observational studies of the kinematics of many individual planetary nebulae.  Unfortunately, to date,  systematic, homogeneous studies of planetary nebula populations to compare with these theoretical efforts are scarce.  \\citet{heap1993} found a correlation between nebular expansion velocity and the terminal velocity of the central star wind.  More recently, \\citet{medinaetal2006} claimed a correlation between expansion velocity and age indicators such as density and central star temperature.  However, both of these studies suffer from heterogeneity, at least regarding object selection.  Even regarding planetary nebulae with Wolf-Rayet (WR) central stars, the precise role that the winds from the central stars play in the kinematics of the nebular shells is not entirely clear \\citep{gesickietal2006, medinaetal2006}.  The studies of the kinematics of planetary nebulae in the Magellanic Clouds are perhaps the most systematic \\citep{dopitaetal1985, dopitaetal1988}.  These studies find correlations between expansion velocity, excitation class, and nebular density, suggesting that the properties of the central star dominate the nebular evolution.  As they point out, however, their results are not general since they have been found only for a particular population of (bright) planetary nebulae.  Therefore, there remains a need for a coherent population study in a different galactic environment.  Here, we present a study of the kinematics for a large sample of planetary nebulae in the bulge of the Milky Way (Bulge).  These objects were selected in a homogeneous way, with the hope of simulating populations of bright extragalactic planetary nebulae in bulge-like systems (bulges of spiral galaxies, dwarf spheroidals, and elliptical galaxies).  The observations were all made with the same instrument and methodology.  Likewise, the data reduction and analysis is homogeneous and entirely independent of model parameters.  Details are given in Section 2.  We find that the expansion properties do indeed depend upon the evolutionary stage of the central star, with planetary nebulae hosting hotter central stars having larger expansion velocities (Section 3). We also find that other parameters that should depend upon age, the nebular size, density, and H$\\beta$ luminosity, also differ significantly between the two groups (Section 4).  We compare these properties with theoretical models and find generally good agreement (Section 5).  We present our conclusions in Section 6. ",
        "conclusions": "We have obtained kinematic data for a large sample of planetary nebulae in the Milky Way bulge, selected with the goal of simulating samples of bright extragalactic planetary nebulae in bulge-like environments.  For most of the sample, our criteria also  included a reddening-corrected H$\\beta$ flux exceeding $\\log F(\\mathrm H\\beta)>-12.0$\\,dex and [\\ion{O}{3}]$\\lambda 5007/\\mathrm H\\beta > 6$, in addition to the usual criterion of projected proximity to the Galactic center.  We measure line widths for H$\\alpha$ and [\\ion{O}{3}]$\\lambda 5007$.  Our sample is the largest that has been selected and observed in a homogeneous way.  We have also analyzed it in a completely model-independent manner.  For half of the sample, the \\ion{He}{2}$\\lambda 6560$ line is observed in the H$\\alpha$ spectra.  Since this line appears only for hotter central stars, it allows us to divide our sample according to the evolutionary stage of the central star, with the more evolved objects found in the subsample containing the hotter central stars.  The kinematics of the two samples are significantly different, with the sample containing the hotter central stars expanding faster.  We also compare the diameters, electron densities (from [\\ion{S}{2}]$\\lambda\\lambda 6716,6731$), and H$\\beta$ luminosities for the two subsamples.  In all cases, there are statistically significant differences.  The planetary nebulae hosting the hotter central stars are larger, less dense, and less luminous than their counterparts with cooler central stars.  All of these differences exceed a statistical significance of 99\\%.  Also, all of these differences are compatible with the results of hydrodynamical models \\citep[e.g.,][]{mellema1994, villaveretal2002, perinottoetal2004}.  Our primary conclusion is that we have clearly observed the acceleration of the nebular shells in planetary nebulae in the bulge of the Milky Way and that this occurs during the early evolution of their central stars.  Our findings, based upon a large, homogeneous sample, constitute the first unequivocal evidence that the nebular kinematics depend upon the evolutionary state of the central star for planetary nebulae in the Milky Way and are far clearer and more convincing than any previous results.  Hence, there is now very clear evidence that this dependence occurs in at least two environments: the Magellanic Clouds and the bulge of the Milky Way.  Given the differences in the stellar populations in these two environments, it is likely that the dependence of the nebular kinematics on the evolutionary state of central star is more general and that it occurs in all environments."
    },
    "0808/0808.3985_arXiv.txt": {
        "abstract": "We present {\\sl Suzaku} spectra of X-ray emission in the fields just off the LMC~X--3 sight line. \\ovii, \\oviii, and \\neix\\ % emission lines are clearly detected, suggesting the presence of an optically thin thermal plasma with an average temperature of 2.4 $\\times 10^6$ K. This temperature is significantly higher than that inferred from existing X-ray absorption line data obtained with \\chandra grating observations of \\xs, strongly suggesting that the gas is not isothermal. We then jointly analyze these data to characterize the spatial and temperature distributions of the gas. Assuming a vertical exponential Galactic disk model, we estimate the gas temperature and density at the Galactic plane and their scale heights as $3.6(2.9, 4.7)\\times10^6$ K and $1.4(0.3, 3.4)\\times10^{-3}~{\\rm cm^{-3}}$ and  $1.4(0.2, 5.2)$ kpc and $2.8(1.0, 6.4)$ kpc, respectively. This characterization can account for all the \\ovi\\ line absorption, as observed in a \\fuse\\ spectrum of LMC~X--3, but only predicts less than one tenth of the \\ovi\\ line emission intensity typically detected at high Galactic latitudes. The bulk of the \\ovi\\ emission most likely arises at interfaces between cool and hot gases. ",
        "introduction": "\\label{sec:intro}  The Galactic diffuse hot gas at temperatures $\\sim10^6$ K can be effectively probed via its emission and absorption features in X-ray and far-ultraviolet (UV) wavelength bands. The previous X-ray emission investigations were largely based on the broadband X-ray background data, e.g., {\\sl ROSAT ALL Sky Survey} (RASS; Snowden \\etal 1997). A high spectral resolution X-ray calorimeter abroad a sounding rocket, though providing little spatial resolution, clearly detected the \\ovii\\ and \\oviii\\ emission lines, confirming that much of the soft X-ray background (SXB) emission is thermal in origin \\citep{mcc02}. Recently, several groups have attempted to study the background emission with X-ray CCDs aboard {\\sl Suzaku X-ray Observatory}, which, compared to those on {\\sl XMM-Newton} and {\\sl Chandra X-ray observatories}, have a significantly improved spectral resolution and a low instrument background (e.g., Smith \\etal 2007; Henley \\& Shelton 2007). These later X-ray observatories, however, also carry the high resolution grating instruments that allow for the detection of the X-ray absorption lines (e.g., from \\ovii, \\oviii, and \\neix\\ K$\\alpha$ transitions) by diffuse hot gas in and around the Galaxy. Indeed, such absorption lines are detected in grating spectra of nearly all Galactic and extragalactic sources as long as the spectral signal-to-noise ratio is high enough (e.g., Futamoto \\etal 2004; Yao \\& Wang 2005; Fang \\etal 2006; Bregman \\& Llyoid-Davis 2007). These X-ray absorption lines trace gas over a broad temperature range of $\\sim 10^{5.5}-10^{6.5}$~K.  Gas at temperatures $\\lsim 10^{5.5}$~K may be traced more sensitively in far-UV, via the detection of the \\ovi\\ lines at $\\lambda\\lambda$1031.96 and 1037.62. Extensive observations of the lines in absorption have been carried out with {\\sl Copernicus} and {\\sl Far Ultraviolet Spectroscopy Explorer} ({\\sl FUSE}; Jenkins 1978; Savage \\etal 2003; Bowen \\etal 2008). In addition, the \\ovi\\ lines have also been detected in emission with {\\sl FUSE}, although the sky is only sampled at various Galactic latitudes (e.g., Shelton \\etal 2001; Otte \\& Dixon 2006). The intensity of the $\\lambda$1031.96 line, for example, ranges from 1800 to 9100 LU (line unit; 1 ${\\rm photon~cm^{-2}~s^{-1}~sr^{-1}}$). However, the interpretation of the \\ovi\\ line(s) alone is not straight forward, because in the collisional ionization equilibrium (CIE) state, the \\ovi\\ population sharply peaks at the intermediate temperature $\\sim 10^{5.5}$~K where gas cools very efficiently \\citep{sut93}. Thus such \\ovi-bearing gas is expected to be rare and may preferentially reside at interfaces between cool gas clouds and thermally more stable hot gas. But this latter hot gas, which should be more abundant, distributed more widely, and effectively traced by the X-ray \\ovii\\ line,  could contribute significantly to the \\ovi\\ line as well. Currently, little is known about the relative contributions from these two origins to the \\ovi, either in absorption or in emission.  Clearly, a combined analysis of the X-ray and far-UV lines, in both emission and absorption, will be the most beneficial. While an absorption line is proportional to the total column density of the gas integrated along a line of sight, an emission line depends on the emission measure (EM) of the gas. Furthermore, for a gas in the CIE state, both the ionic column density and the EM depend on the gas temperature, but in different manners (Fig.~\\ref{fig:absVSemi}). Therefore a joint analysis of multiple emission/absorption lines will enable us to constrain not only the temperature and its distribution but also the size and density of the intervening gas (e.g., Shull \\& Slavin 1994). When transitions from multiple elements (e.g., O and Ne) are detected, their relative abundances can also be estimated (Yao \\& Wang 2006). Such a joint analysis has been tentatively applied to several sources (e.g., Yao \\& Wang 2007; Shelton \\etal 2007); but none of these sources has high resolution emission and absorption data available in both X-ray and far-UV wavelength bands.  \\begin{figure} \\plotone{f1.eps} \\caption{Ionization fraction (a) and the emissivity (b) of oxygen ions as a function of temperature for a gas in the collisional ionization equilibrium state \\citep{sut93}. The emissivity is a summation of the doublet transitions at 1031.96 \\AA\\ and 1037.62 \\AA\\ (\\ovi), triplet transitions at 22.10 \\AA, 21.80 \\AA, and 21.60 \\AA\\ (\\ovii), and K$\\alpha$ plus K$\\beta$ transitions at 18.97 \\AA\\ and 16.0 \\AA\\ (\\oviii). The emissivity of \\ovi\\ is scaled down by a factor of 1000 for demonstration purpose. \\label{fig:absVSemi} } \\end{figure}    In this paper, we report our investigation of the hot gas along the sight line toward LMC~X--3. \\citet{wang05} have reported the detection of the hot gas associated with our Galaxy, based on X-ray and far-UV absorption line spectra from \\chandra\\ and \\fuse\\ observations. Here we present the {\\sl Suzaku} CCD emission spectra of the X-ray background in two fields adjacent to the LMC~X--3 sight line. This unique combination of the X-ray and far-UV spectral data toward essentially the same part of the sky further allows us to constrain the spatial, thermal, and chemical properties of the hot gas.  This paper is organized as follows. We present the {\\sl Suzaku} observations and the data calibration, as well as a brief description of the existing {\\sl Chandra} and {\\sl FUSE} data in \\S~\\ref{sec:obs}. In \\S~\\ref{sec:results}, we first analyze the X-ray emission ({\\S~\\ref{sec:emission}) and absorption ({\\S~\\ref{sec:abs}) data separately, and then build a slab-like hot gas model to jointly analyze these X-ray data (\\S~\\ref{sec:model}). In \\S~\\ref{sec:fuse}, we compare the model predicted \\ovi\\ absorptions with the \\fuse\\ detections, and then use the observed \\ovi\\ absorption line to further constrain our model. We discuss the implications of our results in \\S~\\ref{sec:dis} and summarize our results and conclusions in \\S~\\ref{sec:sum}.  Throughout the paper, we assume the hot emitting/absorbing gas to be optically thin (see \\S~\\ref{sec:scattering} for further discussion) and in the CIE state. We adopt the solar abundances from \\citet{and89} and quote parameter errors at 90\\% confidence levels for a single varying parameter unless otherwise noted. We also refer the hot gas on scales of several kpc as the Galactic disk, in contrast to the Galactic halo on scales of $>10$ kpc (see \\S~\\ref{sec:location} for further discussion). Our spectral analysis uses the software package XSPEC (version 11.3.2). ",
        "conclusions": "\\label{sec:dis}  We have presented the high quality {\\sl Suzaku} emission observations of the diffuse hot gas toward two off-fields of the LMC~X--3 sight line. In particular, the \\ovii\\ and \\oviii\\ emission lines are clearly resolved. Modeling these emission spectra yields a gas temperature that is about two times higher than that inferred from the high resolution absorption data. We find that our non-isothermal thick Galactic gaseous disk model can account for this discrepancy as well as the far-UV \\ovi\\ absorption data. A joint fit to these data indicates that both the X-ray absorption and emission and the far-UV absorption are consistent with being produced from hot gas in a region of several kpc around the Galactic plane. In the following, we first discuss how our results are potentially affected by caveats in our data analysis, and then discuss the implications of our results on the origin and cooling of the \\ovii- and \\ovi-bearing gases.  \\subsection{Uncertainty of the LHB and the SWCX} \\label{sec:LHB}  The biggest uncertainty in modeling the SXB emission is the estimate of the contributions from the LHB and SWCX. Our understanding for both components is still very poor because it is very hard to decompose their contributions. The SWCX emission could in principle be estimated by tracing the solar wind flux and the neutral atoms in the Earth's atmosphere and in the local ISM (e.g., Lallement \\etal 2004). However, such an estimation is model dependent. For instance, for the line of sight with Galactic coordinates $(l, b)=(278.65^\\circ, -45.30^\\circ)$ observed on 2003 March 3, the \\ovii\\ emission is estimated to be 0.83 and 3.7 LU by \\citet{kou07} and by \\citet{hen08}, respectively. Recent comparisons of the observed X-ray emission with models for the SWCX suggest that the foreground emission in the shadowing experiments can be mainly attributed to the SWCX and that the LHB emission is negligible (Koutroumpa \\etal 2007; Rocks \\etal 2007; but also see Shelton 2008). The total \\ovii\\ intensity of the SXB toward a high Galactic latitude sight line ($l, b = 272^\\circ, -58^\\circ$) is 2.7(2.0, 3.4) LU (1$\\sigma$ range), obtained from a recent {\\sl Suzaku} observation \\citep{mit06}. If emission from both the LHB and the SWCX is isotropic and independent of Galactic latitude, this value can be regarded as the upper limit of the combined contribution of these two components.  Our results are not qualitatively affected by above discussed uncertainties. To be more quantitative in assessing the effect, we allow EM of the LHB plus SWCX component to vary from 0 to 0.0113 ${\\rm cm^{-6}~pc}$ (\\S~\\ref{sec:emission}), representing two extreme cases. The low boundary corresponds {\\it zero} emission of the LHB and the SWCX components, whereas the high boundary corresponds to the 3$\\sigma$ upper limit of the observed \\ovii\\ intensity (i.e., 4.8 LU) toward the high Galactic latitude sight line. We then re-perform our spectral analysis described in \\S\\S~\\ref{sec:emission} and \\ref{sec:fuse}. We find that new constrained gas temperature in modeling the emission data alone under the isothermal assumption is still about two times higher than that constrained from the absorption data (refer to the first and the fifth rows in Table~\\ref{tab:results}), which still necessitates the non-isothermal model developed in \\S~\\ref{sec:model}. The new results constrained in the joint analysis are also consistent with the old values, except for with a slightly broader uncertain range (refer to the forth and the sixth rows in Table~\\ref{tab:results}).    \\subsection{Uncertainty of the extragalactic power-law component} \\label{sec:cxb}  In modeling the emission spectra, we used a simple PL function to approximate the extragalactic contribution (\\S~\\ref{sec:emission}). Recent {\\sl Chandra} and {\\sl XMM-Newton} deep surveys for the cosmic X-ray background have resolved $\\gsim80\\%$ of the background emission at 1-8 keV into extragalactic discrete sources (e.g., Hickox \\& Makevitch 2006). Optically bright sources tend to be spectrally hard while faint sources tend to be spectrally soft (e.g., Mushotzky \\etal 2000). To examine the goodness of our approximation, we replace the simple PL with a broken PL and fix the break energy at 1 keV to model the extragalactic component in our joint analysis described in \\S~\\ref{sec:model}. We obtain the photon indices as 1.7(1.3, 2.0) and 1.4(1.3, 1.5) below and above 1 keV, respectively, and find that the parameters of the Galactic hot gas are barely affected (the seventh row of Table~\\ref{tab:results}).  \\subsection{Effect of resonant scattering of the \\ovii\\ line emission} \\label{sec:scattering}  In this work, we have assumed the emitting/absorbing gas to be optically thin, i.e., we have neglected the resonant scattering effect of the hot gas. To qualitatively assess the validity of our assumption, we assume the hot gas density to be uniform for easy of analysis. The observed \\ovii\\ emission line consists of unresolved triplet transitions, resonance, forbidden, and intercombination lines at 21.60 \\AA, 22.10 \\AA, and 21.80 \\AA, respectively. Since the oscillation strength is essentially zero for the latter two transitions, we consider the ``scattering'' only for the resonant line. The absorbed resonant photons should be re-emitted isotropically; half of these re-emitted photons are expected to favor the observing direction. Taking all these into account, we obtain the observed line intensity as \\begin{equation} I = (1 - f) I_0 + 0.5fI_0 \\left( 1 + \\frac { 1 - e^{-\\tau} }{\\tau} \\right ), \\end{equation} where $\\tau$ is the (max) absorption optical depth at center wavelength of the resonant transition, and $I_0$ is the ``intrinsic'' (without the scattering) line intensity. Here, $f$ is the resonant transition fraction of the \\ovii\\ triple, which is a function of gas temperature. Taking the (max) temperature at the Galactic plane $T_0\\sim3\\times10^6$ K (Table~\\ref{tab:results}), which favors more to the resonant transition and gives $f=0.65$, we get $I \\gsim 0.675I_0$ for any value of $\\tau$. This indicates that the observed intensity should not be different from the intrinsic value by a factor larger than 1.5. In reality, the ``background'' gas cloud also scatters the photons emitted by the ``foreground'' cloud to the observing direction. If this effect is further considered, the difference should largely disappear. Therefore we conclude that the resonance scattering should not significantly affect our results, although a more detailed modeling needs to be performed to account for a more realistic geometry of the hot gas distribution.  \\subsection{Origin of the \\ovii-bearing gas in general} \\label{sec:location}  Our analysis indicates that the gas responsible for the observed X-ray absorption/emission along the LMC~X--3 sight line has a pathlength about a few kpc and is consistent with the Galactic disk in origin. This result has important implications for understanding the structure and origin of hot gas around our Galaxy.  As described in \\S~\\ref{sec:intro}, high ionization X-ray absorption lines with zero velocity shift have been observed along many extragalactic sight lines, and the \\ovii\\ emission lines are also detected at various high Galactic latitudes. Several scenarios have been proposed for the origin of the emitting/absorbing gas, including the intergalactic medium (IGM) of the Local Group, the large-scale Galactic halo, and the thick Galactic gaseous disk (see Yao \\& Wang [2007] for a review). These scenarios cannot be distinguished kinematically because of the limited spectral resolution of current X-ray instruments. So currently the most direct information about the location of the gas comes from differential analysis of the X-ray absorption lines toward sources at different distances. \\citet{wang05} find that the \\ovii\\ absorption along the sight line toward LMC~X--3, at a distance of $\\sim50$ kpc, is comparable to those observed toward AGN sight lines. Therefore, if the LMC~X--3 sight line is representative, one can then conclude that the bulk of the X-ray absorption around the Galaxy is within $\\sim50$ kpc. A similar conclusion has been reached based on the estimate of the total baryonic matter contained in the \\ovii-bearing gas, and on the angular distribution of the \\ovii\\ absorption \\citep{fang06, bre07}. A comparison of a Galactic sight line (4U~1957+11) with two extragalactic sight lines toward LMC~X--3 and Mrk~421 further indicates that there is no significant \\ovii\\ absorption in the extended Galactic halo beyond a few kpc \\citep{yao08}, which is consistent with the conclusion drawn in this work.  The \\ovii\\ emission from the hot gas beyond LMC~X--3 is not expected to be important either. \\citet{yao08} obtained a 95\\% upper limit to the \\ovii\\ absorption beyond LMC~X--3 as $N{\\rm _{OVII}<3.7\\times10^{15}~cm^{-2}}$. Assuming this much of \\ovii\\ uniformly distributed in a region with an extension scale of $l$, we calculate its emission as $I_{\\rm OVII}<0.025/(A_{\\rm O0.1}l_{\\rm 100kpc})$ LU for the entire temperature range from $10^5$ to $10^7$ K, where $A_{\\rm O0.1}$ is the oxygen abundance in unit of 10\\% solar value and $l_{\\rm 100kpc}$ is in unit of 100 kpc. In contrast, the inferred the \\ovii\\ intensity of the Galactic disk diffuse gas is 5.0 LU (\\S~\\ref{sec:emission}; Table~\\ref{tab:line}).  We have shown that the X-ray emission and absorption as well as the far-UV \\ovi\\ absorption along the sight line toward \\xs\\ can be explained by the presence of a thick Galactic hot gaseous disk. A similar explanation holds for the sight line toward Mrk~421, for which another joint X-ray absorption/emission and UV-absorption has been carried out (the X-ray emission is, however, mostly based on the RASS broad-band measurements; Yao \\& Wang 2007). The characteristic scale height of the disk is also consistent with the \\neix\\ column density distribution in the Galaxy \\citep{yao05}. The gas in the disk is shown to have a normal metal abundances \\citep{wang05, yao06} and is clearly due to the heating by mechanical energy feedback from stars in the Galaxy (e.g., Ferri\\`ere [1998]). Such an interpretation is also supported by the X-ray emission observations of nearby disk galaxies like our own. The diffuse X-ray emission has been routinely observed in many late-type normal star-forming spirals and is confined to a region around galaxies with vertical extent no more than a few kpc (e.g., Wang \\etal 2001, 2003; T\\\"ullmann \\etal 2006a, 2006b). Both the intensity and the spatial extent appear to be scaled with galactic star formation rate (e.g., T\\\"ullmann \\etal 2006b). Therefore, we conclude that the X-ray emitting/absorbing and \\ovi-absorbing gas, as studied in this paper, arises primarily from the stellar feedback in the Galactic disk.  \\subsection{Location of the \\ovi-bearing gas} \\label{sec:o6location}  Before discussing the \\ovi-bearing gas location, let us look at the \\ovi\\ emission properties first. In \\S~\\ref{sec:fuse} we have shown that the observed \\ovi\\ absorption can be predicted from the X-ray-data-constrained non-isothermal gas model and have then used the \\ovi\\ $\\lambda$1031.96 absorption line to further constrain our spectral model. With the fitted model parameters, we can estimate the intrinsic \\ovi\\ emission intensity of the diffuse hot gas toward LMC~X--3, and then compare it with observations. The best-fit model gives 186 LU in total for the \\ovi\\ doublet $\\lambda\\lambda$1032 and 1038. Taking all the 90\\% boundaries that favor more \\ovi\\ emission, we obtain a firm upper limit of 1612 LU.  There is no direct measurement of the \\ovi\\ emission within 1$^\\circ$ of the LMC~X--3 sight line. Within $\\sim5^\\circ$, some observations show measurable \\ovi~$\\lambda1032$ emission as $2.1-5.7 \\times10^3$ LU, while other observations only yield upper limit even with longer exposures \\citep{ott06, dix06, dix08}. Since the \\ovi\\ emission could vary by a factor of $>2.3$ at an angular scale as small as $25'$ \\citep{dix06}, it is nearly impossible to reliably estimate the \\ovi\\ emission toward LMC~X--3 sight line based on the few existing nearby samples. However, if there were a measurable \\ovi\\ emission toward the LMC~X--3 sight line (but see below), it must be in order of several thousand LU, as typically observed at high galactic latitudes, which is significantly larger than that predicted in the \\ovii-bearing gas and therefore unlikely arises from the diffuse hot gas.  Where is the \\ovi-bearing gas then? To answer this question, it is important to compare the spatial distributions of the \\ovi\\ absorption and emission. In this work, we find that most of the \\ovi\\ absorption can arise from the diffuse gas along a vertical scale of $\\sim$3 kpc around the Galactic plane. Figure~\\ref{fig:OVI_dis} shows the density and column density distributions of \\ovi\\ and \\ovii\\ as a function of the vertical distance, predicted from our constrained non-isothermal disk model. While compared to \\ovii, \\ovi\\ is distributed in a relative narrow region, which is mainly due to the sensitive dependence of \\ovi\\ population on temperature (Fig.~\\ref{fig:absVSemi}), about 80\\% of the \\ovi\\ column still spreads over as large as 1 (from 2.7 to 3.7) kpc. Observationally, \\ovi\\ absorption has been detected toward essentially all sight lines, and the characteristic scale height of the \\ovi\\ column densities is similar to that constrained in this work (e.g., Savage \\etal 2003). It has been suggested that \\ovi\\ could largely arise at interfaces between hot and cool interstellar media. However, for Mrk~421 sight line, \\citet{sav05} counted the number of cool gas velocity components and found that the interfaces can only account for at most 50\\% of \\ovi\\ absorption, leaving a bulk of the \\ovi\\ to other origins. In contrast, the \\ovi\\ emission is measured toward only $\\sim30\\%$ of the surveyed sky \\citep{ott06, dix08}. These results clearly indicate that the ``commonly'' observed \\ovi-absorbing gas in kpc scale is not primarily responsible for the detected \\ovi\\ emission. We propose that the low temperature regime of the kpc-scale diffuse gas could contain a significant or even dominant fraction of the observed \\ovi\\ absorption, as indicated in this work, and that the conductive interfaces between hot and cool gases, though could be another reservoir of \\ovi\\ ion, is mainly responsible for the observed \\ovi\\ emission. This picture not only explains the spatial distribution of the \\ovi\\ absorption, but also naturally explains the sight line to sight line \\ovi\\ variation in both absorption and emission. For LMC~X--3 sight line in particular, since all the observed \\ovi\\ can be attributed to the diffuse gas, the \\ovi\\ emission is not expected to be strong.  \\begin{figure} \\plotone{f7.eps} \\caption{Spatial distribution of density (a) and column density (b) of \\ovi\\ and \\ovii, predicted with our constrained non-isothermal disk model. \\label{fig:OVI_dis} } \\end{figure}  \\subsection{Radiative energy loss rate of the hot gas} \\label{sec:lossrate}  If hot gas toward the LMC~X--3 sight line represents reasonable well the hot ISM of the Galactic disk as a whole, we can then use our non-isothermal gas model to estimate the total radiative energy loss rate of the gas, and compare it to the expected energy input from SN explosions. It has long been observed that in the low $L{\\rm_x}/L{\\rm_B}$ ($L_{\\rm B}$ is the blue luminosity) early-type galaxies, X-ray radiative energy loss rate of the diffuse hot gas is far less than the expected SN energy input (e.g., Canizares \\etal 1987; Brown \\& Bregman 2000). However, all the previous studies were based on imaging observations in which decomposing different emission components is a complicated task (e.g., Li \\etal 2007) and could make the inferred properties of the ``true'' diffuse hot gas very uncertain. The LMC~X--3 sight line is so far the only sight line along which the high resolution X-ray emission/absorption and far-UV absorption measurements with a minimal confusion are available, which allow us to tightly constrain the thermal and chemical properties as well as the global spatial distribution of the Galactic diffuse hot gas, as presented in this work. It is thus worthwhile reexamining the radiative power of the Galactic hot gas based on these characterizations.  The energy loss rate of the hot gas can be expressed as \\begin{equation} \\label{equ:lossrate} \\frac{d^2U}{dsdt} = 2\\times\\int_0^\\infty  1.2n^2\\Lambda(T)~dz, \\end{equation} where the factor 2 accounts for the two sides of the Galactic disk. With our model (Eq.~\\ref{equ:expHT}), we can rewrite the Eq.~\\ref{equ:lossrate} as \\begin{equation} \\label{equ:rate} \\frac{d^2U}{dtds} = 2\\times\\int^{T_0}_{T_{min}} 1.2n_0^2 \\left(\\frac{T}{T_0}\\right)^{2\\gamma-1}\\Lambda(T) (-h_T)~d\\left(\\frac{T}{T_0}\\right). \\end{equation} Assuming the solar abundances, plugging in the best-fit $n_0$, $T_0$, $\\gamma$, and $h_T$ values (\\S~\\ref{sec:model}), and taking the emissivity $\\Lambda(T)$ from the atomic database ATOMDB\\footnote{http://cxc.harvard.edu/atomdb}, we obtain the local surface emissivity of the hot gas as 3.1 and 16.1 $\\times10^{36}~{\\rm ergs~s^{-1}~kpc^{-2}}$ in 0.1-10 and 0.01-10 keV bands, respectively. Including the additional \\ovi\\ emission of $\\sim$3000 LU from the interface component, the total radiative energy loss rate is $\\sim3\\times10^{37}~{\\rm ergs~s^{-1}~kpc^{-2}}$.  The radiative cooling of the Galactic disk hot gas only accounts for less than 10 per cent of the expected energy input of the stellar feedback. At Sun's galactocentric radius, the SN rate is 19 and 2.6 ${\\rm Myr^{-1}~kpc^{-2}}$ for type II and type Ia, respectively \\citep{fer98}. If on average each type II SN progenitor releases $2\\times10^{50}$ ergs of energy before it explodes and each SN explosion releases $10^{51}$ ergs \\citep{fer98, lei92}, the total energy input is then $8\\times10^{38}~{\\rm ergs~s^{-1}~kpc^{-2}}$, which is about 20 times higher than our estimated gas cooling rate. The ``missing'' energy could be either emitted in other wavelength bands (e.g., infrared) or have been consumed in driving other Galactic activities (e.g., galaxy-size out flows).  Obviously, such an estimation is model dependent. \\citet{she07} recently concluded that $\\sim70\\%$ of the SN energy input could be radiated away by the Galactic hot gas. In their model, they required the observed \\ovi\\ emission to be co-spatial with the \\ovi\\ absorption and assumed both the \\ovi-bearing and the hotter \\ovii-bearing gas to be isobaric. This requirement results in that most of the \\ovi-bearing gas is confined within a region of $<100$ pc in size, which favors to produce about ten times more thermal emission in the energy range of 0.01-0.1 keV than our model. In contrast, we constrain our model with X-ray \\ovii\\ and \\oviii\\ emission and absorption, and find that the most of the \\ovi\\ absorption could arise from the kpc-scale diffuse gas. In our model, because of large extent of the diffuse gas (therefore a low density), its emitting power is at most as much as that of the observed/expected \\ovi\\ emission lines, which presumably arise from the interfaces between the hot and cool interstellar media (\\S~\\ref{sec:o6location}). The scale size of the \\ovii- and \\ovi-bearing gas derived in this work and utilized in our estimation, is consistent with that obtained in the recent \\ovi\\ absorption surveys \\citep{sav03, bow08}.  \\subsection{Evidence for oxygen and iron depletion?} \\label{sec:depletion}  In our analysis, we find that, while the relative abundance of Fe/O is about the solar value, the Ne/O is about 2 times higher (Table~\\ref{tab:results}). This overabundance of neon, if true, is consistent with the interpretation of about 50 per cent of oxygen and iron being depleted into dust grains. However, because this overabundance is mainly derived from the emission data in which decomposing various components is still very uncertain, the depletion interpretation therefore may not be unique. For instance, \\citet{hen08} suggest that the SWCX could also produce apparent non-solar abundance ratio. Furthermore, our understanding of the chemical abundances in the solar system is still poor. For example, the oxygen value has recently been revised downward by $\\sim35\\%$ \\citep{asp05}, but this revision is still under debate \\citep{ant06}. The abundance pattern in the ISM is also suggested to be slightly different from that in the solar system \\citep{wil00}. So it is more useful to list the number density ratio of Ne/O derived in this work, which is 0.25(0.19, 0.33). We prefer to defer the interpretation of the apparent overabundance of neon until a better understanding of both the SWCX and the solar chemical abundances is reached."
    },
    "0808/0808.1980_arXiv.txt": {
        "abstract": "We present a new method for constructing equilibrium phase models for stellar systems, which we call the iterative method. It relies on constrained, or guided evolution, so that the equilibrium solution has a number of desired parameters and/or constraints. This method is very powerful, to a large extent due to its simplicity. It can be used for mass distributions with an arbitrary geometry and a large variety of kinematical constraints. We present several examples illustrating it. Applications of this method include the creation of initial conditions for $N$-body simulations and the modelling of galaxies from their photometric and kinematic observations. ",
        "introduction": "In astronomy there are at least two problems where equilibrium phase models of stellar systems need to be constructed. The first one is the construction of phase models for real galaxies from observational data, i.e. the modelling of observational data. The second problem is the construction of initial conditions for $N$-body simulations of stellar systems. It is obvious that these two problems are tightly connected, and yet they have, so far, been solved by different methods. The Schwarzschild method \\citep{S79} and its modifications is often used for modelling of observational data \\citep[e.g][]{H00, vdB06, T07, vdB08, deL08}, but has almost never been used so far to produce initial conditions for simulations. For $N$-body initial conditions, a wide variety of methods has been used, based on the Jeans theorem \\citep[e.g. ][]{Zangthesis, AS86, KD95, WD05, M07}, or on Jeans' equations \\citep[e.g. ][]{H93}. In the case of multi-component systems, e.g. disc galaxies with a bulge and a halo, the components are built separately and then either simply superposed, or the potential of the one is adiabatically grown in the other \\citep[e.g. ][]{Barnes88, M07, A07} before superposition.  For real galaxies the phase space density is generally unknown, but we do have some information about it. For example, we know more or less accurately a distribution of mass for the visible components (notwithstanding uncertainties due to the mass to light ratio) and we often have some constraints on the velocity distribution. It is also reasonable to assume that the galaxy is in an equilibrium state. So in general, the problem of constructing a model in phase space is equivalent to constructing an equilibrium phase model with a given mass distribution and, in many cases, given kinematic constraints. In the case of modelling observational data (first of the two above mentioned problems) the kinematic parameters are the observed velocities integrated along the line of sight. In the case of initial conditions for $N$-body simulations, a wide variety of kinematic parameters is possible, depending on the problem the simulation addresses.  We have developed a new method for constructing equilibrium phase models with a given mass distribution and with given kinematic parameters, which we call the iterative method. It can be applied to systems with arbitrary geometry, so that the requested mass distribution can be arbitrary. The idea and a first implementation was presented in \\citet{RS06}. In \\citet{RO08} we improved it, and applied it to construct an $N$-body model of the stellar disk of our Galaxy for two realistic mass models of the Milky Way. Here we present a final version of this method, fully allowing kinematical constraints. In the previous articles we had concentrated on constructing equilibrium phase models with a given mass distribution, so that kinematic parameters were either not considered or only in terms of auxiliary parameters, such as the total angular momentum \\citep{RS06, RO08}. This, however, limited the applicability of our method, both for initial conditions and for modelling real galaxies. Indeed, initial kinematics play a crucial role in determining the evolution of $N$-body systems, while observational constraints more often than not include kinematics. In this paper we give equal attention to the mass distribution and kinematical constraints, so that the iterative method can now be used for a number of interesting applications. In principle, in our method, both the kinematic constraints and the mass distribution can be arbitrary. But the part of our algorithm that concerns the kinematic constraints is not universal, contrary to the part that handles the mass distribution, but is tailored to the specific constraint. Here we consider several types of constraints. Once these are understood, it is rather easy to extend the algorithm for every new type of kinematic parameter (see below).  The power of the iterative method stems from its simplicity. The iterative method is based on a simple and, in a way, obvious idea, which is implemented in a simple algorithm. The purpose of this article is to fully describe this method. We first introduce the basic concept in Section~\\ref{s_imethod}, where we also explain the different modules of the algorithm and the way they should be applied. In section~\\ref{s_models} we illustrate the use of the method with three examples, namely a triaxial system, a multi-component model of a disk galaxy (including live disk, bulge and halo components) and a disk constructed with given line-of-sight kinematic. We briefly conclude in section~\\ref{s_conc}. ",
        "conclusions": "\\label{s_conc} We presented a new method for constructing equilibrium phase models for stellar systems --- the iterative method. The aim of this method is to construct equilibrium $N$-body models with given parameters, or constraints. More specifically, these are a given mass distribution and, if desired, given kinematic properties, parameters, or constraints. Our method is straightforward both conceptually and in its implementation. We believe that it is this simplicity that makes this method so powerful. It simply relies on a constrained, or guided evolution. We let the system reach equilibrium via a dynamical evolution in a number of successive steps. In between two such steps we make sure that the parameters are set to their desired value and/or that the constraints are fulfilled. This means that the evolution is guided towards an equilibrium with the desired parameters and/or constraints. Setting a mass distribution is of course obligatory, but kinematical constraints are not. If we wish to include them, we have the choice of a large number of possibilities, such as setting the radial profile(s) of one, or more moments of the velocity distribution. In this article we described only a few types of kinematic constraints: the profile of radial velocity dispersion, the profile of velocity anisotropy, a condition of velocity isotropy and line-of-sight kinematics. Procedures for further types of kinematic constraints can be easily found following similar techniques. Furthermore, our implementation of the iterative method can be directly applied to systems with arbitrary geometry, i.e. the given mass distribution can be arbitrary and need not have any symmetries. Thus our method can be used in many different applications.  We used our iterative method to construct several models. The first one is a triaxial system. The second one is a multi-component model of a disk galaxy consisting of a stellar disk with a given radial velocity dispersion profile, a non-spherical bulge and a halo with a given anisotropy profile. We also constructed two disk models with given line-of-sight kinematic. Using self-consistent $N$-body simulations, we made sure that the models we constructed are indeed very close to equilibrium (see figs.~\\ref{fig_3d},~\\ref{fig_disk},~\\ref{fig_bulge}, ~\\ref{fig_halo},~\\ref{fig_svlos_disk} and~\\ref{fig_mvlos_disk}).  The iterative method has a number of further applications. It can of course be used for constructing equilibrium initial conditions for $N$-body modelling of stellar systems. For instance, the iterative method allows one to investigate bar formation in galaxies with a halo having different kinematics. Also the iterative method can be applied for constructing phase models of real galaxies. For example, we can model observational data by constructing phase models with given line-of-sight kinematics, as shown in section~\\ref{s_losdisks}. This paves the way for studies of e.g. the distribution of dark matter in ellipticals, or obtaining phase space models of observed disk galaxies. A further interesting application is the study of the properties of several equilibrium models for a given mass distribution, as for example triaxial systems.  The software necessary for the implementation of this method should be thought of in a very modular way, e.g. with different units for the various kinematical constraints, and is very straightforward to write. Nevertheless, we will make our own software publicly available as soon as this paper is accepted, at the address http://www.astro.spbu.ru/staff/seger/soft/. This package will contain also step-by-step examples for constructing models by using the iterative method, including the models described in this article. Our software uses the $N$-body code gyrfalcON \\citep{D00, D02} and the NEMO package (http://astro.udm.edu/nemo; \\citealt{T95})."
    },
    "0808/0808.1949_arXiv.txt": {
        "abstract": "Stellar jets are normally constituted by chains of knots with some periodicity in their spatial distribution, corresponding to a variability of order of several years in the ejection from the protostar/disk system. A widely accepted theory for the presence of knots is related to the generation of internal working surfaces due to variations in the jet ejection velocity. In this paper we study the effect of variations in the inner disk-wind radius on the jet ejection velocity. We show that a small variation in the inner disk-wind radius produce a variation in the jet velocity large enough to generate the observed knots. We also show that the variation in the inner radius may be related to a variation of the stellar magnetic field. ",
        "introduction": "Stellar jets are observed in the form of chains of emitting nebulae called Herbig-Haro (HH) objects. HH objects are generally thought to be shock-heated density condensations traveling along the outflows formed by star-disk systems during the process of star formation \\citep[e.g.][]{rei01}.  Since their discovery several scenarios have been suggested for the origin of \\emph{steady} outflows from young stellar objects. In the stellar wind model, material is accelerated by thermal pressure gradients \\citep[e.g.][]{can80}. Magnetohydrodynamics (MHD) models rely on the magnetocentrifugal launching mechanism \\citep{bla82}. For the X-wind scenario the jet is magnetically driven from the so-called ``X-annulus'' where the young star's magnetosphere interacts with the disk \\citep{shu00}. In the disk wind scenario the jet is launched from an extended region of the disk surface \\citep[e.g.][]{fer97}. The analytical models mentioned above have focused mainly on the steady-state aspect of the ejection phenomena.  \\emph{Unsteady} periodic ejections with timescales of the order of several rotation periods of the inner disk radius have been obtained by numerical simulations \\citep[e.g.][]{ouy97,goo99,mat02}. Observations indicate a significantly longer timescale is associated with the appearance of knots in stellar outflows. Nearly all observed jets present small scale knots up to 0.1 pc from the central source, with a spacing between the knots corresponding to a timescale of $\\approx$ 1-20 yr (e.g. HH30 - \\citealt{bur96}; HH111 - \\citealt{rei92}; RW-Aur - \\citealt{lop03}). More fragmented knots are observed on a typical timescale of $\\sim 10^{2-3}$ yr at larger distances from the source. While the long term variation of jets can be explained by variations in the accretion rates (e.g. during FU-Orionis phases), the possible origin of small scale knots is still unclear.  A first possibility, suggested by similarities with extra-galactic jets, is that the knots may be formed by hydrodynamics Kelvin-Helmholtz instabilities \\citep{mic98}, MHD Kelvin-Helmholtz reflective pinch mode instabilities \\citep[]{cer99}, or by current-driven instabilities \\citep[]{fra00}. Numerical simulations by these authors have shown that the shocks generated by plasma instabilities are weaker than those seen in observations once radiative cooling is taken into account. Moreover, some jets \\citep[e.g. HH212 -][]{zin98} show a remarkable symmetry on both sides of the central star-disk system. This may be difficult to explain assuming that the instabilities are triggered by small scale perturbations. Additionally, and more importantly, optical images in many cases show that the compact knots have a bow-shock structure (e.g. HH111 - \\citealt{rei92}) that is difficult to obtain supposing that the knots are formed by instabilities.  A widely accepted theory for the presence of knots in stellar jets is related to the generation of internal working surfaces due to supersonic variations of the ejection velocity at the base of the jet \\citep{rag90}. Indeed, numerical simulations, using a velocity variations of about 10-20\\% of the average velocity, have been able to reproduce the morphology and the emission property of the knots in HH objects \\citep[e.g.][]{esq07}.  In this paper we show that velocity variations which lead to the creation of knots similar to the observed ones may be generated by variations of the inner disk-wind radius. Furthermore, we suggest that these variations may be related to changes in the stellar magnetic field.  This paper is organized as follows. In Section \\ref{model} we determine the effect of a variation in the inner disk radius on the jet velocity. In Section \\ref{variable} we suggest that the variation in the inner radius could be connected to a change of the stellar magnetic field. In Section \\ref{discussion} we discuss the approximation used and the limitations of our model. Conclusions are given in Section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this paper we explored the effect of a stellar magnetic field variation on the ejection velocity of protostellar jets. While large scale knots may successfully be explained by a large increase in the accretion rate, we showed that small scale knots may be produced by changes in the inner radius.  We showed that a stellar magnetic field variation, if present, represents a natural candidate to produce stellar knots similar to the observed. In fact, a stellar magnetic field variation produces a change in the inner radius and therefore in the jet velocity. We also estimated that a stellar magnetic field variation $\\lesssim$ 50\\% produces a variation ($\\sim$ 20\\%) in the average velocity large enough to produce knots similar to the observed.  The stellar magnetic field periodic or quasi-periodic variation remains an hypothesis of our scenario, but future large scale temporal observations of magnetic field in young stars, and progress in dynamo theory, will both help to provide an answer to this problem."
    },
    "0808/0808.1208_arXiv.txt": {
        "abstract": "Previous studies of phase mixing of ion cyclotron (IC), Alfv\\'enic, waves in the collisionless regime have established the generation of parallel electric field and hence acceleration of electrons in the regions of transverse density inhomogeneity. However, outstanding issues were left open. Here we use 2.5D, relativistic, fully electromagnetic PIC (Particle-In-Cell) code and an analytic MHD (Magnetohydrodynamic) formulation, to establish the following points: (i) Using the generalised Ohm's law we find that the parallel electric field is supported mostly by the electron pressure tensor, with a smaller contribution from the electron inertia term. (ii) The generated parallel electric field and the fraction of accelerated electrons are independent of the IC wave frequency remaining at a level of six orders of magnitude larger than the Dreicer value and approximately 20 \\% respectively. The generated parallel electric field and the fraction of accelerated electrons increase with the increase of IC wave  amplitude. The generated parallel electric field seems to be independent of  plasma beta, while the fraction of accelerated electrons strongly increases with the decrease of plasma beta (for plasma beta of 0.0001 the fraction of accelerated electrons can be as large as 47 \\%). (iii) In the collisionless regime IC wave dissipation length (that is defined as the distance over which the wave damps) variation with the driving frequency shows a deviation from the analytical MHD result, which we attribute to a possible frequency  dependence  of the effective resistivity. (iv) Effective anomalous resistivity, inferred from our numerical simulations, is at least four orders of magnitude larger than the classical Spitzer value. ",
        "introduction": "Phase mixing is a mechanism of enhanced dissipation of Alfv\\'en waves due to inhomogeneity of Alfv\\'en speed in a direction transverse to a local magnetic field. This mechanism originally was studied in the fusion and laboratory plasma context by a number of authors \\cite{1972PhFl...15.1673U,1973ZPhy..261..203T,1973ZPhy..261..217G,1974PhRvL..32..454H, 1974PhFl...17.1399C,1975JPlPh..13...87T} and subsequently applied to the solar corona \\cite{1983A&A...117..220H}. Most of the large amount of work done in the field of phase mixing was in the resistive MHD (Magnetohydrodynamic) regime. Recently, a few works looked at the same mechanism in the collisionless regime in the context of Earth magnetosphere \\cite{1999JGR...10422649G, 2004AnGeo..22.2081G} and solar corona \\cite{2005A&A...435.1105T,2005NJPh....7...79T}. The main findings of these works include the generation of electric field that is parallel to the ambient magnetic field in the regions of transverse density inhomogeneity, as well as associated electron acceleration. It should be mentioned that these studies considered circularly polarised ion cyclotron (IC) waves which in the low frequency regime become Alfv\\'en waves. We use terms Alfv\\'en or IC interchangeably, but reader should bear in mind we always refer to {\\it waves with frequencies $< \\omega_{ci}$} (with $\\omega_{ci}$ being ion cyclotron frequency). The exact mechanism of generation of the parallel electric field has stimulated a debate \\cite{2006A&A...449..449M,2007NJPh....9..262T}, and even MHD regime option was explored \\cite{2006A&A...455.1073T}. Continuing this investigation, here we apply technique used in the collisionless reconnection \\cite{pritchett01,th08}. Namely, in section 3.1 we use generalised Ohm's law to find out which term generates the parallel electric field.  Solar flare observations \\cite{2005SSRv..121..141F} trigger one's interest in how effectively plasma particles are accelerated. Hence, in section 3.2 we look into how the generated parallel electric field and the fraction of accelerated electrons depend on model parameters such as IC wave frequency, amplitude, and plasma beta.  Ref. \\cite{1983A&A...117..220H} provides a simple analytical expression how Alfv\\'en wave amplitude should decay in space due to the phase mixing. Despite the fact that their formula is derived in the resistive MHD regime, we still apply it to our collisionless, kinetic simulation and see what does the comparison yield (Section 3.3). This is done in the light of previous results of \\citet{2005A&A...435.1105T} who established that in the collisionless, kinetic regime Alfv\\'en wave amplitude in the density gradient regions decays with distance (from where it is driven) according to collisional MHD formula of \\citet{1983A&A...117..220H}. Here we stretch the MHD-kinetic analogy further to test $\\omega_d^2$ dependence under the exponent.  In Sect 3.4 we estimate the effective \"resistivity\" (again the spirit of MHD-kinetic analogy). The quotation marks are needed to signify that PIC (Particle-In-Cell) simulation code is collisionless and hence no resistive effects exist as such. However, scattering of particles by magnetic fields plays effective role of collisions. ",
        "conclusions": "Let us summarise the above findings:  We used the generalised Ohm's law and found that the parallel electric field, which is generated by propagation of IC (Alfv\\'enic) wave in a transversely inhomogeneous plasma, is supported mostly by the electron pressure tensor, with a smaller contribution from the electron inertia term. Surprisingly, this result resembles closely to the previous results on collisionless reconnection both in tearing unstable Harris current sheet \\cite{hesse99,birn01,pritchett01} and stressed X-point collapse \\cite{th07,th08}. However, in the latter two cases, the generated electric field is in the plane perpendicular to the magnetic field. Thus, a universal importance of the  electron pressure tensor in relation to supporting the electric fields in collisionless plasmas should be noted.   We explored physical parameter space of the problem with regards to the efficiency of generation of parallel electric field and acceleration of electrons. We found that the generated parallel electric field and the fraction of accelerated electrons are independent of the IC wave frequency staying at a level that is $10^6$ times larger than the Dreicer value and approximately 20 \\% respectively. The generated parallel electric field and the fraction of accelerated electrons increase with the increase of IC wave  amplitude. The generated parallel electric field seems to be independent of  plasma-$\\beta$. However, the fraction of accelerated electrons strongly increases with the decrease of plasma-$\\beta$, e.g.  for plasma $\\beta=0.0001$ the fraction of accelerated electrons can be as large as 47 \\%.  Previously it was established  that in the collisionless, kinetic regime phase-mixed Alfv\\'en (IC) wave amplitude damps with distance of propagation according to $\\propto \\exp[-(x/L_d)^3]$ \\cite{2005A&A...435.1105T}, which resembles closely to collisional MHD result of \\citet{1983A&A...117..220H}. We tried to stretch this analogy further by investigating how the dissipation length $L_d$ scales with the IC driving frequency. We found that the scaling is different from the MHD result. We have shown that this discrepancy can be attributed to the frequency dependence of the effective resistivity.  We also found that  the effective resistivity, albeit for unrealistic mass ratio, still is as large as $10^4$ times the classical Spitzer value."
    },
    "0808/0808.1722_arXiv.txt": {
        "abstract": "A class of very energetic supernovae (hypernovae) is associated with long gamma-ray bursts, in particular with a less energetic but more frequent population of gamma-ray bursts. Hypernovae also appear to be associated with mildly relativistic jets or outflows, even in the absence of gamma-ray bursts. Here we consider radiation from charged particles accelerated in such mildly relativistic outflows with kinetic energies of $\\sim$10$^{50}$ erg. The radiation processes of the primarily accelerated electrons considered are synchrotron radiation and inverse-Compton scattering of synchrotron photons (synchrotron self-Compton; SSC) and of supernova photons (external inverse-Compton; EIC). In the soft X-ray regime, both the SSC and EIC flux can be the dominant component, but due to their very different spectral shapes it should be easy to distinguish between them. When the fraction of the kinetic energy going into the electrons ($\\epsilon_e$) is large, the SSC is expected to be important; otherwise the EIC will dominate. The EIC flux is quite high, almost independently of $\\epsilon_e$, providing a good target for X-ray telescopes such as {\\it XMM-Newton} and {\\it Chandra}.  In the GeV gamma-ray regime, the EIC would be the dominant radiation process and the {\\it Gamma-ray Large Area Space Telescope (GLAST)} should be able to probe the value of $\\epsilon_e$, the spectrum of the electrons, and their maximum acceleration energy. Accelerated protons also lead to photon radiation through the secondary electrons produced by the photopion and photopair processes. We find that over a significant range of parameters the proton component is generally less prominent than the primary electron component. We discuss the prospects for the detection of the X-ray and GeV signatures of the mildly relativistic outflow of hypernovae. ",
        "introduction": "\\label{sec:Introduction}  While it is recognized that long gamma-ray bursts (GRBs) are associated with very energetic supernovae, sometime referred to as hypernovae, our understanding of the physics that drives and connects these events is far from being established.  Observationally, long GRBs are more complicated than were thought before. For example, they appear to have subgroup that includes GRB 980425 \\citep{Galama1998}, GRB 031203 \\citep{Malesani2004,Soderberg2004} and GRB 060218 \\citep{Campana2006,Cobb2006,Pian2006}. These GRBs occurred relatively nearby and their energies radiated by prompt gamma rays were significantly smaller than the other long GRBs, and they were associated with well-studied hypernovae (SN 1998bw, SN 2003lw and  SN 2006aj, respectively). Radio observations of these events \\citep{Kulkarni1998,Soderberg2006} suggest the presence of mildly relativistic ejecta, which is a different component from the usual nonrelativistic component of the supernova explosion. Their rate of occurrence may be an order of magnitude higher than that of the more energetic GRBs \\citep{Liang2007,Soderberg2006}. See also, e.g., \\citet*{Liang2006,Murase2006,Toma2007,Waxman2007,Gupta2007} for other followup studies. Very recently, a mildly  relativistic outflow component has also been inferred in a supernova of type Ibc unassociated with a GRB, SN 2008D \\citep{Soderberg2008}.  In this paper, we investigate the broadband radiation from the mildly relativistic ejecta associated with hypernovae, focusing especially on the high-energy photon emission in the X-ray and gamma-ray ranges.  We mainly study the radiation from relativistic electrons which are primarily accelerated in shocks, but we also consider the radiation from secondary electrons which are generated by interactions involving accelerated protons. The latter were extensively studied by \\citet{Asano2008} in connection with the origin of Galactic cosmic rays \\citep{Wang2007,Budnik2008}. Our treatment of the proton component includes also the previously neglected effect of photopair production. We show that the primary electron component photon signature generally dominates over the proton component, for a wide range of energies.  In particular the inverse-Compton scattering of hypernova photons due to the primary electrons appears the most promising channel for detection in the X-ray and gamma-ray ranges.  The paper is organized as follows. In \\S~\\ref{sec:Spectrum of Primary Electrons}, we introduce the relevant supernova parameters and the spectrum of the primary accelerated electrons. We present our main results on the radiation from the primary electrons in \\S~\\ref{sec:Radiation from Primary Electrons}, and discuss its dependence on the parameters. Section~\\ref{sec:Radiation from secondary electrons and proton acceleration} is devoted to a discussion of the photon radiation of a proton origin and its comparison to the electron component. We summarize our conclusions in \\S~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  We have studied the radiation from the mildly relativistic ejecta associated with hypernovae. The energy associated with this portion of the ejecta, about $10^{50}$ erg, starts to dissipate at radii of about 10$^{16}$ cm, where electrons and protons are shock accelerated. The electrons radiate photons via synchrotron, synchrotron-self Compton (SSC) of synchrotron photons, and external inverse-Compton (EIC) scattering of hypernova photons. The radiation is most prominent in the X-ray and GeV ranges, detectable for sources at $d < 100$ Mpc. We also studied the radiation from second and third generations of electrons and positrons that are associated with accelerated protons. The interactions that produce these leptons are the photopion and Bethe-Heitler (BH) processes, and both of these would give fluxes above the sensitivity limit of modern X-ray telescopes. However, we find that both of these components are hidden by a large flux due to the primary electron population mentioned above, for most realistic combinations of relevant parameters.  Values of $\\epsilon_e \\ll 10^{-3}$ or very long lasting observations may improve the detection prospects of this component.  The most promising energy range for detecting the electron SSC and EIC emission components is in the soft X-ray range.  If both $\\epsilon_e$ and $\\epsilon_B$ are reasonably large, the SSC component dominates, while otherwise the EIC does (Fig.~\\ref{fig:F1keV}). The regenerated emission component from electron-positron pairs produced by $\\gamma\\gamma$ absorption of $\\gtrsim 10^2$ GeV photons is found to be always negligible. A robust feature of the EIC emission is that it provides fairly large flux in the X-ray regime, for most combinations of the parameters $\\epsilon_e$, $\\epsilon_B$ and $p$, if the supernova occurs at distances $d \\leq 100$ Mpc. When the two components are comparable, it may be easy to distinguish them through their very different spectral shapes. The GeV flux is dominated by the EIC component. Observations with {\\it GLAST} would make it possible to measure $p$ independently, if $\\epsilon_e$ is large enough, and wold also be able to probe for the  spectral cutoff corresponding to the maximum acceleration energy of the radiating electrons.  Finally, we briefly comment on the dependence of the results on the progenitor wind mass-loss rate.  This is interesting in particular because \\citet{Campana2006} estimated a substantial mass-loss rate $\\dot M \\sim 10^{-4} M_{\\sun}$ for the progenitor of GRB 060218/SN 2006aj, which is an order of magnitude larger than the nominal value adopted here. Such a higher mass loss rate decreases the dissipation radius $R$ and the dynamical time scale $t_{\\rm dyn}$ by an order of magnitude. Even if we decrease the supernova luminosity to $L_{\\rm SN} = 10^{42}$ erg s$^{-1}$, an order of magnitude below the nominal value adopted above, while keeping the other parameters the same, we find that both the SSC and EIC photon fluxes increase by about 2--3 orders of magnitude. This more than offsets the decrease by a factor 3 of the sensitivities of X-ray telescopes caused by the smaller $t_{\\rm dyn}$, implying a very significant further improvement in the prospects for detection."
    },
    "0808/0808.0104_arXiv.txt": {
        "abstract": "We examine the implications of recent measurements of the Milky Way (MW) rotation for the trajectory of the Large Magellanic Cloud (LMC). The $\\sim 14\\pm 6\\%$ increase in the MW circular velocity relative to the IAU standard of $220~\\kms$ changes the qualitative nature of the inferred LMC orbit.  Instead of the LMC being gravitationally unbound, as has been suggested based on a recent measurement of its proper motion, we find that the past orbit of the LMC is naturally confined within the virial boundary of the MW. The orbit is not as tightly bound as in models derived before the LMC proper motion was measured. ",
        "introduction": "\\label{sec:intro}  Recently, Kallivayalil et al. (2006; hereafter K06) measured the proper motion of the the Large Magellanic Cloud (LMC), and pioneered a first assessment of its 3D velocity vector.  Based on this measurement, Besla et al. (2007; hereafter B07) concluded that the LMC was unlikely to have passed near the Milky Way (MW) before and was most likely formed outside of the boundaries of the Galaxy, in contrast to traditional scenarios (see references in van der Marel et al. 2002; hereafter vdM02).  This new conclusion is intriguing given that there are no other examples of massive gas-rich galaxies like the LMC within the much bigger volume that separates the MW from its neighboring galaxy, Andromeda (M31).  Other recent studies have also found the B07 results difficult to accept and propose alternate strategies of binding the LMC and MW either through MOND gravity (Wu et al. 2008) or by giving the LMC and its smaller counterpart, the Small Magellanic Cloud, a common halo (Bekki 2008).   A couple of years after K06 published their findings, Piatek et al. (2008; hereafter P08) confirmed independently their results to within one standard deviation, although at the lower end of the inferred range of values.  Also, within the past five years, the circular velocity of the MW and the distance between the Sun and the galactic center have been updated by Reid \\& Brunthaler (2004; hereafter RB04) and Gillessen et al. (2008; hereafter GGE08), respectively. These increased the likely circular velocity of the MW from the IAU standard of $V_{circ}=220$ \\kms~to 251$\\plm15$ \\kms. A value of 220 \\kms~now corresponds to a reduction of the best-fit value by two standard deviations (and equivalent to moving the Sun a total of 1 kpc closer to the center of the Galaxy).  Uemura et al. (2000) used parallax measurements from Hipparcos and SKYMAP to obtain a similar value of $V_{circ}=255\\plm8$ \\kms.  The rotation speed of the MW affects the analysis of the past LMC orbit in two ways. First, because the proper motion of the LMC is measured relative to the solar system which orbits the Galaxy, it is necessary to know the rotational velocity of the Sun in order to transform to the Galactocentric frame (see, e.g. vdM02).  Second, the depth of the gravitational potential well of the MW (involving its estimated mass and scale radius) depends on the normalization of its rotation curve.  B07 adopted the IAU standard in their analysis instead of the modified values for the Milky Way's rotation. In this {\\it Letter}, we examine the implications of the change in the MW parameters for the LMC orbit (also with the updated P08 value).  In the particular geometry of the LMC orbit, both of the above-mentioned effects make the LMC more gravitationally bound to the MW owing to an increase in $V_{circ}$.  Despite the small fractional magnitude of the correction in circular velocity ($\\sim 14\\pm 6\\%$), we find that the qualitative nature of the LMC orbit changes. Instead of the LMC being possibly unbound as suggested by B07, the $\\sim 10\\%$ decrease in the LMC velocity and the $\\sim 50\\%$ increase in the MW mass lead us to conclude that the LMC's past trajectory was probably confined within the virial radius of the MW.  The apogalacticon distance of the orbit is comparable to the MW virial radius, as expected for a satellite that had formed at the outer edge of the Galactic halo.  Traditional studies of the LMC's orbit around the MW (e.g., Murai \\& Fujimoto 1980; Lin \\& Lynden-Bell 1982; Gardiner et al. 1994; Lin et al. 1995; Gardiner \\& Noguchi 1996; vdM02; Bekki \\& Chiba 2005; Mastropietro et al. 2005; Connors et al. 2006) considered the MW as an isolated galaxy.  While this might have been acceptable for studies where the LMC's orbit was confined well within the MW halo, the high LMC velocity measured by K06 and P08 implies (B07) that the apogalacticon could extend beyond the edge of the MW's halo -- where the gravitational influence of M31 is non-negligible (see Table 1 in B07).  Thus, we also include the tidal effect of M31 in our calculations.  In \\S \\ref{sec:method} we describe our adopted model for the mass distribution of the MW halo. We then calculate the past LMC orbit (\\S \\ref{sec:const}) and the effect of M31 on it (\\S \\ref{sec:M31}). Finally, we discuss the implications of our results in \\S \\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  We have found that the $\\sim 14\\pm 6\\%$ increase in the MW circular velocity, relative to the IAU standard of $220~\\kms$, allows the LMC to be gravitationally bound to the MW. Despite its relatively high proper motion (K06; P08), the orbit of the LMC remains confined within the virial radius of the MW.  Since the MW and M31 galaxies account for most of the mass in the LG, the LG itself is not much more spatially extended than the two are individually, having an estimated zero energy surface at a radius of $\\sim0.9~\\Mpc$~(Karachentsev et al. 2008).  If the LMC had followed the path dictated by the P08(220~\\kms) parameters in Table 2, it would have originated from a distance of $\\ga 1.5$ \\Mpc~ away from the MW center (see Figure \\ref{fig:2dpathlg}).  In comparison, if we admit the P08(251~\\kms) parameters, then the LMC originated roughly at the virial radius of the MW halo and well within the boundaries of the LG. The path suggested by P08(236~\\kms) puts the LMC origin outside the LG but closer than the P08(220~\\kms) case.  If the P08(265~\\kms) model is to be believed, the LMC is on its third pass by the MW center and also originated on the edge of the MW halo.  The Galactocentric distance on which this result is based ($\\rsun=9.0$~kpc) is 1--$\\sigma$ above the most recent estimate (8.4\\plm0.5 \\kpc; see GGE08), so it is not excluded.  We note that during the long time-scale between perigee passes ($\\sim6$ Gyr, see Fig. \\ref{fig:distance}), the MW halo mass could evolve by tens of percent (Diemand, Kuhlen \\& Madau 2007). Future studies might use cosmological simulations to incorporate the evolution of both the host (MW-like) galaxy and its most massive (LMC-like) satellite to get a statistical understanding of their likely past interaction.  Our most likely model has the LMC forming at roughly the MW virial radius, as expected from a hierarchical formation of the MW halo (Springel et al. 2008, Figs. 11 \\& 12). Also, all our quoted uncertainty for the LMC orbit stem from the uncertainty in \\rsun (and therefore \\vcirc).  Errors in the proper motion were not taken into account.  The K06 measurements, for example, are slightly over 1--$\\sigma$ off from the P08 measurements of the LMC proper motion.  The difference between these two cases (as seen in Figure \\ref{fig:distance2}) is drastic -- in one case the LMC is clearly bound to the MW (P08), and in the other it is not (K06).  A 1--$\\sigma$ shift in the other direction, however, would alter the path of the LMC into an even tighter orbit, similar to the P08(265 \\kms) model.  With the value of the distance between the Sun and the Galactic center at its IAU value (GGE08, KLB86), $\\sim$8.5$\\plm.5~\\kpc$, the total mass of the MW Galaxy increases by a factor of 1.23--1.75 relative to the values inferred for the lower distance of 8\\plm.5 \\kpc ~(Reid 1993).  B07 correctly rules out a larger mass for the MW but only considers the low mass ($1\\times10^{12}\\msun$) and the high mass ($2\\times10^{12}\\msun$) models of Klypin et al. (2002) which bracket our preferred range.  Xue et al. (2008) describe the uncertainties of the prior assumption that the galactic \\vcirc~ is 220 \\kms.  Table \\ref{tab:MW} provides the corresponding range of masses for the MW halo, all calculated from observational data, fit either to a flat rotation curve or an NFW profile, such as the one we have adopted in this work.  \\begin{table} \\caption{Milky Way Halo Mass in Recent Studies} \\label{tab:MW} \\begin{tabular}{@{}lcc} \\hline Author & Mass & Model \\\\ ~ & ($10^{12}\\msun$) & ~   \\\\ \\hline \\hline Wilkinson \\& Evans (1999)& $1.9^{+3.6}_{-1.7}$ & FRC$^{1}$ \\\\ Sakamoto, Chiba \\& Beers (2003) & $2.5^{+0.5}_{-1.0}$ & FRC \\\\ Sakamoto, Chiba \\& Beers ({\\it w/o} Leo I) & $1.8^{+.04}_{-0.7}$ & FRC \\\\ Smith et al. (2007) & $1.42^{+1.14}_{-0.54}$ & NFW \\\\ Xue et al. (2008) & $0.79\\plm0.15$ -  & NFW\\\\ ~ & ~~~$1.18\\pm0.28$ & \\\\ Li \\& White (2008)& $2.43$ & N-Body$^{2}$\\\\ \\hline \\end{tabular} $^{1}$FRC denotes a Flat Rotation Curve, as in an isothermal sphere. These mass calculations have cutoffs of $\\sim 50~\\kpc$. $^{2}$e.g. the Millennium simulation.  Here the limiting radius is the virial radius, similar to the NFW calculation. \\end{table}  The $\\sim 50 \\%$ increase in mass (from $1\\times10^{12}~\\msun$ to $1.485\\times10^{12}~\\msun$) and the $\\sim 6\\%$ decrease in the velocity of the LMC (P08(220~\\kms) to P08(251~\\kms)) have a comparable influence on making LMC orbit bound to the MW. The combined effect of these changes is not equivalent to a change in the MW mass, as considered by B07.  The K06(251~\\kms) and P08(220~\\kms) lines in Figures \\ref{fig:distance2} and \\ref{fig:distance} correspond to roughly the same LMC velocity (as derived from the IAU standard with the P08 proper motion) but a different MW mass.  The higher mass of the MW contributes about half of the overall difference between these cases.  This parallels the comparison made in B07 between the K06(220~\\kms) and the K06(220~\\kms)$+4\\sigma$ numbers in their {\\it Fiducial} and {\\it High Mass} Models, except that our findings do not require a $4-\\sigma$ deviation from the measured LMC velocity or a $2-\\sigma$ deviation from the mass of the MW in order to make the LMC orbit bound to the MW.  The previous suggestion (B07) of possibly disqualifying the LMC and SMC as MW satellites has undesirable implications for the satellites in the MW halo.  M31 has 18 satellites, 5 of which are gas rich, whereas the MW, if the LMC and SMC are no longer bound to it, has only 12 bound satellites, 2 of which are gas rich (Karachentsev 2005, vdMG08).  M31's satellites range in mass from 0.58 to 500 $\\times10^{8}\\msun$, whereas the remaining MW satellites are in the range of 0.1--1$\\times10^{8}\\msun$ (Mateo 1998).  Excluding the LMC and SMC as satellites of the MW appears unnatural, as there is no reason for the two comparably sized galaxies to have a large disparity in the abundance of massive satellites.  Moreover, the chance of finding massive galaxies like the LMC and the SMC so close to the MW center requires a special coincidence if they are unbound to the MW, since they would have spent most of their orbital time far away from the MW in that case.  Yet, no similar galaxies are known to exist in the much larger volume between M31 and the MW.  Although these statistical arguments are not definitive, they support indirectly the higher updated values of \\vcirc (RB04) and \\rsun (GGE08) for the MW.  No work on the Magellanic Clouds would be complete without a mention of the Magellanic Stream (MS).  There is $<2\\%$ change in the trajectory over the length of the MS (100 degrees) between the P08(220\\kms) and the P08(251\\kms) models, both of which are well within the error margins in Figure 8 of B07.  We reiterate the concerns presented in B07 that neither tidal stripping nor ram pressure can fully explain the orientation of the Stream.  \\noindent {\\bf Acknowledgments.}  We thank Gurtina Besla and Mark Reid for useful discussions.  This work is supported in part by NASA grant NNX08AL43G, by FQXi, and by Harvard University and Smithsonian Astrophysical Observatory funds."
    },
    "0808/0808.0332_arXiv.txt": {
        "abstract": "While the particle hypothesis for dark matter may be very soon investigated at the LHC, and as the PAMELA and GLAST satellites are currently taking new data on charged and gamma cosmic rays, the need of controlling the theoretical uncertainties affecting the possible indirect signatures of dark matter annihilation is of paramount importance. The uncertainties which originate from the dark matter distribution are difficult to estimate because current astrophysical observations provide rather weak dynamical constraints, and because, according to cosmological N-body simulations, dark matter is neither smoothly nor spherically distributed in galactic halos. Some previous studies made use of N-body simulations to compute the $\\gamma$-ray flux from dark matter annihilation, but such a work has never been performed for the antimatter (positron and antiproton) primary fluxes, for which transport processes complicate the calculations. We take advantage of the galaxy-like 3D dark matter map extracted from the HORIZON Project results to calculate the positron and antiproton fluxes from dark matter annihilation, in a model-independent approach as well as for dark matter particle benchmarks relevant at the LHC scale (from supersymmetric and extra-dimensional theories). We find that the flux uncertainties arise mainly from fluctuations of the local dark matter density, and are of $\\sim$ 1 order of magnitude. We compare our results to analytic descriptions of the dark matter halo, showing how the latter can well reproduce the former. The overal antimatter predictions associated with our benchmark models are shown to lie far below the existing measurements, and in particular that of the positron fraction recently reported by PAMELA, and far below the background predictions as well. Finally, we stress the limits of the use of an N-body framework in this context. ",
        "introduction": "\\label{sec:intro}  The idea that dark matter is made of exotic weakly interacting massive particles (WIMP) is very appealing in the sense that the related existing frameworks in particle physics beyond the standard model (BSM) could solve (i) problems inherent to elementary particle theory, e.g. the force unification scheme and/or the stabilization of the theory against quantum corrections, which would thus confer a more fundamental meaning to this theory; and in the sametime (ii) the dark matter issue as characterized in astrophysics and cosmology (see e.g.~\\cite{review_dm_murayama_07} for a recent review). Indeed, the cold dark matter (CDM) paradigm seems to feature a powerful and self-consistent theory of structure formation, which at present is able to reproduce and explain most of large and mid scale observations (see e.g.~\\cite{2007NuPhS.173....1P} and references therein).  While the Large Hadron Collider at CERN (LHC) is about to start hunting BSM particle physics and may soon provide new insights about the dark matter particle hypothesis, high energy astrophysics experiments are also of strong interest to try to determine the actual nature of dark matter. Indeed, if dark matter is made of self-annihilating particles, a property which is found in many BSM models and which provides a natural mechanism to explain the dark matter cosmological abundance as observed today, there should be traces of annihilations imprinting the cosmic ray spectra (indirect detection of dark matter, see e.g.~\\cite{susy_dm_jungman_etal_96,review_dm_bergstrom_00, review_dm_bertone_etal_05,review_dm_carr_etal_06}). Therefore, satellite experiments such as PAMELA and GLAST, which are dedicated to large field of view observations of charged cosmic rays and $\\gamma$-rays, respectively, in the MeV-TeV energy range, should be able to yield additional constraints or smoking guns very soon in this research field~\\cite{2006JCAP...12..003M}. The most promising astrophysical messengers that could trace the annihilation processes are indeed gamma-rays and antimatter cosmic rays, which have been first investigated in~\\cite{1978ApJ...223.1015G} and in~\\cite{1984PhRvL..53..624S} respectively. Confinement and diffusion on magnetic turbulences limit the origin of the latter to our Galaxy, whereas the former, which only experiences the common $r^{-2}$ flux dilution, may be observed even when emitted from extra-galactic regions.    Nevertheless, predictions of these exotic signals are affected by many uncertainties coming from (i) the underlying WIMP model (ii) the distribution of dark matter in the relevant sources (iii) the propagation of charged cosmic rays in the Galaxy in the case of searches in the antimatter cosmic ray spectra. The first point is encoded in the WIMP mass, the annihilation cross-section and the annihilation final states which fully define the injected cosmic ray spectra: this has been widely investigated for years. The second point is currently surveyed by the state-of-the-art N-body experiments, but the use of N-body dark matter maps in the context of indirect detection has only been performed for $\\gamma$-ray predictions~\\cite{2003MNRAS.345.1313S,2007ApJ...657..262D,2008arXiv0801.4673A, 2008arXiv0805.4416K}. Besides, there has been lots of efforts to estimate the effect of varying the dark matter distribution by means of analytical calculations, allowing for instance to study the effects of sub-halos on the $\\gamma$-ray~\\cite{1999PhRvD..59d3506B, 2002PhRvD..66l3502U,berezinsky_etal_03,berezinsky_etal_06, 2008MNRAS.384.1627P} as well as on the antimatter primary fluxes~\\cite{2007A&A...462..827L,2008A&A...479..427L}. Such analytical studies mostly rely on the statistical information supplied by N-body simulations. However, though they provide very nice theoretical bases to unveil the salient effects of any change in the relevant parameter space, and permit to go beyond the current numerical resolution limits, they are usually bound to simplifying hypotheses, such as the sphericity of sources. Especially, no study of the antimatter signatures has been directly performed in the frame of N-body environments up to now. Finally, point (iii), referring to cosmic ray propagation, has also widely been investigated for antiprotons~\\cite{2004PhRvD..69f3501D,2005JCAP...09..010L} and positrons~\\cite{2008A&A...479..427L,2008PhRvD..77f3527D}, and while the transport processes depend on sets of still degenerate parameters within different propagation models, the origins of uncertainties are now rather well understood.     In this paper, we make use of a realistic (in the sense of \\lcdm\\ cosmology) 3D dark matter distribution of a galactic-sized halo extracted from the HORIZON Project simulation results to study the antimatter signals. Our main purposes are to characterize and quantify the uncertainties coming from dark matter inhomogeneities (not necessarily clumps) and departure from spherical symmetry in a virialized Milky-Way-like object; and finally make predictions for different dark matter particle candidates. Despite the rather low resolution of our numerical data compared with the highest-level artillery on the market, as portrayed, e.g., by the Via Lactea simulations~\\cite{2007ApJ...657..262D,2008arXiv0805.1244D}, we study for the first time the production of antimatter cosmic rays from dark matter annihilation in an N-body framework. The resolution issue is however less important for charged cosmic rays than for gamma-rays, because the signal is strongly diluted by diffusion effects, which therefore tends to smooth the yields from high density fluctuations, provided the latter are not too close to the observer. Nevertheless, as very local density fluctuations are expected to affect the antimatter signals, we will quantify and correct the consequence of loss of resolution by analytically extrapolating our results. For the WIMP candidates, we will use a model-independent approach as well as typical models of BSM particle theories, some of them being observable at the LHC.   This article is parted as follows: we recall the salient features of the HORIZON N-body experiment in Sect.~\\ref{sec:horizon}; then, in Sect.~\\ref{sec:crs}, we briefly review the positron and antiproton propagation before sketching the method we adopt to connect the propagation to the N-body source terms; we describe the WIMP models in Sect.~\\ref{sec:wimp}; we finally present and discuss our results and predictions of the primary positron and antiproton fluxes in Sect.~\\ref{sec:res} (where the expert reader is invited to go directly) before concluding. ",
        "conclusions": "\\label{sec:concl}   In this paper, we have studied the theoretical uncertainties associated with dark matter density fluctuations affecting the primary antimatter cosmic ray fluxes (positrons and antiprotons), possibly produced by dark matter annihilation in the Galaxy, by using an N-body simulation together with analytical descriptions of a galaxy-like dark matter halo. This is the first attempt in this research field to directly connect the source term of the cosmic ray propagation equation to a 3D density map coming from a cosmological N-body framework, while this has already often been used for predictions of gamma-ray fluxes. Indeed, the latter situation does not involve cosmic ray propagation, which seriously complicates the calculations, and makes them much more time consuming. The main purposes of this work was to quantify those uncertainties, and in some cases to try to understand them by analytical means. While this was achieved in a WIMP model-independent setting, we also aimed at reviewing the status of predictions for many different well motivated dark matter candidates in particle physics, supersymmetric or not.   The N-body framework led us to study several effects, such as non-spherical dark matter distributions (Sect.~\\ref{subsubsec:ellipt}) and density fluctuations (Sect.~\\ref{subsubsec:inhom}), which are not accounted for when using the traditional spherical smooth halo models. We have shown that, as expected, fluxes are mainly affected at cosmic ray energies corresponding to short propagation characteristic lengths (low/high for antiprotons/positrons) and to local contributions, while when looking at energies of large propagation lengths, the smoothing due to diffusion erases the peculiarities in the spectra, and is well reproduced by a spherical halo model. The main uncertainties on the predicted fluxes come therefore from those on the local dark matter environment, which we have shown to fluctuate a lot in the N-body framework. This can lead to $\\pm$ 1 order of magnitude in terms of flux, at high energy for positrons, and low energy for antiprotons. This could be more constrained with much better limits on the local dark matter density.   Moreover, though our simulation has a too poor spatial resolution to resolve substructures inside the galactic halo, we have extrapolated our results in order to include the potential effects of the presence of sub-halos by using the analytical method detailed in~\\cite{2008A&A...479..427L}. We have considered two situations: one \\emph{maximal}, analytically extending the clump mass spectrum down to $10^{-6}\\Msun$, led to a flux enhancement by an energy-dependent factor reaching $\\sim 4$ at maximum; a second, \\emph{Via-Lactea-like}, involved clumps down to $10^6\\Msun$ only and did not result in any flux enhancement.    Consequently to the previous points, the smooth and spherical modeling of the dark matter halo leads to a very good approximation of a more complete and complex situation, as derived from an N-body simulation, at large characteristic propagation lengths for antimatter cosmic rays (low/high energies for positrons/antiprotons). Nevertheless, though predictions in the spectral regions associated with those large characteristic scales are also the less affected by the uncertainties coming from dark matter density fluctuations, we remind that they are still very sensitive to the used propagation model.    Beside this trial for directly \\emph{measuring} the theoretical uncertainties with an N-body experiment, we reviewed the predictions of the positron and antiproton primary contribution associated with some specific particle physics dark matter candidates (Sect.~\\ref{subsec:dm_models}). We have not only included some popular supersymmetric and extra-dimensional models, but also some more specific ones based on minimality arguments, such as the little Higgs model, or the inert doublet model. We have shown that predictions are well lower than the existing data, except perhaps for the antiproton flux associated with the lightest neutralino model, just because of the more favorable $1/\\mchi^2$ factor appearing in the flux expression. Nevertheless, even if the latter appears in tension with the current data, its contribution to the antiproton flux occurs at energies where solar modulation effects are expected to be important, and it would be hard to find a clear signature in this spectral region. The other models are well below the experimental measures, but higher energy measurements could perhaps unveil some spectral transitions followed by an energy cut-off which could possibly be attributed to dark matter. In this case, uncertainties associated with the dark matter density fluctuations would be much less stringent for the antiproton signal than for the positron because of propagation scale arguments. On the contrary, a nearby clump would favor a detection with positrons at high energy, much more than with antiprotons. Anyway, higher energy data are necessary in this field, for these antimatter species as well as for standard nuclei species which will yield much better constraints to the propagation modeling. Obviously, indications of a WIMP candidate mass, as expected from the LHC, would provide the relevant spectral region where to concentrate the astrophysical searches, and would allow to concentrate on much more subtle effects.    Finally, we made a self-criticism exercise -- Sect.~\\ref{subsec:dm_limits} -- by comparing our different halo models to the star kinematic data available for our Milky-Way Galaxy, after subtraction of the baryonic contribution as modeled by~\\cite{2006CeMDA..94..369E}. We show that the dark matter contribution to the velocity field as inferred from our N-body simulation, as well as more general ones, are far to reproduce the kinematic data, systematically leading to a mass excess in the central regions ($\\lesssim 4$ kpc) of the Galaxy. This clearly calls for restraint in our potential claims, and clearly points towards some large amount of ignorance in the intimate (and gravitational) relations between dark matter and baryons."
    },
    "0808/0808.2047_arXiv.txt": {
        "abstract": "We provide a new derivation of the anisotropies of the cosmic microwave background (CMB), and find an exact expression that can be readily expanded perturbatively.  Close attention is paid to gauge issues, with the motivation to examine the effect of super-Hubble modes on the CMB.  We calculate a transfer function that encodes the behaviour of the dipole, and examine its long-wavelength behaviour.  We show that contributions to the dipole from adiabatic super-Hubble modes are strongly suppressed, even in the presence of a cosmological constant, contrary to claims in the literature.  We also introduce a naturally defined CMB monopole, which exhibits closely analogous long-wavelength behaviour.  We discuss the geometrical origin of this super-Hubble suppression, pointing out that it is a simple reflection of adiabaticity, and hence argue that it will occur regardless of the matter content. ",
        "introduction": "The anisotropies in the cosmic microwave background (CMB) reveal a great deal about our Universe, since they persist essentially unscathed from the epoch when fluctuations were well described by simple linear theory.  The comparison of CMB observations with theory has become a mature subject, and has played an important role in forming our current understanding of the Universe (see, \\eg, \\cite{wmap5}).  Critical to that comparison is the accurate theoretical calculation of the anisotropies.  Since the pioneering work of Sachs and Wolfe \\cite{sw67}, the theoretical anisotropies have been refined and recalculated using different formalisms many times (see, \\eg, \\cite{panek86,magueijo92,rsxd93,wh97,dunsby97,cl98,hn99}).  Accurate calculations are now readily available via public code packages such as \\textsc{camb} \\cite{lcl00,notecamb}.  In the present work, we revisit the calculation of anisotropies. While our results may not lead to more accurate or efficient calculations, we hope that they will help to clarify some of the conceptual issues surrounding the calculations.  In particular, our approach makes explicit the physical meaning of the various contributions to the anisotropies.  Crucial to this is our use of the covariant approach to cosmology (see \\cite{ee98,tcm08} for reviews), which is ideal for writing exact solutions and for physical clarity. We present a remarkably simple but exact expression for the anisotropy, which applies to arbitrary spacetimes and includes the effects of tensor as well as scalar \\perts\\ and any line-of-sight integrated Sachs-Wolfe (ISW) effect.  This general result can be readily expanded perturbatively, and here we turn to the metric formalism for computational efficiency and show that we recover previous results in the literature.  A main motivation for our work is in examining the behaviour of the anisotropies due to super-Hubble fluctuations (we use the terms ``super-Hubble'', ``long-wavelength'', and ``large-scale'' interchangeably, to mean scales larger than the current Hubble or last scattering radius), where gauge issues are paramount.  This question has been examined before in the context of the Grishchuck-Zel'dovich effect \\cite{gz78}, which describes the large-angular-scale anisotropies that result from super-Hubble modes.  In the context of a matter-dominated universe with adiabatic fluctuations, it was shown that the CMB dipole receives strongly suppressed contributions from long-wavelength modes.  A claim was made in Ref.~\\cite{turner91} that this suppression would not occur in models with cosmological constant, so that we could ``see'' very long-wavelength structure in the dipole.  It was also found that in the presence of {\\em isocurvature} \\perts, the suppression may not occur (see \\cite{langlois96} and references therein).  To study this issue, we construct a transfer function that describes the scale dependence of contributions to the dipole.  Working by analogy, we carefully define a CMB {\\em monopole} \\pert, and find its transfer function.  As is well known, such a monopole cannot be observable, but we show that its {\\em variance} is well defined theoretically.  Our definitions have simple interpretations:  the dipole measures the departure of radiation and matter comoving worldlines, while the monopole measures how well radiation and matter constant-density hypersurfaces coincide.  The usefulness of the monopole will be in examining its long-wavelength behaviour, where it will help to clarify the dipole case.  We show that the contributions to both dipole and monopole vanish for large scale sources, even in the presence of a cosmological constant.  We close by pointing out that this is a direct consequence of adiabaticity, and hence that this result is expected to hold regardless of the matter content of the Universe, unless isocurvature modes are present.  Another potential reason that a careful treatment of the monopole may be of interest involves the measurement of the mean CMB temperature and its relation to constraints on other cosmological parameters. The mean temperature is currently measured to a precision of a few parts in $10^4$ \\cite{mfsmw99}.  It has been pointed out that the precision of this measurement could be improved by nearly two orders of magnitude with currently available technology \\cite{fm02}.  Such a measurement would reach the naive cosmic variance limit of a part in $10^5$ as suggested by the observed amplitude of fluctuations (see \\cite{wz08} for a related discussion).  It would then become necessary to be very careful about exactly what information the mean temperature measurement is giving us, and about the nature of monopole fluctuations.  Relevantly, recent studies have examined the importance of the mean temperature measurement to our ability to constrain the cosmological parameters \\cite{cs08,hw08}.  Since the covariant approach to cosmology is essential to this work, we begin in Sec.~\\ref{covarsec} with a summary of the required formalism. Next, in Sec.~\\ref{swsec} we present the derivation of the Sachs-Wolfe effect, beginning with an exact result before specializing to first order and recovering previous results.  In Sec.~\\ref{secmonodi} we present calculations of the dipole and monopole transfer functions, and we examine their long-wavelength behaviour in Sec.~\\ref{seclongwl}. Finally we discuss our results in Sec.~\\ref{discusssec}.  The Appendices summarize relevant material in the metric formalism, and demonstrate both the gauge invariance and the gauge dependence of our results. We use signature $(-,+,+,+)$, and greek indices indicate four-tensors, while latin indices indicate spatial three-tensors. ",
        "conclusions": "\\label{discusssec}   Our result in Sec.~\\ref{seclongwl} that the dipole and monopole receive suppressed contributions from large scales, even in the presence of a dominant cosmological constant, strongly suggests that the cancellations involved are not accidental.  To understand the origin of this behaviour, consider first the monopole case.  Physically, the suppression as $k \\ra 0$ in the monopole transfer function shown in Figs.~\\ref{transfnc} and \\ref{monotrans} means that surfaces of uniform total and radiation energy density coincide on the largest scales, according to our definition of the monopole in Sec.~\\ref{secmono}.  But this is just the statement of adiabaticity: an adiabatic matter-radiation fluid is characterized by the condition \\beq \\fr{\\del\\rho}{\\dot\\rho} = \\fr{\\del\\rho_{(\\gam)}}{\\dot\\rho_{(\\gam)}}, \\label{adcond} \\eeq on the largest scales.  Therefore, \\eq(\\ref{uedgtrns}) shows that the same gauge transformation takes us to both constant total matter and constant radiation \\hsfs; in other words, those surfaces must coincide.  It is important to point out that adiabatic long-wavelength modes remain adiabatic under evolution \\cite{mfb92}, and indeed the comoving curvature \\pert\\ $\\R$ remains constant in time (see, \\eg, \\cite{ll00}).  This trivial evolution on super-Hubble scales means that when constant matter and radiation density surfaces coincide on large scales at last scattering, due to adiabatic initial conditions, they must also coincide today.  An analogous situation holds for the dipole.  In this case, the adiabaticity condition implies that, on large scales, the radiation and total matter comoving worldlines coincide (\\ie, there is no ``peculiar velocity'' isocurvature mode between the two components). According to our definition of the dipole in Sec.~\\ref{dipsubsec}, this is simply the statement that the dipole is suppressed on large scales.  This leads us to the important conclusion that this insensitivity to long-wavelength sources must apply regardless of the matter content of the Universe, as long as adiabaticity holds.  Therefore the suppression we found for the specific case of cosmological constant (or Einstein-de~Sitter) universes must in fact occur in general.  Note that this may have relevance to very recent discussions regarding a potential power asymmetry in the CMB \\cite{ekc08}.  The exquisite cancellations visible at large scales in Figs.~\\ref{ditrans} and \\ref{monotrans} between the SW and ISW components illustrate a previously unrecognized relation between the two, which is enforced by the condition of adiabaticity.  In brief, the dipole and monopole as defined here are just measures of departures from the adiabaticity condition, \\eq(\\ref{adcond}), which generally occur on small scales.  Are these results sensitive to the definitions of dipole and monopole used?  As long as the dipole and monopole are defined {\\em physically}, \\ie\\ in relation to locally measureable quantities, then the results must still hold.  An important example which does {\\em not} satisfy this criterion is the zero-shear, or longitudinal gauge, since it is not defined in terms of local, observable matter quantities.  If the dipole (or monopole) is defined with respect to a zero-shear frame, then it may appear from theoretical calculations that the dipole {\\em is} sensitive to long-wavelength modes.  However, this sensitivity cannot be observable, since zero-shear frames cannot be {\\em uniquely} constructed locally. (A linear boost or ``tilt'' of a zero-shear frame is still a zero-shear frame.)  Thus great care must be taken when considering behaviour on large scales with longitudinal gauge.  Indeed, another way to understand this result is to realize that, in the limit $k/(aH) \\ra 0$, an adiabatic \\pert\\ mode locally becomes essentially pure gauge, and can be removed within any {\\em sub-Hubble} region by a simple boost coordinate transformation.  Such a mode is indistinguishable locally from a homogeneous background, and hence cannot have any observational consequences such as a dipole anisotropy \\cite{unruh98}.  Finally, we note that when the condition of adiabaticity is relaxed, then our conclusions no longer hold \\cite{langlois96}.  In the presence of isocurvature \\perts, it is possible that a {\\em physically} defined dipole be sensitive to super-Hubble modes, since the extra freedom allows for a relative tilt between comoving matter and radiation comoving worldlines on large scales.  While we have emphasized here the consequences of adiabaticity, it is hoped that other applications will follow from the exact formalism for CMB anisotropies that we have developed.  One possibility is the evaluation of anisotropies, in particular the ISW effect, in void models of acceleration \\cite{celerier99} (see \\cite{enqvist08} for a brief review), which have not yet been confronted with observations at the perturbative level \\cite{zibin08}.  Note however that the present approach is not limited to calculating CMB anisotropies, and that more generally it is applicable to calculating redshifts in arbitrary spacetimes.  Potential uses include calculating the redshift-luminosity distance relation in perturbed spacetimes (see, \\eg, \\cite{hg06}).  {\\em Note added:}  When this work was essentially complete, a related paper appeared \\cite{eck08}, which appears to support our conclusion that long-wavelength \\perts\\ cannot effect the CMB dipole."
    },
    "0808/0808.3088_arXiv.txt": {
        "abstract": "{Disk galaxies with a spheroidal component are known to host Supermassive Black Holes (SMBHs) in their center. Unequal-mass galaxy mergers have been rarely studied despite the fact that they are the large majority of merging events by number and they are associated with the typical targets of gravitational wave experiments such as LISA. We perform N-body/SPH simulations of disk galaxy mergers with mass ratios 1:4 and 1:10 at redshifts z=0 and z=3. They have the highest resolution achieved so far for merging galaxies, and include star formation and supernova feedback. Gas dissipation is found to be necessary for the pairing of SMBHs in these minor mergers. Still, 1:10 mergers with gas allow an efficient pairing only at high z when orbital times are short enough compared to the Hubble time. ",
        "introduction": "Observations of nearby galaxies show that Supermassive Black Holes (SMBHs) ranging between $\\sim10^6$ and $\\sim10^9 M_\\odot$ inhabit the centers of virtually all massive galactic spheroids, from massive ellipticals to pseudo-bulges of late-type galaxies. Their masses, as inferred from dynamical measurements, appear to be correlated with various properties of their hosts, e.g. bulge luminosity and mass (e.g. \\citet{kormendyrichstone95, magorrian98, marconihunt03}), velocity dispersion \\citep{tremaine02, ferrarese00}, concentration of the light profile \\citep{graham01}. In the $\\Lambda$CDM model, galaxies build up their masses hierarchically starting from initial, gravitationally amplified fluctuations (e.g. \\citet{whiterees78}); therefore, every time two galaxies merge, the remnant is expected to host two (or more) SMBHs. The formation of a binary SMBH sistem has been shown to proceed quickly once both the compact objects are embedded in a circumnuclear gaseous disk (see \\citet{lucio07}; see also Mayer, Kazantzidis \\& Escala in this volume), but whether the large scale merger can lead them to such a favorable configuration is still a matter of debate: while mergers between galaxies of equal mass have been quite extensively explored in literature and seem to lead to the formation of a SMBH pair \\citep{stelios05,springel05}, little attention has been paid to the fate of SMBHs in events which are much more frequent in the typical history of a $\\Lambda$CDM galaxy, i.e. minor mergers with mass ratios from 1:4 to 1:10 and less \\citep{stewart08}. ",
        "conclusions": "Understanding the formation of SMBH binaries is of fundamental importance for the search of gravitational waves as well as for all studies of black hole demography and host galaxy coevolution. SMBH pairing in unequal-mass mergers depends crucially on gasdynamical effect: satellites are disrupted too quickly, leaving wandering SMBHs in all the collisionless cases we studied. {\\it The presence of gas seems to be necessary for the pairing of SMBHs}. While it also appears sufficient for $q=0.25$, lower mass objects are more heavily affected by feedback from star formation, and their outcome could be very sensitive to the details of the physical processes involved."
    },
    "0808/0808.0161_arXiv.txt": {
        "abstract": "Ultra-high energy cosmic rays (UHECRs) accelerated in the jets of active galactic nuclei can accumulate in high magnetic field, $\\sim 100$ kpc-scale regions surrounding powerful radio galaxies. Photohadronic processes { involving UHECRs and photons of the extragalactic background light} make ultra-relativistic electrons and positrons that initiate electromagnetic cascades, leading to the production of a $\\gamma$-ray synchrotron halo. We calculate the halo emission in the case of Cygnus A and show that it should be detectable with the { Fermi Gamma ray Space Telescope} and possibly detectable with ground-based $\\gamma$-ray telescopes if radio galaxies are the sources of UHECRs. ",
        "introduction": "Active galactic nuclei (AGN) and gamma-ray bursts are considered as two of the most plausible classes of astrophysical accelerators of extragalactic ultra-high energy cosmic rays \\citep[UHECRs; see, e.g.,][]{hh02}. The recent report of the \\citet{Auger07} about clustering of the arrival directions of UHECRs with energies $E \\gtrsim 6\\times 10^{19}$ eV within $\\approx 3^\\circ$ of the directions to AGN at distances $d\\lesssim 75\\, \\rm Mpc$ strongly suggests that effective production of cosmic rays with energies  $E\\sim 10^{20}\\,\\rm eV$ takes place in at least one of these source classes. Because of photohadronic GZK interactions of protons or ions with the CMB radiation, the study of super-GZK UHECRs from sources at $d\\gtrsim 100\\,\\rm Mpc$ becomes impossible with cosmic-ray detectors like the Pierre Auger Observatory \\citep{hmr06}. Powerful AGN, as well as GRBs, are mostly located at larger distances.  { Relativistic beams of energy from the central nuclei of AGN are thought to power the multi-kpc scale radio lobes of powerful galaxies and form an extended cavity \\citep{sch74}. Acceleration of UHECRs in the compact inner jets of the radio galaxy on the sub-parsec scale, followed by production of collimated beams of ultra-high energy neutrons and gamma-rays, provides a specific mechanism to transport energy to the radio lobes and cavity \\citep{ad03}.} When the neutron-decay UHECR protons interact with the extragalactic background light (EBL), which is dominated by the CMB radiation, ultra-relativistic electrons (including positrons) and $\\gamma$ rays with $E\\gtrsim 10^{18}\\,\\rm eV$ are produced. Such secondaries can initiate pair-photon cascades to form large multi-Mpc scale halos of GeV/TeV radiation due to Compton and synchrotron processes \\citep{acv94,aha02,ias05}.   In this Letter, we predict that synchrotron GeV fluxes from UHECR AGN sources are detectable with the Fermi Gamma ray Space Telescope (FGST; formerly the Gamma ray Large Area Space Telescope, GLAST) if UHECRs are captured in the vicinity of radio galaxies for sufficiently long times. Magnetic fields at the $\\gtrsim \\mu$G level in the kpc -- Mpc vicinity from the AGN core  are required to isotropize UHECRs accelerated by jets of radio galaxies.  Indeed, for protons with energy $E\\equiv 10^{20} E_{20} \\,\\rm eV$, the mean magnetic field $B$ required to provide gyroradii smaller than size $r$ is $B\\gtrsim 10^{-4} E_{20} r_{\\,\\rm kpc}^{-1}\\,\\rm G$, where $r_{\\rm kpc} =r/1\\,\\rm kpc$ is the spatial scale where the isotropization occurs.  Here we consider the specific case of the powerful radio galaxy Cygnus A, { where the mean magnetic field in the surrounding cavity could reach 10 -- 100 $\\mu$G at 100 kpc scales}. Its properties are considered in Section 2, and calculations are presented in Section 3. We summarize in Section 4. ",
        "conclusions": "If the radio lobes of Cyg A are powered by UHECR production from the inner pc-scale jets, then trapping of these particles in the surrounding strong magnetic-field region  leads to the production of secondary $\\gamma$ rays that should be significantly detected with the FGST in one year of observation if the jet injection age is $\\gtrsim 100$ Myr. If radio galaxies are not the sources of UHECRs, then Cygnus A will not be detected by the FGST. Cyg A could also be detected with VERITAS in a 50 hour pointing, depending on the duration of activity of the central engine and the level of the EBL.  Detection of GeV $\\gamma$ rays from Cyg A with the FGST might also be expected to arise from other processes. The $\\sim 10$ -- 100 GeV radio-emitting electrons from the lobes of radio galaxies will Compton-scatter CMB photons to MeV -- GeV energies \\citep[e.g.,][]{che07}.  For the strong magnetic field, $\\approx 60~\\mu$G, in the lobes of Cyg A, however, the ratio of the magnetic-field to CMBR energy densities is $\\approx 400$. Thus the total energy flux of Compton-scattered CMBR from Cyg A is $\\approx 10^{45}$ erg s$^{-1}/[400(4\\pi d^2)] \\cong 4\\times  10^{-13}$ ergs cm$^{-2}$ s$^{-1}$, with the $E\\cdot F(E)$ flux a factor of $\\sim 5$ -- 10 lower. As can be seen from Fig.\\ \\ref{f2}, this process is almost two orders of magnitude below the UHECR-induced synchrotron flux, and well below the FGST sensitivity.  { \\citet{ias05} considered fluxes expected from the $\\lesssim 1 \\,\\rm Mpc$ halos of clusters of galaxies with weaker magnetic fields, $B\\simeq 0.1$-$1\\,\\mu$ G in a model where acceleration of UHECRs occurs in accretion shocks in the cluster. Because of lower maximum energies of accelerated protons, $E\\lesssim 10^{19}\\,\\rm eV$, and lower magnetic fields, this model predicts hard spectral fluxes of Compton origin peaking at TeV energies}. \\citet{ga05} predicted that synchrotron radiation from $\\gtrsim 10^{18}\\,\\rm eV$ electrons is produced by secondaries of UHECRs that leave the acceleration region and travel nearly rectilinearly through weak intergalactic magnetic fields at the level $B \\sim 10^{-7}$ -- $10^{-9}\\,$G. These sources would appear as point-like quiescent GeV -- TeV sources with spectra in the GeV domain as hard as $\\alpha \\cong -1.5$, and very soft, $\\alpha \\lesssim -3$ spectra in the 100 GeV -- TeV domain.  { In contrast to both these models}, we predict soft 0.1 -- 1 GeV spectra with $\\alpha \\cong -2.5$ and hard, $\\alpha \\cong -2$ spectra at TeV energies due to the much higher magnetic field in the confinement region. These models can be clearly distinguished if Cygnus A is resolved by the Fermi Gamma ray Space Telescope or the ground-based $\\gamma$-ray telescopes, as the emission region in our model subtends an angle $\\approx 10^\\prime$.  Because Cygnus A lies outside the GZK horizon, only UHECRs with energy below the GZK energy could be correlated with this source. Other closer FRII radio galaxies that are correlated with the arrival directions of UHECRs are, however, candidate sources of $\\gamma$ rays made through the mechanism proposed here. IGR J21247+5058 at $z = 0.02$ or $d\\approx 80$ Mpc, recently  discovered with INTEGRAL \\citep{mol07}, is 2.1 degrees away from a HiRes Stereo event with $E > 56$ EeV (C.\\ C.\\ Cheung, private communication, 2008).\\footnote{This radio galaxy was identified as such only recently and would not have appeared in the list of AGN used by the HiRes collaboration in their correlation study \\citep{abb08}.} PKS 2158-380 at $\\approx 140$ Mpc is also within 3.2$^\\circ$ degrees of an Auger UHECR with $E> 57$  EeV \\citep{mos08}. By comparison with Cyg A, these are low luminosity FRIIs, and their predicted flux level will require detailed modeling for each source, as done here for Cyg A. Variability of $\\gamma$-ray emission would rule out our model."
    },
    "0808/0808.2813_arXiv.txt": {
        "abstract": "We make a case for the existence for ultra-massive black holes (UMBHs) in the Universe, but argue that there exists a likely upper limit to black hole masses of the order of $M \\sim 10^{10} \\msun$. We show that there are three strong lines of argument that predicate the existence of UMBHs: (i) expected as a natural extension of the observed black hole mass bulge luminosity relation, when extrapolated to the bulge luminosities of bright central galaxies in clusters; (ii) new predictions for the mass function of seed black holes at high redshifts predict that growth via accretion or merger-induced accretion inevitably leads to the existence of rare UMBHs at late times; (iii) the local mass function of black holes computed from the observed X-ray luminosity functions of active galactic nuclei predict the existence of a high mass tail in the black hole mass function at $z = 0$. Consistency between the optical and X-ray census of the local black hole mass function requires an upper limit to black hole masses. This consistent picture also predicts that the slope of the $M_{\\rm bh}$-$\\sigma$ relation will evolve with redshift at the high mass end. Models of self-regulation that explain the co-evolution of the stellar component and nuclear black holes naturally provide such an upper limit. The combination of multi-wavelength constraints predicts the existence of UMBHs and simultaneously provides an upper limit to their masses. The typical hosts for these local UMBHs are likely the bright, central cluster galaxies in the nearby Universe. ",
        "introduction": "Observations of black hole demographics locally is increasingly providing a strong constraint on models that explain the assembly and growth of black holes in the Universe. The existence of a tight relation between the velocity dispersion of bulges and the mass of the central black hole has been reported by several authors (Merritt \\& Ferrarese 2001; Tremaine et al. 2002; Gebhardt et al. 2003). This correlation is tighter than that between the luminosity of the bulge and the mass of the central black hole (Magorrian et al. 1998).  The physical processes that set up this correlation are not fully understood at the present time, although there are several proposed explanations that involve the regulation of star formation with black hole growth and assembly in galactic nuclei (Haehnelt, Natarajan \\& Rees 1998; Natarajan \\& Sigurdsson 1998; Silk \\& Rees 1999; Murray, Quataert and Thompson 2004; King 2005).   Recent work by several authors has suggested that UMBHs\\footnote{Black holes with masses in excess of $5 \\times 10^9\\,M_{\\odot}$ are hereafter referred to as UMBHs.} ought to exist: Bernardi et al. (2006) show that the high velocity dispersion tail of the velocity distribution function of early-type galaxies constructed from the Sloan Digital Sky Survey (SDSS) had been under-estimated in earlier work suggestive of a corresponding high mass tail for the central black hole masses hosted in these nuclei. As first argued by Lauer et al. (2007a) and subequently by Bernardi et al. (2007) and Tundo et al. (2007), even when the scatter in the observed $M_{\\rm bh} - \\sigma$ correlation is taken into account it predicts fewer massive black holes compared to the $M_{\\rm bh} - L_{\\rm bulge}$ relation. While Bernardi et al. (2007) argue that this is due to the fact that the $\\sigma - L_{\\rm bulge}$ relation in currently available samples is inconsistent with the SDSS sample from which the distributions of $L_{\\rm bulge}$ or $\\sigma$ are based. From an early-type galaxy sample observed by HST, Lauer et al. (2007b) argue that the relation between $M_{\\rm bh} - L_{\\rm bulge}$ is likely the preferred one for BCGs (Brightest Cluster Galaxies) consistent with the harboring of UMBHs as evidenced by their large core sizes. The fact that the high mass end of the observed local black hole mass function is likely biased is a proposal that derives from optical data. Deriving the mass functions of accreting black holes from optical quasars in the Sloan Digital Sky Survey Data Release 3 (SDSS DR3), Vestergaard et al. (2008) also find evidence for UMBHs in the redshift range $0.3 \\leq z \\leq 5$.  In this paper, we show that UMBHs exist using X-ray and bolometric AGN luminosity functions and for consistency with local observations of the BH mass density, an upper limit to their masses is required. To probe the high mass end of the BH mass function, in earlier works the AGN luminosity functions were simply extrapolated. This turns out to be inconsistent with local estimates of the BH mass function. Here we focus on the high mass end of the predicted local black hole mass function, i.e.  extrapolation of the $M_{\\rm bh} - \\sigma$ relation to higher velocity dispersions and demonstrate that a self-limiting cut-off in the masses to which BHs grow at every epoch reconciles the X-ray and optical views.  The outline of this paper is as follows: in Section 2, we briefly summarise the current observational census of black holes at high and low redshift including constraints from X-ray AGN. The pathways to grow UMBHs are described in Section 3. Derivation of the local black hole mass function from the X-ray luminosity functions of AGN is presented in Section 4. The arguement for the existence of an upper limit to black hole masses from various lines of evidence is presented in Section 5; the prospects for detection of this population is presented in Section 6 followed by conclusions and discussion. We adopt a cosmological model that is spatially flat with $\\Omega_{\\rm matter} = 0.3$; $H_0 = 70\\,{\\rm km~s^{-1}}/{\\rm Mpc}$. ",
        "conclusions": "The interplay between the evolution of BHs and the hierarchical build-up of galaxies appears as scaling relations between the masses of BHs and global properties of their hosts such as the BH mass vs. bulge velocity dispersion - the $M_{\\rm bh} - \\sigma_{\\rm bulge}$ relation and the BH mass vs. bulge luminosity $M_{\\rm bh} - L_{\\rm Bulge}$ relation. The low BH mass end of this relation has recently been probed by Ferrarese et al. (2006) in an ACS survey of the Virgo cluster galaxies. They find that galaxies brighter than $M_B \\sim -20$ host a supermassive central BH whereas fainter galaxies host a central nucleus, referred to as a central massive object (CMO). Ferrarese et al. report that a common $M_{\\rm CMO} - M_{\\rm gal}$ relation leads smoothly down from the scaling relations observed for more more massive galaxies. Extrapolating observed scaling relations to higher BH masses to the UMBH range, we predict that these are likely hosted by the massive, high luminosity, central galaxies in clusters with large velocity dispersions. The velocity dispersion function of early-type galaxies measured from the SDSS points to the existence of a high velocity dispersion tail with $\\sigma > 350\\,{\\rm kms}^{-1}$ (Bernardi et al. 2006). If the observed scaling relations extend to the higher mass end as well, these early-types are the most likely hosts for UMBHs.  Recent simulation work that follows the merger history of cluster scale dark matter halos and the growth of BHs hosted in them by Yoo et al.(2007) also predict the existence of a rare population of local UMBHs.  However, theoretical arguments suggest that there may be an upper limit to the mass of a BH that can grow in a given galactic nucleus hosted in a dark matter halo of a given spin. Clearly the issue of the existence of UMBHs is intimately linked to the efficiency of galaxy formation and the formation of the largest, most luminous and massive galaxies in the Universe.  Possible explanations for the tight correlation observed between the velocity dispersion of the spheroid and black hole mass involve a range of self-regulated feedback prescriptions. An estimate of the upper limits on the black hole mass that can assemble in the most massive spheroids can be derived for all these models and they all point to the existence of UMBHs.  In this paper, we have argued that while rare UMBHs likely exist, there is nevertheless an upper limit of $\\sim 10^{10}\\,\\msun$ for the mass of BHs that inhabit galactic nuclei in the Universe. We first show that our current understanding of the accretion history and mass build up of black holes allows and implies the existence of UMBHs locally. This is primarily driven by new work that predicts the formation of massive black hole seeds at high redshift (Lodato \\& Natarajan 2007) and their subsequent evolution (Volonteri, Lodato \\& Natarajan 2008). Starting with massive seeds and following their build-up through hierarchical merging in the context of structure formation in a cold dark matter dominated Universe, we show that a viable pathway to the formation of UMBHs exists. There is also compelling evidence from the observed evolution of X-ray AGN for the existence of a local UMBH population. Convolving the observed X-ray LF's of AGN, with a simple accretion model, the mass function of black holes at $z = 0$ is estimated. Mimic-ing the effect of self-regulation processes that impose an upper limit to BH masses and incorporating this into the X-ray AGN LF we find that the observed UMBH mass function at $z = 0$ is reproduced. This self-regulation limited growth is implemented by steepening the high luminosity end of the AGN LF at the bright end. We estimate the abundance of UMBHs to be $\\sim \\,7 \\times 10^{-7}\\, Mpc^{-3}$ at $z = 0$. The key prediction of our model is that the slope of the $M_{\\rm bh} - \\sigma$ relation likely evolves with redshift at the high mass end. Probing this is observationally challenging at the present time but there are several bright, massive early-type galaxies that are promising host candidates from the SDSS survey as well as a survey of bright central galaxies of nearby clusters. Observational detection of UMBHs will provide key insights into the physics of galaxy formation and black hole assembly in the Universe."
    },
    "0808/0808.2214_arXiv.txt": {
        "abstract": "We present an analysis of XMM-Newton and RXTE data from three observations of the neutron star LMXB 4U~1636-536. The X-ray spectra show clear evidence of a broad, asymmetric iron emission line extending over the energy range 4--9 keV. The line profile is consistent with relativistically broadened Fe K-$\\alpha$ emission from the inner accretion disk. The Fe K-$\\alpha$ line in 4U~1636-536 is considerably broader than the asymmetric iron lines recently found in other neutron star LMXBs, which indicates a high disk inclination. We find evidence that the broad iron line feature is a combination of several K-$\\alpha$ lines from iron in different ionization states. ",
        "introduction": "Relativistically broadened, asymmetric Fe~K-$\\alpha$ lines from the inner accretion disk have been observed in many supermassive and stellar-mass black holes \\citep[e.g.][]{2006AN....327..943F}. In neutron star binaries, however, the iron lines are weaker, and until recently observations did not clearly reveal a relativistic line profile. \\citet{2007ApJ...664L.103B} found an asymmetric Fe~K-$\\alpha$ line in {\\it XMM-Newton} spectra of the low-mass X-ray binary (LMXB) Serpens X-1 and showed that the line profile is consistent with fluorescent iron line emission from the inner accretion disk. Similar asymmetric Fe~K-$\\alpha$ lines were found by \\citet{2008ApJ...674..415C} in {\\it Suzaku} spectra of the neutron star LMXBs Serpens X-1, 4U~1820-30, and GX~349+2. In this paper we present {\\it XMM-Newton} and {\\it RXTE} observations of the neutron star LMXB 4U~1636-536. We analyze the X-ray spectra to determine the profile of the relativistic Fe~K-$\\alpha$ line and constrain the properties of the inner accretion disk.   4U~1636-536 (V801~Ara) is a well-studied, bursting LMXB consisting of a neutron star in a 3.8-hr orbit with a 0.4 solar mass, 18th magnitude star \\citep{1990A&A...234..181V} and is located at a distance of $\\sim$6~kpc \\citep{2006ApJ...639.1033G}. The X-ray timing properties of the binary have been studied extensively. \\citet{1996ApJ...469L..17Z} and \\citet{1997ApJ...479L.141W} discovered quasi-periodic oscillations (QPOs) at kHz frequencies. The source also exhibits highly coherent burst oscillations at 581~Hz which are likely related to the rotation of the neutron star \\citep{1997IAUC.6541....1Z,2002ApJ...577..337S}. \\citet{1999ApJ...514L..31K} found that the soft X-ray emission, modulated at the kHz QPO frequency, lags behind the hard X-ray emission. A possible explanation for this phase lag is the reprocessing of hard X-rays in a cooler Comptonizing corona with a size of at most a few kilometers. \\citet{2007MNRAS.376.1139B} interpreted observations showing a decline of the QPO coherence and rms amplitude at high QPO frequencies as evidence that the inner radius of the accretion disk in 4U~1636-536 is usually larger than but sometimes approaches the innermost stable circular orbit (ISCO). However, \\citet{2006MNRAS.371.1925M} argued that this decline is not caused by effects related to the ISCO.    In this paper we report on three X-ray observations of 4U~1636-536 that were carried out simultaneously with {\\it XMM-Newton} and {\\it RXTE}. We present an analysis of the X-ray spectrum over the 0.5--100~keV energy range and our results from modeling the continuum and relativistic Fe~K-$\\alpha$ line emission. We use the Fe~K-$\\alpha$ line profile to derive constraints on the disk inclination and inner disk radius. We show that the line profile is not consistent with a single relativistic Fe~K-$\\alpha$ line from a neutron star accretion disk and that at least two lines from different ionization states of iron are needed to adequately describe the line profile. Finally, we discuss the implications of our findings for measurements of neutron star radii based on Fe~K-$\\alpha$ line profiles. ",
        "conclusions": "We have analyzed X-ray spectra of the neutron star LMXB 4U~1636-536 obtained with {\\it XMM-Newton} and {\\it RXTE} in 2005, 2007, and 2008. The very high signal-to-noise ratio of the spectra allowed us to clearly detect a broad, relativistic Fe~K-$\\alpha$ line from the inner accretion disk. The line is significantly broader than the asymmetric Fe~K-$\\alpha$ lines recently found in other neutron star LMXBs \\citep{2007ApJ...664L.103B,2008ApJ...674..415C}. The broader line profile is likely the result of a high disk inclination in 4U~1636-536. The inclination angles derived from iron line profiles in other neutron star LMXBs have so far been comparatively low. As pointed out by \\citet{2008ApJ...674..415C}, this is likely a selection effect because the narrower lines in low inclination systems are more easily detectable. With the high signal-to-noise ratio of the 4U~1636-536 spectra, we have now been able to measure the broader line profile in a high inclination LMXB. Our analysis of the Fe~K-$\\alpha$ line profile places a lower limit of $64^\\circ$ on the disk inclination in 4U~1636-536. This limit is consistent with the 36--74$^\\circ$ constraint on the orbital inclination by \\citet{2006MNRAS.373.1235C} and with the non-detection of eclipses.    When fitting the iron line profile with a single relativistic line component, we find a significant difference in the rest-frame line energy between the three observations. This difference is likely caused by a change in the ionization profile of the disk related to the change in X-ray luminosity between the three observations. The line energy derived for the second observation is close to 6.97~keV, the K-$\\alpha$ line energy of Fe~XXVI, which suggests that Fe~XXVI contributes significantly to the observed line profile. According to our fit with a disk-blackbody model, the highest temperature in the disk is $\\sim$0.9~keV. Because plasma in thermal equilibrium at this temperature does not contain a significant fraction of Fe~XXVI, is it evident that the ionization profile in the disk is strongly affected by photoionization.    We find that the Fe~K-$\\alpha$ line profile for two of the observations is too broad to be adequately described by a single relativistic emission line with physically reasonable values of disk inclination and inner disk radius. The broader than expected line profile can be explained by overlapping K-$\\alpha$ lines from iron in different ionization states. The presence of multiple ionization states is also indicated by the difference in the fitted line energy between the three observations. It is evident that multiple line components need to be considered to adequately model the relativistic iron line profiles in neutron star LMXBs. We obtained reasonable line parameters when fitting two iron lines with different rest-frame energies. This is obviously an oversimplification, since many ionization states are probably contributing to the iron line profile. However, because the lines are broad and overlap considerably, it is not possible to constrain the parameters of all line components independently. Even with only two line components, the fit parameters are already strongly correlated, leading to larger parameter uncertainties than a fit with a single line component. In order to adequately model the contribution from multiple ionization states when fitting relativistic Fe~K-$\\alpha$ line profiles, improved models are needed that can predict the ionization profile in the accretion disk.    The Fe~K-$\\alpha$ line profiles in neutron star LMXBs can in principle be used to place upper limits on the neutron star radius by constraining the inner disk radius \\citep{2008ApJ...674..415C}. Previously, these line profiles have been fitted with only a single line component. However, if more than one ionization state contributes significantly to the Fe~K-$\\alpha$ emission, the line profile will be broader than for a single relativistic line, and a fit with a single line model may underestimate the inner disk radius and thus the limit on the neutron star radius. This can be clearly seen for two of our observations for which the upper limit on $R_{in}$ increases from $6.3R_g$ to $11.9R_g$ and $10.9R_g$, respectively, when the iron line profile is fitted with two line components instead of a single line component. In contrast, the upper limit on $R_{in}$ for the second observation decreases from $13.3R_g$ to $9.8R_g$. It is evident that the constraints on the inner disk radius strongly depend on the assumptions made about the contributing ionization states of iron and that a better understanding of the ionization profile in the disk is needed to obtain reliable limits on the neutron star radius. The contribution of multiple ionization states may also be important for the interpretation of some iron line profiles in accreting black holes. We note that the disk inclination and inner disk radius are strongly anti-correlated when fitting broad iron line profiles, which can lead to large uncertainties of the two parameters. Prior knowledge of the disk inclination can significantly reduce the uncertainty of the inner disk radius.    It was shown by \\citet{2007ApJ...656.1056L} that broad iron line features can also be produced by Compton scattering of line photons in a strong outflow. The expected line profiles for this process are generally characterized by a narrow line at $\\sim$6.4--6.6~keV from fluorescence in the outflow and a broad, redshifted component below 7~keV from downscattering of the line photons. These line profiles differ qualitatively from those found in 4U~1636-536 which show significant emission above 7~keV and no narrow line component (Figure \\ref{spec2}). We also note that Compton scattering in an outflow only contributes significantly to the iron line emission if the optical depth is at least of order unity, which requires a mass outflow rate on the order of the Eddington mass accretion rate. The X-ray luminosity we observed in 4U~1636-536 was only $\\sim$5\\% of the Eddington luminosity, which suggests that the rate of any outflow in 4U~1636-536 was significantly below the Eddington mass accretion rate. It therefore seems unlikely that a large fraction of the observed iron line emission was produced in an outflow."
    },
    "0808/0808.3815_arXiv.txt": {
        "abstract": "% NGC\\,2362 is a richly populated Galactic cluster, devoid of natal molecular gas and dust. The cluster represents the final product of the star forming process and hosts an unobscured and near-complete initial mass function. NGC\\,2362 is dominated by the O9 Ib multiple star, $\\tau$ CMa, as well as several dozen unevolved B-type stars. Distributed throughout the cluster are several hundred suspected intermediate and low-mass pre-main sequence members. Various post-main sequence evolutionary models have been used to infer an age of $\\sim$5 Myr for the one evolved member, $\\tau$ CMa. These estimates are in close agreement with the ages derived by fitting pre-main sequence isochrones to the contracting, low-mass stellar population of the cluster. The extremely narrow sequence of stars, which extends more than 9 mag in the optical color-magnitude diagram, suggests that star formation within the cluster occurred rapidly and coevally across the full mass spectrum. Ground-based near infrared and H$\\alpha$ emission surveys of NGC\\,2362 concluded that most ($\\sim$90\\%) of the low-mass members have already dissipated their optically-thick, inner ($\\ll$1 AU) circumstellar disks. {\\it Spitzer} IRAC observations of the cluster have confirmed these results, placing an upper limit on the primordial, optically thick disk fraction of the cluster at $\\sim$7$\\pm$2\\%. The presence of circumstellar disks among candidate members of NGC\\,2362 is also strongly mass-dependent, such that no stars more massive than $\\sim$1.2 M$_{\\odot}$ exhibit significant infrared excess shortward of 8 $\\mu$m. NGC\\,2362 will likely remain a favored target of ground-based and space-based observations. Its well-defined upper main sequence, large population of low-mass, pre-main sequence stars, and the narrow age spread evident in the color-magnitude diagram ensure its role as a standard model of cluster as well as stellar evolution. ",
        "introduction": "The young cluster NGC\\,2362 in Canis Majoris (CMa) is dominated by the 4$^{th}$ mag O9 Ib multiple star, $\\tau$ CMa and several dozen B-type stars (McSwain \\& Gies 2005), spherically distributed within a volume $\\sim$3 pc in radius. Shown in Figure~1 is a 15\\arcmin$\\times$15\\arcmin\\ Second Palomar Observatory Sky Survey (POSS-II) red image of NGC\\,2362 obtained from the Digitized Sky Survey. The cluster is free of molecular gas and nebular emission and suffers very little interstellar reddening despite an accepted distance of nearly 1.5 kpc.  With an age of $\\sim$5 Myr, only $\\tau$ CMa has evolved significantly away from the cluster zero-age main sequence (ZAMS). Although probably relaxed, the evaporation timescale for the cluster is significantly greater than its age, implying that few members have dispersed. In essence, an unobscured and near-complete stellar population remains around $\\tau$ CMa, making NGC\\,2362 an ideal target for initial mass function (IMF) studies (Moitinho et al. 2001). The rich history of the cluster begins with its discovery by Fr. Giovanni Battista Hodierna in Sicily during the 17$^{th}$ century using a Galilean-type refractor.  His observations of nebulous objects which includes NGC\\,2362 were published in 1654. With only one member ($\\tau$ CMa) visible to the unaided eye, the cluster quickly returned to obscurity for over a century before Sir William Herschel noted its presence, entry H VII.17 in his catalog of stellar clusters and nebulae, one of the pre-cursors to Dreyer's New General Catalog (NGC). Dreyer's notes within the original NGC summarize NGC\\,2362 as a pretty large, rich cluster centered upon 30 CMa ($\\tau$). The compact nature of NGC\\,2362 is quite striking when viewed on the Palomar Observatory Sky Survey (POSS) plates. The large number of early-type stars form a luminous halo around $\\tau$ CMa, $\\sim$10\\arcmin\\ in diameter (Figure~\\ref{f1}).  \\begin{figure}[h!] \\hspace{-1.0cm} \\plotfiddle{f1.ps}{11.6cm}{90.0}{58.0}{58.0}{70.0}{-15.0} \\caption[f1.ps]{A 15\\arcmin$\\times$15\\arcmin\\ Second Palomar Observatory Sky Survey (POSS-II) red image of NGC\\,2362 obtained from the Digitized Sky Survey. $\\tau$ CMa, the 4$^{th}$ mag O9 Ib star (center), is the most massive cluster member and the only star that has evolved significantly away from the ZAMS. Given the abundance of unevolved B-type members from B1 to B9, the early population of NGC\\,2362 has often been used to define the upper section of the empirically-derived ZAMS. The pre-main sequence population of NGC\\,2362 is symmetrically distributed around $\\tau$ CMa.\\label{f1}} \\end{figure}    Although lacking nebulosity in the immediate cluster vicinity, just over one degree east of NGC\\,2362 extensive H~II emission is apparent on the POSS plates. This emission is part of the giant H~II region Sharpless 310, described by Sharpless (1959) as an incomplete ring 8$^{\\circ}$ in diameter.  The ionizing sources for this gas are most likely 29 CMa, $\\tau$ CMa, and the early B-type members of NGC\\,2362. IRAS images of the NGC\\,2362 region reveal that the cluster is located within an evacuated cavity approximately 30\\arcmin\\ in diameter. A partial ring of dust emission, most prominent at 60~$\\mu$m, is evident to the east. The nearby dark nebula, L1667, also believed to be 1.5~kpc distant (Lada \\& Reid 1978), lies southeast of NGC\\,2362, near the variable M5 supergiant VY~CMa. Tenuous dark clouds appear to follow the contours of H~II emission on the POSS plates, perhaps remnants of the molecular cloud complex from which NGC\\,2362 formed. The OB stellar population of NGC\\,2362 may have triggered a second generation of star formation within L1660 to the northeast, which hosts the Herbig-Haro object HH~72 (Reipurth \\& Graham 1988).  Shown in Figure~2 is a reproduction of a Schmidt plate from Reipurth \\& Graham (1988) centered upon L1660 and with the locations of HH~72 and three known H$\\alpha$ emission stars identified. Reipurth \\& Graham (1988) conclude that L1660 is slowly being eroded away by intense UV radiation from the OB stellar population of NGC\\,2362 and other nearby massive stars.   \\begin{figure} \\centering \\includegraphics[angle=0,width=0.85\\textwidth]{f2.eps} \\caption[f2.eps]{The L1660 dark cloud and HH 72 on the SERC-J Schmidt plate from Reipurth \\& Graham (1988).  The field shown is approximately 13\\arcmin$\\times$17\\arcmin. The locations of three H$\\alpha$ emission stars are also identified, one of which (labeled 1) is possibly a HAeBe star. The other two H$\\alpha$ emitters are candidate low-mass T Tauri stars. \\label{f2}} \\end{figure} ",
        "conclusions": "Over the last half-century, NGC\\,2362 has played an extraordinary role in our understanding of the star formation process. Early perceptions of an anomalous mass function were overturned with the advent of modern detectors. Deep CCD surveys of the cluster by Wilner \\& Lada (1991) and Moitinho et al. (2001) revealed a well-populated, pre-main sequence extending more than 9~mag to nearly the substellar limit.  There are new indications, however, that a deficit of low-mass stars is present within the cluster, implying that the IMF issue has yet to be fully resolved.  Recent {\\it Chandra} X-ray observations of NGC\\,2362 have added several hundred more pre-main sequence candidates to the $\\sim$100$+$ suspected members exhibiting H$\\alpha$ emission or strong Li~I $\\lambda$6708 absorption. Deeper optical and infrared surveys of the cluster will also push the source detection threshold into the brown dwarf regime, permitting a closer examination of the cluster IMF. Interest in NGC\\,2362 has also shifted to the remaining optically thick circumstellar disks around low-mass cluster members. Near infrared and H$\\alpha$ emission surveys suggest that the inner disk regions have dissipated for most ($\\sim$90\\%) of the suspected cluster members as evidenced by the decay of near infrared excess and strong H$\\alpha$ emission. {\\it Spitzer} observations are beginning to resolve the remaining questions of disk frequency within the cluster. It is perhaps somewhat ironic that in the process of identifying and characterizing the low-mass population of NGC\\,2362 over the last decade, the OB stars have been somewhat neglected. McSwain \\& Gies' (2005) recent Str\\\"omgren photometric survey of 41 OB-type stars in the cluster region found only one candidate classical Be star. A more thorough spectroscopic analysis of the B-star population is needed to confirm spectral types, evaluate membership, and to examine questions of binarity, critical to understanding the placement of stars on the ZAMS. NGC\\,2362 will likely remain a favored target for ground-based and space-based observations. Its large, statistically-significant population of low-mass, pre-main sequence stars, its well-defined upper main sequence, compact structure, and lack of circumstellar and interstellar gas and dust relative to similarly aged clusters all contribute to the cluster's unique nature.  \\vspace{0.5cm}  {\\bf Acknowledgments.}  I wish to thank the referee for this paper, Francesco Damiani, and the editor, Bo Reipurth, for many helpful comments that significantly improved the manuscript.  SED is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-0502381.  The Digitized Sky Surveys were produced at the Space Telescope Science Institute under U.S. Government grant NAG W-2166. The images of these surveys are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the UK Schmidt Telescope.  The plates were processed into the present compressed digital form with the permission of these institutions. The Second Palomar Observatory Sky Survey (POSS-II) was made by the California Institute of Technology with funds from the National Science Foundation, the National Geographic Society, the Sloan Foundation, the Samuel Oschin Foundation, and the Eastman Kodak Corporation."
    },
    "0808/0808.3448_arXiv.txt": {
        "abstract": "This paper reports on the results of a numerical investigation designed to address how the initially anisotropic appearance of a GRB remnant is modified by the character of the circumburst medium and by the possible presence of an accompanying supernova (SN). Axisymmetric hydrodynamical calculations of light, impulsive jets propagating in both uniform and inhomogeneous external media are presented, which show that the resulting dynamics of their remnants since the onset of the non-relativistic phase is different from the standard self-similar solutions.  Because massive star progenitors are expected to have their close-in surroundings modified by the progenitor winds, we consider both free winds and shocked winds as possible external media for GRB remnant evolution. Abundant confirmation is provided here of the important notion that the morphology and visibility of GRB remnants are determined largely by their circumstellar environments. For this reason, their detectability is highly biased in favor of those with massive star progenitors; although, in this class of models, the beamed component may be difficult to identify because the GRB ejecta is eventually swept up by the accompanying SN. The number density of asymmetric GRB remnants in the local Universe could be, however, far larger if they expand in a tenuous interstellar medium, as expected for some short GRB progenitor models. In these sources, the late size of the observable, asymmetric remnant could extend over a wide, possibly resolvable angle and may be easier to constrain directly. ",
        "introduction": "\\label{int} Relativistic jets are common in the astrophysical environment. Objects known or suspected to produce them include radio galaxies and quasars \\citep{b84}, microquasars \\citep{mr99} and gamma-ray bursts \\citep{g06}. An important difference between jets of gamma-ray bursts (GRBs) and the better studied radio jets of quasars or microquasars is that active quasars often inject energy over extended periods of time into the jet while GRB sources are impulsive.  Although quasar jets remain highly collimated throughout their lifetimes, GRB jets decelerate and expand significantly once they become nonrelativistic. Expansion into a uniform medium has been well studied \\citep{ap01}, but the interaction of a GRB remnant with a nonuniform medium remains poorly understood.  Much of our effort in this paper is therefore dedicated to determining how the morphology and dynamics of young GRB remnants is modified by the character of the circumburst medium.  Some of the questions at the forefront of attention include the effects of the external medium and the degree to which GRB remnant dynamics and structures are modified by the presence of an accompanying supernova. We address both of these issues here. Because massive stars are expected to have their close-in surroundings modified by the progenitor winds, we consider both free winds and shocked winds as possible surrounding media for the GRB remnant evolution. Detailed hydrodynamic simulations of this interaction are presented in \\S\\S \\ref{wind} and \\ref{bubble}, while a brief description of the numerical methods and the initial models is giving in \\S \\ref{nm}. For completeness, the interaction with a constant-density medium is discussed in \\S \\ref{const}. The role of supernova explosions in shaping the evolution and morphology of GRB remnants is discussed in \\S \\ref{sn}.  The effects of a nonspherical circumburst medium are briefly addressed in \\S \\ref{wind}. Discussion and conclusions are presented in \\S \\ref{dis}. ",
        "conclusions": ""
    },
    "0808/0808.0577_arXiv.txt": {
        "abstract": "{} {The stability of dissipative Taylor-Couette flows with an axial stable density stratification and a prescribed  azimuthal magnetic field is considered. } { Global nonaxisymmetric solutions of the linearized MHD equations with toroidal magnetic field, axial density stratification and differential rotation are found for both insulating and conducting cylinder walls.} {Flat rotation laws such as the quasi-Kepler  law are unstable against the nonaxisymmetric stratorotational instability (SRI). The influence of a current-free toroidal magnetic field   depends on the magnetic Prandtl number Pm: SRI is supported  by $\\rm Pm > 1$ and it is suppressed by  $\\rm Pm \\lsim 1$.  For too flat rotation laws a smooth transition exists to the instability  which the toroidal magnetic field produces in combination with the differential rotation. This nonaxisymmetric azimuthal magnetorotational instability (AMRI) has been computed   under the presence of    an axial density gradient.  If the magnetic field between the cylinders is not current-free then also the  Tayler instability  occurs and the transition from the hydrodynamic SRI to the magnetic Tayler instability  proves to be rather complex. Most spectacular is the `ballooning' of the stability domain by the density stratification: already a rather small rotation stabilizes  magnetic fields against the Tayler instability.   An  azimuthal component of the resulting  electromotive force    only exists for density-stratified flows. The related  alpha-effect  for magnetic SRI of  Kepler rotation appears   to be positive for negative $d\\rho/dz <0$. } {} ",
        "introduction": " ",
        "conclusions": "The stability of the dissipative Taylor-Couette flow under the joint influence of a stable vertical density stratification and an azimuthal magnetic field is considered. The problem is of interest for future laboratory experiments but also within the frame of  accretion disk physics. The Kepler rotation generates strong toroidal magnetic fields   dominating the poloidal components. The standard MRI which works with only axial fields may be of minor relevance for the stability of the Kepler rotation law compared with the azimuthal MRI in connection with the density stratification and Tayler instability. Mainly nonaxisymmetric `kink' modes ($m=1$) are considered but there are also examples where the axisymmetric modes   are the most unstable ones.  We started with a    discussion of  the SRI  without magnetic fields. For a flat rotation law a  stratification value Fr exists for  which the critical Reynolds number of the rotation has a minimum (see Fig. \\ref{NN}). The SRI is basically stabilized by both too weak or  too strong stratification. The instability only appears if the characteristic buoyancy time approaches the rotation period. Also if the rotation is too flat the instability disappears.   The limiting ratio $\\mu_\\Omega$  strongly depends on the gap width (see Fig. \\ref{gaps}).  For small gaps  rotation laws with $\\mu_\\Omega>\\hat\\eta$ even prove to be unstable while for wide  gaps the condition  $\\mu_\\Omega<\\hat\\eta$ results for exciting SRI. Our previous finding that $\\mu_\\Omega \\lsim \\hat\\eta$ limits the SRI is reproduced for medium-sized gaps. In all cases, however, the quasi-Kepler rotation proves to be unstable.  New experiments with different gap sizes and larger Reynolds numbers could  verify these results.  Figures~\\ref{amri} and \\ref{amri1} yield the basic results for SRI subject to toroidal fields. The magnetic field is assumed as current-free in the fluid between the cylinders, i.e. $B_\\phi \\propto 1/R$ excluding  Tayler instability. Without density  stratification  no Rayleigh instability exists for quasi-Kepler rotation but nonaxisymmetric modes with $m=1$ are unstable as a result of the interaction of differential rotation and magnetic field (`AMRI').  The magnetic influence strongly depends on the magnetic Prandtl number. The instability needs higher Reynolds number  for ${\\rm Pm}\\lsim 1$ and it needs lower Reynolds number  for ${\\rm Pm}\\gsim 1$. After  our experiences with MHD instabilities this is not a surprise. It is a surprise, however, that always the magnetic influence is only weak.  Up to Hartmann numbers of ${\\rm Ha}\\simeq 100$ only a magnetic-induced factor of two plays a role\\footnote{${\\rm Ha}=100$ for gallium corresponds to $B\\approx 2200/R_0\\ [{\\rm Gauss}]$ with $R_0$ in cm}. Hence, our conclusion is that the SRI  survives for rather high magnetic fields. For large Pm it is even supported by the toroidal magnetic field.   The combination of rotation, density stratification and magnetic field leads to complex results.  However, a basic observation is that the critical Reynolds numbers, if  existing, above which the flow becomes unstable without a magnetic field, are increased by the stable density stratification. In contrast, the critical Hartmann numbers, if  existing by the Tayler instability, above which the field becomes unstable without a rotation, do not depend on the stratification. Different transition are possible between the two limits.  Generally speaking, the  density stratification `balloons' the stability region and in this sense it stabilizes the flow. For slow rotation  the maximal stable  magnetic field exceeds the critical magnetic field without rotation while for faster rotation the   maximal stable  magnetic field is smaller than this critical value. Even rather slow values of the Reynolds numbers lead to a stabilization of those fields which are unstable for $\\rm Re=0$. The effect is strong for $\\rm Pm=1$ but it becomes smaller for decreasing magnetic Prandtl numbers.  For steep radial profiles of the magnetic field (i.e. strong axial currents) and magnetic Prandtl numbers $\\rm Pm \\gsim 1$ one also finds  the magnetic field destabilizing the Rayleigh instability, i.e. the critical  Reynolds numbers with magnetic field are lower than the Reynolds numbers without magnetic field.  This magnetic destabilization only exists for not too small magnetic Prandtl numbers. It  exists for both uniform and density-stratified fluids (Fig. \\ref{gen}).  Finally, the $\\phi$-component of the electromotive force representing  the $\\alpha$-effect of the mean-field dynamo theory has been computed. The computations require an extreme degree of accuracy. The results demonstrate the importance of the density stratification for the existence of the $\\alpha$-effect. Without density stratification the correlations are vanishing in the radial average. With an axial density stratification  the calculations model the polar region of a rotating sphere or disk. This interpretation accepted we found the  $\\alpha_{\\phi\\phi}$ at the northern pole as {\\em positive} (Fig. \\ref{emf}). Again, the basic ingredient of this $\\alpha$-effect model is the density stratification."
    },
    "0808/0808.2744_arXiv.txt": {
        "abstract": "We present 3D radiation-gasdynamical simulations of an ionization front running into a dense clump. In our setup, a B0 star irradiates an overdensity which is at a distance of $10\\,\\mathrm{pc}$ and modelled as a supercritical $100\\,\\msol$ Bonnor-Ebert sphere. The radiation from the star heats up the gas and creates a shock front that expands into the interstellar medium. The shock compresses the clump material while the ionizing radiation heats it up. The outcome of this ``cloud-crushing'' process is a fully turbulent gas in the wake of the clump. In the end, the clump entirely dissolves. We propose that this mechanism is very efficient in creating short-living supersonic turbulence in the vicinity of massive stars.  \\noindent \\pacs{47.40.Nm,97.10.Bt,98.38.Am} ",
        "introduction": "Massive stars strongly influence the environment in which they have formed by stellar winds, ionizing radiation and supernova explosions. An ionization front which expands into the ambient interstellar medium (ISM) and hits a dense clump may compress it so heavily that gravitational collapse might be triggered. On the other hand, ionization heats up the material and can photoevaporate the clump. The remaining material of this competition may form low-mass stars~\\cite{hest05} or brown dwarfs~\\cite{whit04}.  The interaction of a shock with a dense clump is called ``cloud-crushing''. The cloud-crushing scenario has been studied numerically both for supernova shocks and ionization fronts. The first extensive studies of the fate of the shocked cloud already showed that strong vortex rings can be produced~\\cite{klein94}. The mixing properties of the cloud depend sensitively on the initial density distribution~\\cite{naka06}. Furthermore, simulations of dense clumps exposed to an ionizing flux but without strong shocks show the generation of kinetic energy~\\cite{kessel03} and fragmentation of the clump~\\cite{esquivel07}. Radiation-gasdynamical simulations have also been used to match observations in H II regions, especially the Eagle Nebula~\\cite{williams01, miao06}.  Although ionizing radiation injects a significant amount of energy into the ISM, it does not seem to be an important driving mechanism of interstellar turbulence on a global scale~\\cite{maclow04}. However, the cloud-crushing process generates a considerable amount of turbulence locally in the wake of the cloud. We find that the motion of the cloud material is mostly supersonic while the ambient gas behind the front moves only subsonically. The continuous heating limits the lifetime of the dense material, but the supersonic motions are maintained until the cloud disperses. This is contrary to the situation in jet-clump interactions, where the situation is less clear with some studies showing mostly subsonic motions~\\cite{baner07} while others claim supersonic velocity fields~\\cite{li07}. ",
        "conclusions": "We have seen that cloud-crushing by ionization fronts can lead to short-living supersonic turbulence. Altough only a minute fraction of the energy input is converted into kinetic energy, up to $60\\%$ of the affected gas is supersonic. While it is mainly the cold gas that is highly supersonic, it is the hot gas that moves the fastest. The bulk motion of the shock is an important contribution to the supersonic flow, since the transversal fluctuations are at best slightly supersonic."
    },
    "0808/0808.3118_arXiv.txt": {
        "abstract": "We have developed a new two-dimensional hydrostatically-balanced isobaric hydrodynamic model for use in simulation of exoplanetary atmospheres.  We apply this model to the infrared photosphere of the hot Jupiter HD 189733 b, for which an excellent 8-$\\mu$m light curve has been obtained.  For reasonable parameter choices, the results of our model are consistent with these observations.  In our simulations, strongly turbulent supersonic flow develops, with wind speeds of approximately 5~km~s$^{-1}$.  This flow geometry causes chaotic variation of the temperature distribution, leading to observable variations in the light curve from one orbit to the next. ",
        "introduction": "Since the earliest dynamical models of strongly irradiated exoplanetary atmospheres were published \\citep{sho02,cho03}, it has been apparent that the expected large temperature gradients and the resulting high wind speeds could have such a large impact on these planets' appearance that the effects could be observable even at a distance of dozens of parsecs.  As the models have grown increasingly sophisticated over the past six years, research has progressed along two distinct lines.  \\citet{sho02}, \\citet{coo05}, \\citet{dob08} and \\citet{sho08} have produced three-dimensional models, in an effort to simulate the greatest possible range of relevant physical processes.  Because of the high computational costs of such models, they have been run at comparatively low resolution.  In contrast to this approach, \\citet{cho03}, \\citet{lan07}, \\citet{lan08}, and \\citet{lan08b} have chosen to employ two-dimensional models which can be run at higher resolutions.  On highly irradiated planets, the radiative zone is believed to extend deep into the atmosphere, to a pressure depth of hundreds of bars \\citep{coo05, iro05, sho08}.  The flow is therefore expected to be strongly stratified, with vertical motion comparatively unimportant.  The assumption inherent in a two-dimensional model is that the motion at small scales which can be captured due to the finer grid spacing is more important than the vertical flow which a two-dimensional model must neglect.  Nevertheless, a serious concern for two-dimensional models is the possible existence of crucial three-dimensional processes which do not require strong vertical motion in order to become significant.  Furthermore, previous attempts to develop two-dimensional models have been hampered by questions regarding the validity of the physics involved.  \\citet{lan07} employ a shallow-water model which is at best a first-order approximation to realistic atmospheric dynamics.  While the barotropic-equivalent model used by \\citet{cho03}, formally very similar to the shallow-water equations, is on a firmer physical footing, there is some indication that the model becomes numerically unstable at high wind speeds \\citep{rau07}.  In an attempt to achieve greater realism than is possible using a shallow-water model, \\citet{lan08} developed a model employing fully-compressible two-dimensional hydrodynamics.  As we will show in \\S2, however, such models produce features which necessarily violate hydrostatic balance.  In this paper, then, we present a two-dimensional model which must maintain hydrostatic equilibrium. It is our hope that this new model will be able to approach the rigor of a three-dimensional model, while maintaining the adaptability and speed of a two-dimensional model.  While these modeling efforts have been underway since 2002, the ability to constrain these models using observations is a more recent development.  Of particular interest, \\citet{knu07} have produced a map of the planet's longitudinal temperature variation based on their \\textit{Spitzer} observations of its flux in the 8-$\\mu$m band.  From the depth of the secondary eclipse, they find a hemispherically-averaged day-side brightness temperature of $1205.1 \\pm 9.3$ K, while they estimate a cooler hemispherically-averaged night-side brightness temperature of $973 \\pm 33$ K, based on the flux curve.  Interestingly, the hottest part of the planet is \\emph{not} directly beneath the star: the temperature maximum is offset some $30^\\circ$ east of the substellar point.  Perhaps even more interestingly, the temperature minimum is offset $30^\\circ$ \\emph{west} of the antistellar point.  It is clear, then, the temperature distribution is strongly influenced by planetary winds; at first glance it would appear that the flow is eastward on the day-side, while westward on the night-side.  To date, models have not reproduced this flow geometry \\citep{cho03, coo05, lan07, cho08, dob08, lan08, sho08}.  Additionally, these observations do not seem to be generally applicable to hot Jupiters.  \\citet{har06}, for example, observed the 24-$\\mu$m flux of the giant exoplanet $\\upsilon$ Andromedae b, finding a much larger temperature contrast between the illuminated and dark hemispheres.  Additionally, they found a slight \\emph{westward} displacement of the hot spot from the substellar point, although their observations did not exclude the possibility that there was no offset at all.  The mid-infrared observations of HD 179949 b by \\citet{cow07} also appear to be consistent with zero phase offset.  It is not clear to what extent these discrepancies result from different dynamics on the planets themselves, and to what extent they are caused by the different pressure depths under observation.  In other respects, HD 189733 b remains a fairly typical hot Jupiter; its period, however, is shorter than most, with $P=2.218573 \\pm 0.000020$ d.  The transit has allowed determinations of several other parameters: $R=1.154 \\pm 0.032 R_J$, $i=85.79^\\circ \\pm 0.24^\\circ$ \\citep{bak06}, $a=0.0313 \\pm 0.0004$, and $M=1.15 \\pm 0.04 M_J$ \\citep{bou05}.  The orbit is assumed to be circular.  HD 189733 b has not received the attention from modelers which has been bestowed on its more famous cousin, HD 209458 b.  Nevertheless, the well-resolved flux curve produced by \\citet{knu07} has made it a most interesting target for simulation, and \\citet{sho08} apply their atmospheric model to HD 189733 b, with mixed success.  They recover the correct eastward phase offset for the hot spot, but are unable to produce a flow pattern which causes the observed westward offset for the cold spot from the antistellar point.  Their flow patterns in the upper atmosphere are in general supersonic, with very low pressures characterized by longitudinal and latitudinal flow from the day side towards the night side, while deeper layers are dominated by a supersonic eastward jet at the equator.  Interestingly, despite the high wind speeds ($|\\mathbf{v}| \\gtrsim 3$ km/s) and the resulting large wind shear, no turbulence develops in their simulations.  As we shall see, this does not match the results of the two-dimensional model presented in this paper.  It is possible that their relatively low horizontal resolution (144x90) combined with the finite-difference differentiation employed by the ARIES/GEOS dynamical core used in their model conspire to produce sufficient numerical dissipation to prevent the development of turbulence; it is also possible (and possibly more likely) that the development of turbulence is prevented by three-dimensional effects which our two-dimensional model is unable to capture.  An analytical study to constrain the conditions under which turbulence is expected to arise would surely be profitable; however, this, and indeed, the more general question of the conditions necessary for planetary-scale turbulent flow, must remain a topic for future investigation.  In any case, it is clear that the \\citet{knu07} time-series is of sufficiently high quality to provide an excellent benchmark for both existing and future simulations of exoplanetary atmospheres.  The ARIES/GEOS core used in \\citet{coo05} and \\citet{sho08} has been applied to the terrestrial atmospheres of Earth and Mars \\citep{sho08}, and the equivalent barotropic formulation of \\citet{cho03} -- also used in \\citet{rau07} -- has enjoyed some success in reproducing the primary features of Jupiter's dynamics.  However, the conditions on many extrasolar planets are so utterly unlike anything seen in our solar system that it is advisable to include data from exoplanet observations when testing the models.  This paper is organized as follows: In \\S 2, we derive in some detail the hydrodynamical core of our model, as well as providing a treatment of the radiative forcing scheme.  In \\S 3, we apply our model to the atmosphere of HD 189733 b, comparing our results to those of \\citet{sho08} and to the data obtained by \\citet{knu07}.  We conclude in \\S4. ",
        "conclusions": "While this model has not yet provided an optimized fit to the light-curve of HD 189733 b obtained by \\citet{knu07}, it does provide a reasonable explanation for a dynamical configuration which could give rise to the observed light curve.  In general, other models have been able to reproduce the observed phase offset of the flux maximum, corresponding to an eastward shift of the hottest temperatures from the substellar point \\citep{coo05, sho08}.  It has been more difficult to account for the apparent westward shift of the coldest temperatures from the anti-stellar point.  However, the model proposed in this paper offers a physically reasonable mechanism by which the observed light curve may arise: at low pressures $p<200$ mbar, the flow tends to run from the substellar point to the antistellar point at supersonic speeds.  This results in a collision between eastward winds and westward winds, which produces chaotic, turbulent flow on a large scale.  As a result, there is no steady-state temperature distribution; the enormous winds cause unpredictable shifts in the temperature distribution which, in turn, alter the position of the flux minima and maxima from orbit to orbit.  The phase offset of the flux minimum observed by \\citet{knu07} is not inconsistent with these results.  Furthermore, this hypothesis is readily testable: turbulence on the scale predicted by our model is expected to induce shifts of several hours or more in the timing of the flux minima and flux maxima; the light-curve obtained by \\citet{knu07} indicates that it is possible to measure the timing of these flux extrema with an error of $\\lesssim \\pm 1$ hour.  Therefore, if subsequent observations of HD 189733 b showed significant variations in the infrared light curve from one orbit to the next, the presence of significant large-scale turbulence would be strongly supported.  Conversely, the absence of such variations would imply that turbulence on a scale necessary to explain the \\citeauthor{knu07} results is suppressed by other factors for which this simple two-dimensional model is unable to account.  Some support for the model presented here may be found from the recent 24-$\\mu$m phase curve obtained by \\citet{knu08}.  These observations were taken over the same portion of the orbit as the earlier 8-$\\mu$m observations \\citep{knu07}.  In both cases, the flux maximum precedes secondary eclipse.  In contrast to the 8-$mu$m data, however, the 24-$\\mu$m curve increases monotonically throughout the observing window -- there is no flux minimum following the transit.  It is unclear whether this is due to a qualitatively different flow at the 24-$\\mu$m photosphere, or due to the type of turbulent variation suggested by the model presented in this paper.  It would be useful to obtain more observations in both the 8- and 24-$\\mu$m bands so that a comparison can be made between flows on multiple orbits, but at the same atmospheric depth.   In any case, it is quite clear that extensive further observations are necessary to obtain a detailed characterization of the atmospheric behavior of HD 189733 b.  The model presented here represents a significant forward step from previous two-dimensional models.  Although the number of free parameters in our treatment of the radiative forcing mitigates our success in fitting the \\citet{knu07} light curve, the quality of the fit is sufficient to conclude that the dynamics produced by our model offer a reasonable explanation for the observations, while a poor fit would imply that the simulated flows are excluded by the data.  We can therefore say with some confidence that our model is not excluded by the observations currently available.  It is also the first model -- in either two or three dimensions -- to provide an explanation for the unexpected phase offset of the flux minimum in the HD 189733 b 8-$\\mu$m curve.  Despite these encouraging results, much room for improvement exists.  A more sophisticated treatment of the radiative forcing is necessary.  However, true radiative transfer is difficult to approximate in a two-dimensional model; among other issues, the depth to which incident radiation penetrates depends rather strongly on the wavelength.  It is therefore likely that improvements in this area will require a shift to a fully three-dimensional code.  Furthermore, a three-dimensional treatment is necessary to ensure that departures from stratified flow do not significantly affect the results.  However, the ubiquity of turbulent flow obtained in two-dimensional simulations suggests that small-scale motion is non-negligible, so that the low horizontal resolutions seen current three-dimensional models are less than ideal.  An immediate goal is therefore the production of a high-resolution three-dimensional model, including a realistic multi-wavelength treatment of radiative transfer.  With the recent decision to fund a non-cryogenic \\textit{Spitzer} mission motivated at least in part by the promise of useful observations of exoplanets, the efforts of modelers to assist in the identification of interesting targets will be critical.  Whatever uncertainties still plague these efforts, it is certain that the influx of \\textit{Spitzer} data over the next few years should provide rigorous tests on both existing models and on those yet to be developed."
    },
    "0808/0808.3560_arXiv.txt": {
        "abstract": "{During the OGLE-2 operation, Soszynski et al.~(2003) found 3 LMC candidates for an RR Lyr-type component in an eclipsing binary system. Two of those have orbital periods that are too short to be physically plausible and hence have to be optical blends. For the third, \\targetnts, we developed a model of the binary that could host the observed RR Lyr star. After being granted HST/WFPC2 time, however, we were able to resolve 5 distinct sources within a 1.3'' region that is typical of OGLE resolution, proving that \\target is also an optical blend. Moreover, the putative eclipsing binary signature found in the OGLE data does not seem to correspond to a physically plausible system; the source is likely another background RR Lyr star. There are still no RR Lyr stars discovered so far in an eclipsing binary system.} ",
        "introduction": "Eclipsing binary systems (EBs) have long been recognized as one of the most astrophysically rewarding targets for the reliable determination of masses, radii, luminosities and other physical stellar properties. The precision of the derived parameters is typically better than a few percent; it enables the studies of stellar structure and evolution, and yields reliable distances that are independent of any calibrations or empirical relations. EBs thus serve as \\emph{astrophysical laboratories} suitable for the study of individual component properties. The importance of finding intrinsic variables in EBs is thus obvious: being able to obtain the fundamental properties of such stars provides improved theoretical understanding, better calibrations, and in some cases leads to improving the cosmological distance ladder \\citep{guinan1998,fitzpatrick2003}. Cepheids, for example, have been found in galactic EBs \\citep{freyhammer2005,antipin2007} as well as in the LMC \\citep{alcock2002,lepischak2004,guinan2005}; there are more than 30 known $\\delta$-Scuti type components in EBs \\citep[see e.g.][]{dallaporta2002,rodriguez2004,christiansen2007}, and several other types such as slowly pulsating B stars \\citep{pigulski2007,pilecki2007}. To date, though, no RR Lyr type star has been found in an EB. Recently, however, an accurate parallax and relative proper motion of RR Lyr itself was obtained with the HST Fine Guidance Sensor \\citep{benedict2002} that estimated its absolute magnitude of $M = 0.61 \\pm 0.1$.  The Optical Gravitational Lensing Experiment (OGLE) carried out a 4.5 square degree survey of the LMC during the second phase of operations (Jan 1997 -- Nov 2000; \\citealt{soszynski2003}). They discovered 7612 RR Lyr-type objects: 5455 fundamental mode pulsators, 1655 first-overtone, 272 second-overtone, and 230 double-mode pulsators, along with several dozen other short-period pulsating variables. Three objects in their sample exhibited a superimposed RR Lyrae type variability on an eclipsing binary light curve: \\targetnts, OGLE051822.60-691817.3, and OGLE050731.10-693010.3. These could be either the results of blending, or genuine RR Lyr stars in eclipsing binaries. In the latter case it would be possible, with additional spectroscopic data, to determine a reliable mass and radius estimate of an RR Lyr star for the first time. Given their (nearly) constant mean luminosities ($\\langle L\\rangle \\sim 45L_\\odot$ for RR Lyr type ab) and easily recognizable light curves, RR Lyr stars are important for the studies of the formation and evolution of population II stars, and for determining the galactic distance scale (globular clusters, LMC, and the local group; see \\citealt{brown2004,sarajedini2006}). Although considerable effort went into calibrating the luminosity function for RR Lyr stars, it is compromised by dependences on metallicity, reddening, and possibly on period and/or pulsational modes \\citep{sollima2006}.  Of the three candidates, \\citet{soszynski2003} concluded that, based on the period of the binary, the latter two are likely to be blends, leaving \\target as the remaining candidate for an EB system hosting an RR Lyr component. We conducted a thorough feasibility study for all three candidates confirming their preliminary results: only the \\target light curve could be plausibly attributed to an RR Lyr component in the EB system. The OGLE $I_C$ data (archive ID \\verb|LMC_SC6_I_433519|) were passed through the period analysis algorithm PDM \\citep{stellingwerf1978} and two distinct minima were detected: one pertaining to the RR Lyr ($P_\\mathrm{RRLyr} = 0.564876$-d) and the other to the putative EB ($P_\\mathrm{EB} = 8.92371$-d). The disentangling of the two contributions was done by \\emph{polyfit} \\citep{prsa2008} to the data phased at the $P_\\mathrm{RRLyr}$ period (cf.~Fig.~\\ref{lcs}, left). The residuals, folded at $P_\\mathrm{EB}$, exhibit an EB-like shape (Fig.~\\ref{lcs}, right). Assuming canonical values of parameters for the RR Lyr component: $\\mathscr M_\\mathrm{RRLyr} = 0.53 \\mathscr M_\\odot$ \\citep{sandage2004} and $\\mathscr R_\\mathrm{RRLyr} = 5.4 \\mathscr R_\\odot$ \\citep{marconi2005}, we were able to derive a solution consistent with the observed OGLE color of the system ($(B-V)_0 = 0.27$) and the expected absolute magnitude of an RR Lyr star ($M_V \\sim 0.5$). To solve the light curve we used {\\tt PHOEBE} \\citep{prsa2005}, an analysis suite based on the Wilson-Devinney algorithm \\citep{wd1971}. Our preliminary results indicated that the observed light curve could be plausibly explained by an EB hosting a horizontal branch (HB) primary and an RR Lyr secondary. Table \\ref{params} lists the preliminary parameters obtained, which were used to compute the model light curve depicted in Fig.~\\ref{lcs}.  \\begin{figure*} \\centering \\includegraphics[height=\\textwidth,angle=-90]{figs/oglephot.eps} \\caption{Left: light curve of \\target phased with the RR Lyr pulsation period ($P_\\mathrm{RRLyr} = 0.564876$-d). The solid line represents a 2nd order polynomial chain fit obtained by \\emph{polyfit} \\citep{prsa2008}. The apparent scatter under the curve is due to the putative EB signature. Right: residual light curve after subtracting the theoretical RR Lyr pulsation fit, phased at the EB period ($P_\\mathrm{EB} = 8.92371$-d). The solid line depicts a model solution derived by {\\tt PHOEBE}.} \\label{lcs} \\end{figure*}  \\begin{table} \\caption{Preliminary parameters of \\target HB--RR Lyr EB obtained by fitting the residual light curve depicted in Fig.~\\ref{lcs} (right). Values denoted with an asterisk (${}^*$) are assumed. The errors in the table are \\emph{formal}, derived from the covariance matrix.} \\label{params} \\centering \\begin{tabular}{lccc} \\hline\\hline Parameter:           &                 & System           &                   \\\\ & Primary (HB)    &                  & Secondary (RRLyr) \\\\ \\hline $a$ [$R_\\odot$]      &                 & 18.5${}^\\dagger$ &                   \\\\ $i$ [${}^\\circ$]     &                 & $69.0 \\pm 1.45$  &                   \\\\ $T_\\mathrm{eff}$ [K] & $8300 \\pm 113$  &                  & $6700 \\pm 66$     \\\\ $\\mathscr M$ [$\\mathscr M_\\odot$]      & 0.53${}^*$       &                 & 0.53${}^*$        \\\\ $\\mathscr R$ [$\\mathscr R_\\odot$]      & $4.12 \\pm 0.15$  &                 & $5.41 \\pm 0.24$   \\\\ $M_\\mathrm{bol}$     & $0.14 \\pm 0.04$ &                  & $0.5^*$           \\\\ $L (I_C)$ []         & $6.36 \\pm 0.09$ &                  & $6.44 \\pm 0.08$   \\\\ \\hline \\multicolumn{4}{l}{\\rule{0pt}{2.6ex} ${}^\\dagger$derived from $\\mathscr M_1 + \\mathscr M_2 = 1.06 \\mathscr M_\\odot$ and period $P=8.924$-d.} \\end{tabular} \\end{table} ",
        "conclusions": "This paper establishes the nature of \\target as a \\emph{single} RR Lyr type star. Moreover, it shows the two-fold danger of third light contamination: 1) the binary model based on OGLE data alone provided a perfectly plausible physical description with virtually no means of resolving it without additional observations, and 2) although the residuals appeared to clearly point to the binary star signature, it seems that, somewhat ironically, the source of this secular variation is a distinct short period RR Lyr type c star. Being aware of these traps is all the more important in the era of fully automatic surveys and missions, where human supervision and a detailed object-by-object study is no longer possible. False positives present a serious challenge not only for ground-based observations, but for the upcoming space missions like Kepler and Gaia as well, emphasizing the importance of follow-up observations."
    },
    "0808/0808.1279_arXiv.txt": {
        "abstract": "Depending on the density reached in the cores of neutron stars, such objects may contain stable phases of novel matter found nowhere else in the Universe. This article gives a brief overview of these phases of matter and discusses astrophysical constraints on the high-density equation of state associated with ultra-dense nuclear matter. ",
        "introduction": "A forefront area of modern research concerns the exploration of the properties of ultra-dense nuclear matter and the determination of the equation of state (EoS)--the relation between pressure, temperature and density--of such matter. Experimentally, relativistic heavy-ion collision experiments enable physicists to cast a brief glance at hot and ultra-dense matter for times as short as about $10^{-22}$~seconds. This is different for neutron stars, which are observed with radio and X-ray telescopes as radio pulsars and X-ray pulsars. The matter in the cores of such objects is compressed permanently to densities that may be more ten times higher than the densities inside atomic nuclei, which make neutron stars natural astrophysical laboratories that allow for a wide range of (astro) physical studies and astrophysical phenomena (Fig.\\ \\ref{eos-blaschke-fig:multifaceted2}) linked to the properties of ultra-dense nuclear matter and its associated EoS.\\cite{glen97:book,Weber:1999qn,blaschke01:trento,% weber05:ppnp,Sedrakian:2006mq,page06:review,weber07:erice,Lattimer07:a,% Schaffner07:a} \\begin{figure}[tb] \\begin{center} \\includegraphics[keepaspectratio,width=0.65\\textwidth,angle=0] {multifaceted3-eos-blaschke.eps} \\caption[]{The multifaceted connection between high-density nuclear matter and neutron (compact) star phenomena.\\cite{weber05:ppnp}} \\label{eos-blaschke-fig:multifaceted2} \\end{center} \\end{figure} Of particular interest are neutron stars whose observed properties deviate significantly from the norm. Examples of such neutron stars are PSR J0751+1807 whose mass is $2.1 \\pm 0.2~M_\\odot$,\\cite{NiSp05} neutron star RX J1856.5-3754 whose radius may be $\\gsimB 13$~km,\\cite{Trumper:2003we} and XTE J1739-285 whose rotation period may be as small as 0.89~ms.\\cite{Kaaret:2006gr} As discussed in this paper, such neutron star data provide an excellent opportunity to gain profound insight into the properties of nuclear matter at most extreme conditions of density.\\cite{Klahn:2006ir,Klahn:2006iw} ",
        "conclusions": ""
    },
    "0808/0808.1423_arXiv.txt": {
        "abstract": "We present in this paper the first results of a spectropolarimetric analysis of a small sample ($\\sim20$) of active stars ranging from spectral type M0 to M8, which are either fully-convective or possess a very small radiative core. This study aims at providing new constraints on dynamo processes in fully-convective stars.  The present paper focuses on 5 stars of spectral type $\\sim$M4, i.e. with masses close to the full convection threshold ($\\simeq 0.35~\\msun$), and with short rotational periods. Tomographic imaging techniques allow us to reconstruct the surface magnetic topologies from the rotationally modulated time-series of circularly polarised profiles. We find that all stars host mainly axisymmetric large-scale poloidal fields. Three stars were observed at two different epochs separated by $\\sim$1~yr; we find the magnetic topologies to be globally stable on this timescale.  We also provide an accurate estimation of the rotational period of all stars, thus allowing us to start studying how rotation impacts the large-scale magnetic field. ",
        "introduction": "\\label{sec:intro} Magnetic fields play a key role in every phase of the life of stars and are linked to most of their manifestations of activity. Since \\cite{Larmor19} first proposed that electromagnetic induction might be the origin of the Sun's magnetic field, dynamo generation of magnetic fields in the Sun and other cool stars has been a subject of constant interest. The paradigm of the $\\alpha\\Omega$ dynamo, i.e. the generation of a large-scale magnetic field through the combined action of differential rotation ($\\Omega$ effect) and cyclonic convection ($\\alpha$ effect), was first proposed by \\cite{Parker55} and then thoroughly debated and improved \\citep[e.g.,][]{Babcock61, Leighton69}. A decade ago, helioseismology provided the first measurements of the internal differential rotation in the Sun and thus revealed a thin zone of strong shear at the interface between the radiative core and the convective envelope. During the past few years, theoreticians pointed out the crucial role for dynamo processes of this interface -- called the tachocline -- being the place where the $\\Omega$ effect can amplify magnetic fields (see \\citealt{Charbonneau05} for a review of solar dynamo models).  Among cool stars, those with masses lower than about 0.35~\\msun\\ are fully-convective \\citep[e.g.,][]{Chabrier97}, and therefore do not possess a tachocline; some observations further suggest that they rotate almost as rigid bodies \\citep{Barnes05}. However, many fully-convective stars are known to show various signs of activity such as radio, Balmer line, and X-ray emissions \\citep[e.g.,][]{Joy49, Lovell63, Delfosse98, Mohanty03, West04}. Magnetic fields have been directly detected thanks to Zeeman effect on spectral lines, either in unpolarised light \\citep[e.g.,][]{Saar85, Johns96, Reiners06}, or in circularly polarised spectra \\citep[][]{Donati06}.  The lack of a tachocline in very-low-mass stars led theoreticians to propose non-solar dynamo mechanism in which cyclonic convection and turbulence play the main roles while differential rotation only has minor effects \\citep[e.g.,][]{Durney93}. During past few years, several semi-analytical approaches and MHD simulations were developed in order to model the generation of magnetic fields in fully-convective stars. Although they all conclude that fully-convective stars should be able to produce a large-scale magnetic field, they disagree on the properties of such a field, and the precise mechanisms involved in the dynamo effect remain unclear. Mean-field modellings by \\cite{Kuker05} and \\cite{Chabrier06} assumed solid-body rotation and found $\\alpha^2$ dynamo generating purely non-axisymmetric large-scale fields. Subsequent direct numerical simulations diagnose either ``antisolar''  differential rotation (i.e. poles faster than the equator) associated with a net axisymmetric poloidal field \\citep[e.g.,][]{Dobler06}; or strongly quenched ``solar'' differential rotation (i.e. the equator faster than the poles) and a strong axisymmetric toroidal field component \\citep[e.g.,][]{Browning08}.  The first detailed observations of fully-convective stars do not completely agree with any of these models. Among low-mass stars, differential rotation appears to vanish with increasing convective depth \\citep{Barnes05}. This result is further confirmed by the first detailed spectropolarimetric observations of the very active fully-convective star V374~Peg by \\cite{Donati06} and \\cite{Morin08} (hereafter M08) who measure very weak differential rotation (about 1/10th of the solar surface shear). These studies also report a strong mostly axisymmetric poloidal surface magnetic field stable on a timescale of 1~yr on V374~Peg, a result which does not completely agree with any of the existing theoretical predictions. V374~Peg being a very fast rotator, observations of fully-convective stars with longer rotation periods are necessary to generalise these results.  In order to provide theoretical models and numerical simulations with better constraints, it is necessary to determine the main magnetic field properties -- topology and time-variability -- of several fully-convective stars, and to find out their dependency on stellar parameters -- mass, rotation rate, and differential rotation. In this paper, we present and analyse the spectropolarimetric observations of a small sample of stars just around the limit to full convection (spectral types ranging from M3 to M4.5), collected with ESPaDOnS and NARVAL between 2006 Jan and 2008 Feb. Firstly, we briefly present our stellar sample, and our observations are described in a second part. We then provide insight on the imaging process and associated physical model. Afterwards, we present our analysis for each star of the sample. Finally, we discuss the global trends found in our sample and their implications in the understanding of dynamo processes in fully-convective stars. ",
        "conclusions": "\\label{sec:disc}  \\begin{table*} \\caption[]{Magnetic quantities derived from our study. For each star, different observation epochs are presented separately. In columns 2--5 we report quantities from Table~\\ref{tab:sample}, respectively the stellar mass, the rotation period (with an accuracy of 2 digits), the effective Rossby number and the X-ray to bolometric luminosity ratio. Columns 6, 7 and 8 mention the Stokes $V$ filling factor, the reconstructed magnetic energy and the average magnetic flux. Columns 9--13 list the percentage of reconstructed magnetic energy respectively lying in poloidal, dipole (poloidal and $\\ell=1$), quadrupole (poloidal and $\\ell=2$), octupole (poloidal and $\\ell=3$) and axisymmetric modes ($m=0$ / $m < \\ell/2$).} \\begin{tabular}{ccccccccccccc} \\hline Name & Mass & \\Prot & $Ro$ & log$R_X$ & $f_V$ & $<B^2>$ & $<B>$ & pol. & dipole & quad. & oct. & axisymm. \\\\ & (\\msun) & (d) & ($10^{-2}$) & & & ($\\rm10^5\\,G^2$) & (kG) & (\\%) & (\\%) & (\\%) & (\\%) & (\\%) \\\\ \\hline EV~Lac (06) & 0.32 & 4.38 & 6.8 & -3.3 & 0.11 & 4.48 & 0.57 & 87 & 60 & 13 & 3 & 33/36\\\\ \\phantom{EV~Lac} (07) & -- & -- & -- & -- & 0.10 & 3.24 & 0.49 & 98 & 75 & 10 & 3 & 28/31\\\\ YZ~CMi (07) & 0.31 & 2.77 & 4.2 & -3.1 & 0.11 & 5.66 & 0.56 & 92 & 69 & 10 & 5 & 56/61\\\\ \\phantom{YZ~CMi} (08) & -- & -- & -- & -- & 0.11 & 4.75 & 0.55 & 97 & 72 & 11 & 8 & 85/86\\\\ AD~Leo (07) & 0.42 & 2.24 & 4.7 & -3.2 & 0.14 & 0.61 & 0.19 & 99 & 56 & 12 & 5 & 95/97\\\\ \\phantom{AD~Leo} (08) & -- & -- & -- & -- & 0.14 & 0.61 & 0.18 & 95 & 63 & 9 & 3 & 85/88\\\\ EQ~Peg~A (06) & 0.39 & 1.06 & 2.0 & -3.0 & 0.11 & 2.73 & 0.48 & 85 & 70 & 6 & 6 & 69/70\\\\ EQ~Peg~B (06) & 0.25 & 0.40 & 0.5 & -3.3 & na & 2.38 & 0.45 & 97 & 79 & 8 & 5 & 92/94\\\\ V374~Peg (05) & 0.28 & 0.45 & 0.6 & -3.2 & na & 6.55 & 0.78 & 96 & 72 & 12 & 7 & 75/76\\\\ \\phantom{V374~Peg} (06) & -- & -- & -- & -- & na & 4.60 & 0.64 & 96 & 70 & 17 & 4 & 76/77\\\\  \\hline \\label{tab:syn} \\end{tabular} \\end{table*}  \\begin{figure*} \\center{% \\includegraphics[scale=0.60]{fig/plotMP.eps} \\caption[]{Properties of the magnetic topologies of M dwarfs as a function of rotation period and stellar mass. Larger symbols indicate larger magnetic fields while symbol shapes depict the different degrees of axisymmetry of the reconstructed magnetic field (from decagons for purely axisymmetric fields to sharp stars for purely non axisymmetric fields). Colours illustrate the field configuration (dark blue for purely toroidal fields, dark red for purely poloidal fields and intermediate colours for intermediate configurations). Solid lines represent contours of constant Rossby number $Ro=0.1$ and $0.01$ respectively corresponding approximately to the saturation and super-saturation thresholds \\citep[e.g.,][]{Pizzolato03}. The theoretical full-convection limit ($\\mstar \\simeq0.35\\msun$, \\citealt{Chabrier97}) is plotted as a horizontal dashed line.} \\label{fig:plotMP}} \\end{figure*}  Spectropolarimetric observations of a small sample of active M dwarfs around spectral type M4 were carried out with ESPaDOnS at CFHT and NARVAL at TBL between 2006 Jan and 2008 Feb. Strong Zeeman signatures are detected in Stokes $V$ spectra for all the stars of the sample. Using ZDI, with a Unno-Rachkovsky's model modified by two filling factors, we can fit our Stokes $V$ time series. It can be seen on Fig.~\\ref{fig:zdi_spec_adleo}, \\ref{fig:zdi_spec_evlac}, \\ref{fig:zdi_spec_yzcmi} and \\ref{fig:zdi_spec_eqpegab} that rotational modulation is indeed mostly modelled by the imaging code.  From the resulting magnetic maps, we find that the observed stars exhibit common magnetic field properties. (a) We recover mainly poloidal fields, in most stars the observations can be fitted without assuming a toroidal component. (b) Most of the energy is concentrated in the dipole modes, i.e. the lowest order modes. (c) The purely axisymmetric component of the field ($m=0$ modes) is widely dominant except in EV~Lac. These results confirm the findings of M08, i.e. that magnetic topologies of fully-convective stars considerably differ from those of warmer G and K stars which usually host a strong toroidal component in the form of azimuthal field rings roughly coaxial with the rotation axis \\citep[e.g.,][]{Donati03a}.  Table~\\ref{tab:syn} gathers the main properties of the reconstructed magnetic fields and Figure~\\ref{fig:plotMP} presents them in a more visual way. We can thus suspect some trends: (a) The only partly-convective stars of the sample, AD~Leo, hosts a magnetic field with similar properties to the observed fully-convective stars. The only difference is that compared to fully-convective stars of similar $Ro$, we recover a significantly lower magnetic flux on AD~Leo, indicating that the generation of a large-scale magnetic field is more efficient in fully-convective stars. This will be confirmed in a future paper by analysing the early M stars of our sample. (b) We do not observe a growth of the reconstructed large-scale magnetic flux with decreasing Rossby number, thus suggesting that dynamo is already saturated for fully-convective stars having rotation periods lower than 5~\\d, in agreement with \\cite{Pizzolato03} and \\cite{Kiraga07}. Further confirmation from stars with $\\Prot \\gtrsim10~\\d$ is needed. This is supported by the high X-ray fluxes we report, all lying in the saturated part of the rotation-activity relation with log$R_X\\simeq-3$ \\citep[e.g.,][]{James00}. AD~Leo also exhibits a saturated X-ray luminosity despite a significantly weaker reconstructed magnetic field, indicating that the coronal heating is not directly driven by the large-scale magnetic field. (c) The only star showing strong departure from axisymmetry is EV~Lac, i.e. the slowest rotator (though lying in the saturated regime with $Ro=0.07$). Further investigation is needed to check if this a general result for fully-convective stars having $\\Prot \\gtrsim4~\\d$.  The large-scale magnetic fluxes we report here range from 0.2 to 0.8~\\kG. For AD~Leo, EV~Lac and YZ~CMi, previous measurements from Zeeman broadening of atomic or molecular unpolarised line profiles report significantly higher overall magnetic fluxes (several \\kG) \\citep[e.g.,][]{Saar85, Johns96, Reiners07}. We therefore conclude that a significant part of the magnetic energy lies in small-scale fields. Even for the fast rotators EQ~Peg~A and B and V374~Peg for which ZDI is sensitive to scales corresponding to spherical harmonics up to order $\\ell=\\,$12, 20 and 25 (cf. M08), respectively, we reconstruct a large majority of the magnetic energy in modes of order $\\ell\\leq3$. This suggests that the magnetic features we miss with ZDI lie at scales corresponding to $\\ell > 25$ in the reconstructed magnetic fields of mid-M dwarfs.  Three stars of the sample have been observed at two different epochs separated by about 1~yr. AD~Leo, EV~Lac, and YZ~CMi exhibit only faint variations of their magnetic topology during this time gap, the overall magnetic configuration remained stable similarly to the behaviour of V374~Peg (cf. M08). This is at odds with what is observed in more massive active stars, whose magnetic fields reportedly evolve significantly on time-scales of only a few months \\citep[e.g.,][]{Donati03a}.  For three stars of our sample we are able to measure differential rotation and find that our data are compatible with solid-body rotation. In addition, for EV~Lac and YZ~CMi we infer that differential rotation is at most of the order of a few \\mrpd\\ i.e. significantly weaker than in the Sun and apparently lower than in V374~Peg (cf.~M08). This is further confirmed by the fact that the rotation periods we find are in good agreement with photometric periods previously published in the literature (whenever reliable).  This result is consistent with the conclusions of the latest numerical dynamo simulations in fully convective dwarfs with $Ro\\simeq0.01$ \\citep{Browning08} showing that (i) strong magnetic fields are efficiently produced throughout the whole star (with the magnetic energy being roughly equal to the convective kinetic energy as expected from strongly helical flows, i.e., with small $Ro$) and that (ii) these magnetic fields successfully manage to quench differential rotation to less than a tenth of the solar shear (as a result of Maxwell stresses opposing the equatorward transport of angular momentum due to Reynolds stresses). However, these simulations predict that dynamo topologies of fully convective dwarfs should be mostly toroidal, in contradiction with our observations showing strongly poloidal fields in all stars of the sample;  the origin of this discrepancy is not clear yet.   Our study of Stokes $I$ and $V$ time-series allows to measure both the rotational period (\\Prot) and the projected equatorial velocity (\\vsini) of the sample, from which we can straightforwardly deduce the \\rsini. \\Prot\\ is well constrained by our data sets (see the error-bars in Tab.~\\ref{tab:sample}), therefore the incertitude on \\rsini\\ essentially comes from the determination of \\vsini\\ ($\\sigma\\simeq1~\\kms$). This leads to an important incertitude on the \\rsini\\ deduced for slowly rotating stars. As explained in M08, for V374~Peg we find a \\rsini\\ significantly greater than the predicted radius . Here (except for AD~Leo which is seen nearly pole-on) we find $\\rsini \\simeq \\rstar$ (cf. Tab.~\\ref{tab:sample}), suggesting radii larger than the predicted ones. This is consistent with the findings of \\cite{Ribas06} on eclipsing binaries, further confirmed on a sample of single late-K and M dwarfs by \\cite{Morales08}, that active low-mass stars exhibit significantly larger radii and cooler \\teff\\ than inactive stars of similar masses. \\cite{Chabrier07} proposed in a phenomenological approach that a strong magnetic field may inhibit convection and produce the observed trends. This back-reaction of the magnetic field on the star's internal structure may be associated with the dynamo saturation observed in our sample (see above), and with the frozen differential rotation predicted by \\cite{Browning08} when the magnetic energy reaches equipartition (with respect to the kinetic energy).  We also detect significant RV variations in our sample (with peak-to-peak amplitude of up to 700~\\ms). We observe the largest RV variations on the star having the strongest large-scale magnetic field (YZ~CMi). This suggests that although the relation between magnetic field measurements and RV is not yet clear, these smooth fluctuations in RV are due to the magnetic field and the associated activity phenomena. Therefore, if we can predict the RV jitter due to a given magnetic configuration, spectropolarimetry may help in refining RV measurements of active stars, thus allowing to detect planets orbiting around M dwarfs.  The study presented through this paper aims at exploring the magnetic field topologies of a small sample of very active mid-M dwarfs, i.e. stars with masses close the full-convection threshold. Forthcoming papers will extend this work to both earlier (partly-convective) and later M dwarfs, in order to provide an insight on the evolution of magnetic topologies with stellar properties (mainly mass and rotation period). We thus expect to provide new constraints and better understanding of dynamo processes in both fully and partly convective stars."
    },
    "0808/0808.1109_arXiv.txt": {
        "abstract": "The efficiency of star formation governs many observable properties of the cosmological galaxy population, yet many current models of galaxy formation largely ignore the important physics of star formation and the interstellar medium (ISM). Using hydrodynamical simulations of disk galaxies that include a treatment of the molecular ISM and star formation in molecular clouds (Robertson \\& Kravtsov 2008), we study the influence of star formation efficiency and molecular hydrogen abundance on the properties of high-redshift galaxy populations.  In this work, we focus on a model of low-mass, star forming galaxies at $1\\lesssim z\\lesssim2$ that may host long duration gamma-ray bursts (GRBs). Observations of GRB hosts have revealed a population of faint systems with star formation properties that often differ from Lyman-break galaxies (LBGs) and more luminous high-redshift field galaxies.  Observed GRB sightlines are deficient in molecular hydrogen, but it is unclear to what degree this deficiency owes to intrinsic properties of the galaxy or the impact the GRB has on its environment. We find that hydrodynamical simulations of low-stellar mass systems at high-redshifts can reproduce the observed star formation rates and efficiencies of GRB host galaxies at redshifts $1\\lesssim z \\lesssim2$. We show that the compact structure of low-mass high-redshift GRB hosts may lead to a molecular ISM fraction of a few tenths, well above that observed in individual GRB sightlines. However, the star formation rates of observed GRB host galaxies imply molecular gas masses of $10^{8}-10^{9}~M_{\\odot}$ similar to those produced in the simulations, and may therefore imply fairly large average H$_{2}$ fractions in their ISM. ",
        "introduction": "To improve the physical description of star formation in hydrodynamical simulations of galaxies, \\cite[Robertson \\& Kravtsov (2008)]{robertson2008a} implemented a new model for the ISM that includes low-temperature ($T<10^{4}$K) cooling, directly ties the star formation rate to the molecular gas density, and accounts for the destruction of molecular hydrogen by an interstellar radiation field (ISRF) from young stars. They used simulations to study the relation between star formation and the ISM in galaxies and demonstrated that, for the first time, their new model simultaneously reproduces the molecular gas and total gas Kennicutt-Schmidt (KS) relations, the connection between star formation and disk rotation, and the relation between interstellar pressure and the fraction of gas in molecular form \\cite[(e.g. Wong \\& Blitz 2002, Blitz \\& Rosolowsky 2006)]{wong2002a,blitz2006a}. The capability of this model to reproduce both the star formation efficiency and molecular abundance of nearby systems makes it useful for simulating low-mass galaxies that have suppressed H$_{2}$ abundances (and whose star formation rates would be overestimated in common treatments of star formation based on the KS relation) and high-redshift galaxies whose structural properties may vary substantially from local systems (and may therefore not have the same KS relation normalization).  The model should be especially useful for studying low-mass galaxies at high-redshift, such as long duration gamma-ray burst (GRB) host galaxies at $1\\lesssim z \\lesssim2$, which is the focus of this work.  The highly-energetic phenomena known as GRBs were discovered over forty years ago \\cite[(Klebesadel et al. 1973)]{klebesadel1973a}, but their extragalactic origin was confirmed only in the last decade \\cite[(e.g., Metzger et al. 1997)]{metzger1997a}.  Since then, the properties of the cosmological population of galaxies that host GRBs have been increasingly well-studied \\cite[(e.g., Bloom et al. 2002, Le Floc'h et al. 2006, Prochaska et al. 2006, Berger et al. 2007a,b)]{bloom2002a,le_floch2006a,prochaska2006a,berger2007a,berger2007b}. Recently, interest in long duration GRB galaxy hosts as possible tracers of the global star formation history of the universe has motivated systematic studies of their star formation efficiencies and stellar masses \\cite[(Castro Cer\\'on et al. 2008, Savaglio et al. 2008)]{castro_ceron2008a,savaglio2008a}. These studies have found that high-redshift GRB hosts have small stellar masses ($\\log M_{\\star}\\sim 9.3$) and moderate star formation rates ($\\mathrm{SFR}\\sim2.5~M_{\\odot}~\\mathrm{yr}^{-1}$).  Compared with other high-redshift galaxy populations, GRB hosts tend to have lower star formation rates at fixed stellar mass compared with Lyman-break galaxies and lower stellar masses at fixed star formation rate compared with field galaxies \\cite[(for details, see Savaglio et al. 2008)]{savaglio2008a}.  Spectroscopic studies of GRB sightlines have provided additional information about the post-explosion character of the host galaxy ISM.  \\cite[Tumlinson et al. (2007)]{tumlinson2007a} failed to detect H$_{2}$ in five GRB sightlines and suggested that low metallicity and large far ultraviolet ISRF strengths ($10-100\\times$ the Milky Way value) were responsible for destroying molecular hydrogen in GRB hosts.  They interpreted the lack of vibrationally excited H$_{2}$ lines as evidence against the GRB destroying its parent molecular cloud, but noted various caveats to this conclusion such as the parent cloud size or cloud photodissociation before to the GRB. \\cite[Whalen et al. (2008)]{whalen2008a} used one-dimensional radiative hydrodynamical calculations to show that GRBs can ionize nearby neutral hydrogen, but suggested that an additional ISRF is necessary to remove molecular hydrogen from the nearby ISM. \\cite[Prochaska et al. (2008)]{prochaska2008a} studied NV absorption in GRB sightlines, and argued that if nitrogen ionization by GRB afterglows leads to NV absorption then the observations support a scenario where dense, molecular cloud-like environments serve as the sites of GRBs.  Given the increasingly detailed studies of GRB hosts, their interesting ISM and star formation properties, and their low stellar masses, a theoretical study of GRB host galaxy analogues using hydrodynamical simulations that include a treatment of the molecular ISM is warranted.  Below, we present simulations of a model GRB host galaxy that include a prescription for the molecular ISM and star formation in molecular clouds \\cite[(Robertson \\& Kravtsov 2008)]{robertson2008a}.  We use the simulations to examine the star formation efficiency and molecular hydrogen content of galaxies with structural properties similar to those expected for low-mass galaxies at $1\\lesssim z \\lesssim 2$. Below, we discuss our methodology and present some initial results. ",
        "conclusions": ""
    },
    "0808/0808.1615_arXiv.txt": {
        "abstract": "Viable alternatives to astrophysical black holes include hyper-compact objects without horizon, such as gravastars, boson stars, wormholes and superspinars. The authors have recently shown that typical rapidly-spinning gravastars and boson stars develop a strong instability. That analysis is extended in this paper to a wide class of horizonless objects with approximate Kerr-like geometry. A detailed investigation of wormholes and superspinars is presented, using plausible models and mirror boundary conditions at the surface. Like gravastars and boson stars, these objects are unstable with very short instability timescales. This result strengthens previous conclusions that observed hyper-compact astrophysical objects with large rotation are likely to be black holes. ",
        "introduction": "Astrophysical Black Holes (BHs) are believed to be common objects in galaxies. Their mass is expected to span many orders of magnitude, from a fraction of the solar mass (primordial BHs in the galactic halo) to few solar masses (stellar BHs in the galactic plane) up to several billions of solar masses (supermassive BHs in galactic centers). Their angular momentum should be close to the extremal limit due to accretion and mergers \\cite{Gammie:2003qi, Merritt:2004gc}. For example, if quasars are powered by supermassive BHs, astrophysical observations suggest that they should be rotating near the Kerr bound \\cite{Wang:2006bz}.  Unquestionable observational evidence of the existence of BHs is still lacking \\cite{Narayan:2005ie, Abramowicz:2002vt, Lasota:2006jh}. Current astrophysical data cannot rule out ``BH forgeries'', i.e.\\ hyper-compact objects with redshift and geodesics similar to those of BHs, but lacking an event horizon. Several models of hyper-compact objects with these characteristics have been known in the literature for some time. Among these models, gravastars \\cite{Chapline:2000en,Mazur:2001fv} and boson stars \\cite{bosonstars, Berti:2006qt} have been proposed as the most viable alternatives to astrophysical BHs. The authors recently showed that rapidly spinning gravastars and boson stars may develop a strong ergoregion instability \\cite{Cardoso:2007az}. Their typical instability timescales are of order of $0.1$ seconds to 1 week for objects with mass $M = 1 - 10^6 M_{\\odot}$. Therefore, observed astrophysical hyper-compact objects are likely not to be gravastars nor boson stars.  The purpose of this paper is to compute the ergoregion instability for other horizonless, Kerr-like hyper-compact objects: wormholes and superspinars. Wormholes can be objects even simpler than BHs \\cite{Morris:1988cz, visserbook, Lemos:2003jb}. They are infinitesimal variations of the Schwarzschild space-time which may be indistinguishable from BHs \\cite{Damour:2007ap}. In a string theory context, the fuzzball model replaces BHs by horizonless structures \\cite{Mathur:2005zp}. The BH-like geometry emerges in a coarse-grained description which ``averages'' over horizonless geometries and produces an effective horizon at a radius where the individual microstate geometries start to differ. Superspinars are solutions of the gravitational field equations that violate the Kerr bound.  These geometries could be created by high energy corrections to Einstein gravity such as those present in string-inspired models \\cite{Gimon:2007ur, Matsas:2007bj}. Superspinars are expected to have compactness of the order of extremal rotating Kerr BHs and to exist in any mass range.  A rigorous analysis of the ergoregion instability for these models is a non-trivial task; known wormhole solutions are special non-vacuum solutions of the gravitational field equations. Thus their investigation requires a case-by-case analysis of the stress-energy tensor. Exact solutions of four-dimensional superspinars are not known. To overcome these difficulties, the following analysis will focus on a simple model which captures the essential features of most Kerr-like horizonless hyper-compact objects. Superspinars and rotating wormholes will be modeled by the exterior Kerr metric down to their surface, where Dirichlet boundary conditions are imposed. This problem is very similar to Press and Teukolsky's ``BH bomb'' \\cite{bhbombPress, Cardoso:2004nk}, i.e. a rotating BH surrounded by a perfectly reflecting mirror with its horizon replaced by a reflecting surface. These boundary conditions are perfect mirror conditions and require a reflection coefficient $R=1$. In a more realistic model $R<1$ and a certain transmittance $T=1-R$ should be taken into account, which will in principle decrease the strength of the ergoregion instability. We argue that the qualitative behavior of the instability is the same as long as the reflection or the superradiant amplification are large enough. Letting $\\rho$ be a superradiant factor, one expects an ergoregion instability to develop whenever $\\rho (1-T)>1$. In the perfect mirror limit $T=0$ and the superradiant condition is simply $\\rho>1$. The general case can be handled using both the analytical and numerical techniques presented here.  In Section \\ref{sec:superspinars} we introduce the class of objects we will deal with in this work. They are general approximations to superspinars and wormholes with the basic key features retained. In Section \\ref{sec:wh} we show how to solve for the instability analytically in two different regimes. The details of these computations are left for Appendices \\ref{app:slow rotation} and \\ref{app:extremal}. These approximations are compared with numerical results in Section \\ref{sec:num}, where we also show that another kind of instability sets in for general naked singularities. This ``algebraic'' instability can be computed algebraically in the Kerr geometry. We close with a brief discussion of our results. ",
        "conclusions": "} This paper presented a general method for investigating the ergoregion instability of ultra-compact, horizonless  Kerr-like objects. The essential features of these objects have been captured by a simple model whose physical properties are largely independent from the dynamical details of the gravitational system. The method has been applied to superspinars and rotating wormholes.  Numerical and analytic results show that the ergoregion instability of these objects is extremely strong for any value of their angular momentum, with timescales of order $10^{-5}$ seconds for a $1 M_\\odot$ star and $10$ seconds for a $M=10^6 M_{\\odot}$ star. The above investigation confirms previous results for gravastars and boson stars \\cite{Cardoso:2007az}, namely that exotic objects without event horizon are likely to be ruled out as viable candidates for astrophysical hyper-compact objects."
    },
    "0808/0808.0519_arXiv.txt": {
        "abstract": "I describe the IR and X-ray observational campaign we have undertaken for the purpose of determining the nature of the faint discrete X-ray source population discovered by Chandra in the Galactic Center (GC). Data obtained for this project includes a deep Chandra survey of the Galactic Bulge; deep, high resolution IR imaging from VLT/ISAAC, CTIO/ISPI, and the UKIDSS Galactic Plane Survey (GPS); and IR spectroscopy from VLT/ISAAC and IRTF/SpeX.  By cross-correlating the GC X-ray imaging from Chandra with our IR surveys, we identify candidate counterparts to the X-ray sources via astrometry.  Using a detailed IR extinction map, we are deriving magnitudes and colors for all the candidates.  Having thus established a target list, we will use the multi-object IR spectrograph FLAMINGOS-2 on Gemini-South to carry out a spectroscopic survey of the candidate counterparts, to search for emission line signatures which are a hallmark of accreting binaries.  By determining the nature of these X-ray sources, this FLAMINGOS-2 Galactic Center Survey will have a dramatic impact on our knowledge of the Galactic accreting binary population. ",
        "introduction": "The unprecedented sensitivity and angular resolution of \\Chandra has been utilized by Wang \\etal\\ (2002; hereafter W02 \\cite{wang}) and Muno \\etal\\ (2003; hereafter M03 \\cite{muno03}) to investigate the X-ray source population of the Galactic Center (GC).  The W02 ACIS-I survey of the central 0.8\\degs$\\times$2\\degs of the GC revealed a population of $\\sim$800 previously undiscovered discrete weak sources with X-ray luminosities of $10^{32}-10^{35}$\\ergs.  M03 imaged the central 40 pc$^{2}$ (at 8 kpc) around Sgr A*, finding an additional $\\sim$2300 discrete point sources down to a limiting flux of $10^{31}$ erg/s.  More recently, a deeper \\Chandra survey of the central $\\sim$1\\degs around Sgr A* has been obtained; the combination of all of these surveys has revealed a total of $>$10,000 discrete sources (Figure 1; Muno \\etal~ 2008, hereafter M08 \\cite{muno08}).  The harder ($\\geq$3 keV) X-ray sources are likely to be at the distance of the GC, while the softer sources are likely to be foreground X-ray active stars or CVs within a few kpc of the Sun.  Some individual sources have been identified as X-ray transients, high-mass stars, LMXBs, and CVs.  However, the nature of the majority of these newly detected sources is as yet unknown.  \\begin{figure} \\includegraphics[width=0.99\\textwidth]{rmb_fig1.jpg} \\caption{ACIS-I mosaic of the GC, combining the surveys of W02, M03, and M08, as well as additional archival \\Chandra data in the region \\cite{muno08}.  Red is 1-3 keV, green is 3-5 keV, and blue is 5-8 keV.} \\end{figure} ",
        "conclusions": ""
    },
    "0808/0808.2807_arXiv.txt": {
        "abstract": "Charm production gives rise to a flux of very high energy neutrinos from astrophysical sources with jets driven by central engines, such as gamma ray bursts or supernovae with jets. The neutrino flux from semi-leptonic decays of charmed mesons is subject to much less hadronic and radiative cooling than the conventional flux from pion and kaon decays and therefore has a dominant contribution at higher energies, of relevance to future ultrahigh energy neutrino experiments. ",
        "introduction": "Large underground or underwater experiments like IceCube~\\cite{Ahrens:2003ix} and KM3NeT~\\cite{Katz:2006wv} are designed with the goal of observing high energy neutrinos produced in astrophysical sources. The highest energy neutrinos, with energies of $10^{9}$~GeV and higher may be observed in radio detection experiments \\cite{radio}, and with an even higher energy threshold of $10^{12}$~GeV with acoustic detection experiments ~\\cite{Vandenbroucke:2004gv}. We consider astrophysical sources driven by a relativistic jet outflow, accelerated by a central engine such as a black hole~\\cite{Woosley:1993wj,grbreview}. Shock accelerated protons in the jet outflow may give rise to a high-energy neutrino flux~\\cite{Waxman:1997ti}. These neutrinos are potentially produced in hadronic interactions: proton--proton interactions produce charged pions and kaons which subsequently decay into muons and neutrinos. Above the threshold for $\\Delta^+$ production, proton interactions with ambient photons also produce charged pions, and at higher energies, kaons. The relative importance of the $pp$ and $p\\gamma$ contributions to the neutrino fluxes depends on the characteristics of the astrophysical environment. Several types of astrophysical sources have been studied in e.g.~\\cite{Waxman:1997ti,RMW,Ando:2005xi,Koers:2007je,Meszaros:2001ms,Razzaque:2003uv,Horiuchi:2007xi,Wang:2008zm,Murase:2008sp,kt1,kt2}.  High energy pions and kaons are relatively long-lived and therefore subject to both hadronic and radiative cooling before they decay, which downgrades the neutrino energies. Charm production and decay in astrophysical jets is also a source of neutrinos \\cite{kt1}. In this paper, we show that production of charmed mesons in $pp$ collisions gives a large contribution to the neutrino flux at the highest energies, since high energy charmed hadrons ($D^\\pm, D^0$) have short lifetimes and therefore predominantly decay before they interact. Moreover, since the amount of radiative cooling scales as $m^{-4}$, the larger masses of the charmed hadrons lead to less cooling. The neutrino flux from charm is therefore less suppressed up to higher energies. Even though the production cross section is orders of magnitude smaller than for pions and kaons, neutrinos from charm decays become the dominant contribution at high energies.   The energy at which charm begins to dominate depends on the detailed properties of the astrophysical source, and the extent to which charm contributes depends on the maximum energy of the accelerated protons in the jet.  An example of an astrophysical source model is the slow-jet supernova (SJS) model, proposed  by Razzaque, M\\'esz\\'aros and Waxman  (RMW) \\cite{RMW}. It is characterized by a mildly relativistic jet propagating in a collapsing star. This jet does not emerge from the source, and is sometimes referred to as a `choked jet.' This environment has a large optical depth \\cite{Ando:2005xi,Koers:2007je}, so neutrinos may be the only high energy signals to emerge. On the other hand, the jet in a Gamma Ray Burst (GRB) (see e.g.~\\cite{grbreview} for a review) is highly relativistic and the optical depth might be such that photons escape during the burst and do not thermalize. Other GRBs may be choked, so that only neutrinos escape~\\cite{Meszaros:2001ms}.  We present here a new analysis of the neutrino flux for two types of source environments: the SJS and a GRB model, both of which have large magnetic fields in the jets. Our treatment of the proton--proton collisions accounts for charmed meson production, as well as pion and kaon production. The shorter lifetimes of charmed mesons  allow this channel to potentially dominate the neutrino flux from $pp$ collisions. ",
        "conclusions": ""
    },
    "0808/0808.2346_arXiv.txt": {
        "abstract": "The dynamical effects of magnetic fields in models of radiative, Herbig-Haro (HH) jets have been studied in a number of papers. For example, magnetized, radiative jets from variable sources have been studied with axisymmetric and 3D numerical simulations. In this paper, we present an analytic model describing the effect of a toroidal magnetic field on the internal working surfaces that result from a variability in the ejection velocity. We find that for parameters appropriate for HH jets the forces associated with the magnetic field dominate over the gas pressure force within the working surfaces. Depending on the ram pressure radial cross section of the jet, the magnetic field can produce a strong axial pinch, or, alternatively, a broadening of the internal working surfaces. We check the validity of the analytic model with axisymmetric numerical simulations of variable, magnetized jets. ",
        "introduction": "It is now relatively certain that some Herbig-Haro (HH) jets have knot structures which are the result of a time-variability in the ejection. For example, the observations of some jets with organized structures of knots of different sizes (e.~g., HH 30, 34 and 111, see \\citealt{esq07}, \\citealt{rag02} and \\citealt{mas02}) can be reproduced surprisingly well with variable ejection jet models. In the present paper, we study the effect of the presence of a magnetic field on the evolution of a variable jet.  It is still an open question to what extent magnetic fields are important in determining the dynamics of HH jets. The associated problem of radiative, MHD jets has been explored in some detail in the existing literature. \\citet{cer97}, and \\citet{cer99} computed 3D simulations of radiative, MHD jets with different magnetic field configurations (at the injection point). \\citet{fra98} carried out axisymmetric simulations of similar flows.  The problem of an MHD, radiative jet ejected with a time-variable velocity was explored with axisymmetric simulations by \\citet{gara00, garb00, sto00, osu00, fra00, dec06} and \\citet{har07}. Variable, MHD jets were also explored with 3D simulations by \\citet{cer01a,cer01b}. The general conclusions that can be obtained from these simulations is that the internal working surfaces produced by the ejection variability are not affected strongly by a poloidal magnetic field. On the other hand, if the magnetic field is toroidal (or, alternatively, has a strong toroidal component), the material within the working surfaces of the jet flow has a stronger concentration towards the jet axis.  \\citet{gara00} showed that in a variable ejection velocity jet the ``continuous jet beam'' sections in between the working surfaces have a low toroidal magnetic field, which grows in strength quite dramatically when the material goes through one of the working surface shocks into one of the knots. In the present paper, we present a simple, analytic model from which we obtain the conditions under which the toroidal magnetic field produces an axial compression of the internal working surfaces. This analytic model is presented in \\S 2. In \\S 3, we present axisymmetric numerical simulations in which we compare the working surfaces with and without a toroidal magnetic field, showing the effect described by the analytic model. Finally, in \\S 4 we present our conclusions. ",
        "conclusions": "It is a known result that internal working surfaces in radiative, MHD jets with a toroidal magnetic field configuration form dense, axial structures, which do not appear in unmagnetized jets. We present a simple, analytic model with which we show that the strong jump conditions (applied to one of the working surface shocks) imply that the magnetic force dominates over the gas pressure force within a radiative working surface and that the gas pressure force is dominant for a non-radiative working surface (provided that one has a shock Mach number of at least $M_w\\sim 10$ and an Alv\\'enic Mach number which does not exceed $M_w^2$).  Interestingly, the radial dependence of the toroidal magnetic field within a radiative working surface depends only on the cross section of the pre-shock ram pressure $p_{ram}(r)=\\rho(r)v^2(r)$ impinging on the shocks. From equation (\\ref{ffm}), we can see that if we have a $p_{ram}(r)$ that decreases towards the edge of the jet faster than $1/r^2$, the magnetic force within the working surface will be directed outwards, and will tend to increase the width of the working surface.  We have run four simulations (with a top hat cross section for $p_{ram}$, that results in an axially directed magnetic pinch within the working surfaces), therefore in complete consistency with our analytic model. We find that in the non-radiative case the presence of a toroidal magnetic field has very little effect on the structure of the internal working surfaces. We also find that for the radiative case, the presence of a toroidal magnetic field produces a strong axial compression of the material within the internal working surfaces (see Figure 4).  The analytic model presented in this paper can then be used to decide what ram pressure and toroidal magnetic field cross section to use in a magnetized, radiative, variable jet simulation in order to produce internal working surfaces that show narrower or broader structures than what is obtained in non-magnetized jet simulations. This might be a valuable tool when trying to model the knots in specific HH jets, and might provide a possible method for constraining the strength and the configuration of magnetic fields within such objects."
    },
    "0808/0808.2988_arXiv.txt": {
        "abstract": "The shape of the primordial matter power spectrum encodes critical information on cosmological parameters. At large scales, in the linear regime, the observable galaxy power spectrum $\\pobs(k)$ is expected to follow the shape of the linear matter power spectrum $\\plin(k)$, but on smaller scales the effects of nonlinearity and galaxy bias make the ratio $\\pobs(k)/\\plin(k)$ scale-dependent. We develop a method that can extend the dynamic range of the primordial matter power spectrum recovery, taking full advantage of precision measurements on quasi-linear scales, by incorporating additional constraints on the galaxy halo occupation distribution (HOD) from the projected galaxy correlation function $\\wpp$. We devise an analytic model to calculate observable galaxy power spectrum $\\pobs(k)$ in real-space and redshift-space, given $\\plin(k)$ and HOD parameters, and we demonstrate its accuracy at the few percent level with tests against a suite of populated $N$-body simulations. Once HOD parameters are determined by fitting $\\wpp$ measurements for a given cosmological model, galaxy bias is completely specified, and our analytic model predicts both the shape and normalization of $\\pobs(k)$. Applying our method to the main galaxy redshift samples from the Sloan Digital Sky Survey (SDSS), we find that the real-space galaxy power spectrum follows the shape of the nonlinear matter power spectrum at the $1-2\\%$ level up to $k=0.2\\hmpci$ and that current observational uncertainties in HOD parameters leave only few percent uncertainties in our scale-dependent bias predictions up to $k=0.5\\hmpci$. These uncertainties can be marginalized over in deriving cosmological parameter constraints, and they can be reduced by higher precision $\\wpp$ measurements. When we apply our method to the SDSS luminous red galaxy (LRG) samples, we find that the linear bias approximation is accurate to 5\\% at $k\\leq0.08\\hmpci$, but the strong scale-dependence of LRG bias prevents the use of linear theory at $k\\geq0.08\\hmpci$. Our HOD model prediction is in good agreement with the recent SDSS LRG power spectrum measurements at all measured scales ($k\\leq0.2\\hmpci$), naturally explaining the observed shape of $\\pobs(k)$ in the quasi-linear regime. The phenomenological ''$Q$-model'' prescription is a poor description of galaxy bias for the LRG samples, and it can lead to biased cosmological parameter estimates when measurements at $k\\geq0.1\\hmpci$ are included in the analysis. We quantify the potential bias and constraints on cosmological parameters that arise from applying linear theory and $Q$-model fitting, and we demonstrate the utility of HOD modeling of high precision measurements of $\\pobs(k)$ on quasi-linear scales, which will be obtainable from the final SDSS data set. ",
        "introduction": "\\label{sec:int} In the linear regime, the power spectrum of matter fluctuations encodes information about the physics of early universe (e.g., the potential of the field that drives inflation) and about the matter and energy contents of the cosmos. The power spectrum of galaxies can be biased relative to the power spectrum of matter \\citep{nick,BBKS}, but fairly general theoretical arguments imply that the shape of galaxy power spectrum should approach the shape of the linear matter power spectrum $\\plin(k)$ at sufficiently large scales, i.e., \\beeq \\pR(k)=b_0^2~\\plin(k)+N_0, \\label{eq:bias0} \\eneq where $b_0$ is a constant galaxy bias factor and $\\pR(k)$ denotes the real-space galaxy power spectrum \\citep{peter,fry,dhw95,mann,bob2,vijay,andreas,alexia}. The additive ``shot noise'' term $N_0$ reflects both galaxy discreteness and small scale clustering \\citep{bob2,patMc,smith2006}; in general, it can differ from a simple Poisson sampling correction. In the linear regime, distortions of redshift-space structure by peculiar velocities also alter the amplitude but not the shape of galaxy power spectrum \\citep{nick2}. There have therefore been great efforts to measure the galaxy power spectrum on large scales from angular catalogs and redshift surveys and to use the results to test cosmological models (e.g., \\citealt{yupee,baum,feldman,cbp,lcrs,pscrs}). The enormous size of the Two Degree Field Galaxy Redshift Survey (2dFGRS; \\citealt{colless}) and the Sloan Digital Sky Survey (SDSS; \\citealt{sloan}) relative to earlier samples allows much higher precision measurements. The current state-of-the-art power spectrum measurements are Cole~et~al.'s (2005) analysis of the power spectrum in the 2dFGRS, \\citet{nikhil}, and Blake~et~al.'s (2007) analyses of luminous red galaxies (LRGs) with photometric redshifts in the SDSS, and \\citet{will2} and Tegmark~et~al.'s (2006) measurements from the SDSS redshift survey of main sample galaxies and LRGs.  This paper investigates the problem of going from the galaxy power spectrum to the linear matter power spectrum, and hence to cosmological conclusions. The latest observational analyses yield impressive statistical precision on scales near the transition from the linear to the nonlinear regime, e.g., typical 1-$\\sigma$ errors of 5$-$10\\% in $P(k)$ at $k\\simeq0.15\\hmpci$. The critical uncertainty in cosmological interpretation is therefore the accuracy of equation~(\\ref{eq:bias0}) on these scales. The effects of nonlinearity and redshift-space distortions on the $matter$ power spectrum can be computed using numerical simulations or tuned analytic models \\citep[and references therein]{smith}, but details of galaxy formation physics can influence the relation between galaxy and matter power spectra in this regime. \\citet{will2} find that linear theory fits imply different cosmological parameters if applied to measurements with $k\\leq0.06\\hmpci$ or with $k\\leq0.15\\hmpci$, indicating that nonlinear effects have become significant in this regime. Furthermore, \\citet{asp} analyze the SDSS and 2dFGRS galaxy samples and find that the measured shapes of galaxy power spectra differ at a level that cannot be explained by the expected cosmic variance. They show that the likely source of the discrepancy is different scale-dependence of galaxy bias, originating from the different color distributions of galaxies in the SDSS and 2dFGRS samples.  \\citet{2dfn}, \\citet{max2} and \\citet{nikhil} approach this problem by fitting a parametrized model of scale-dependent bias, \\beeq P_\\up{gal}(k)=b_0^2\\plin(k){1+Qk^2\\over1+Ak}, \\label{eq:qmodel} \\eneq where we use $P_\\up{gal}(k)$ to represent the galaxy power spectrum, which can be either in real-space or redshift-space. The functional form is devised for convenience to approximate the scale-dependent bias of galaxy samples obtained by populating the Hubble volume simulation \\citep{hubblevol} using a semi-analytic model of galaxy formation \\citep{andrew1}. Here $A=1.4\\hmpc$ or $1.7\\hmpc$ for real-space power spectrum $\\pR(k)$ or angle-averaged redshift-space power spectrum $P_0(k)$ measurements, respectively, and $Q$ is treated as a free parameter that is marginalized over in deriving cosmological parameter constraints. This approach is adequate {\\it if} equation~(\\ref{eq:qmodel}) is a sufficiently accurate description of scale-dependent bias for some value of $Q$, but it could yield biased parameter estimates or incorrect error bars if the actual scale-dependence is different. It also gives up on extracting cosmological information from scales where bias might be mildly scale-dependent. For example, \\citet[hereafter, T06]{max2} find that cosmological parameters remain unaffected by changes in power spectrum measurements at $k\\geq0.1\\hmpci$ once they marginalize over the value of $Q$. This implies that the statistical constraining power on cosmological parameters is lost at $k\\geq0.1\\hmpci$ by the marginalization process.  In this paper, we present an alternative approach to recovering the shape of the linear matter power spectrum, both more aggressive and more robust than ``marginalizing over $Q$.'' Our approach is based on the halo occupation distribution (HOD) framework, which describes the nonlinear relation between galaxies and matter by specifying the probability $P(N|M)$ that a halo of mass $M$ hosts $N$ number of galaxies of a given type, together with specification of the relative spatial and velocity distributions of galaxies within halos.\\footnote{Throughout this paper, the term ``halo'' refers to a dark matter structure of overdensity $\\rho/\\bar \\rho_m\\simeq200$, in approximate dynamical equilibrium.} The HOD formalism has emerged as a powerful method of modeling galaxy bias \\citep{jing2,uros,chung,john2,roman,andreas} because the dynamics of dark matter halos can be accurately calculated using analytic approximations or $N$-body simulations, and the effects of galaxy formation physics can be parametrized in terms of an HOD and inferred by fitting observational data.  Our strategy for extending recovery of the primordial matter power spectrum is to use complementary information from the measurements of the projected correlation function $\\wpp$ as a constraint to obtain HOD parameters given a cosmological model. We then predict the galaxy power spectrum $P_\\up{gal}(k)$ and study the scale-dependent bias \\beeq b^2(k)\\equiv P_\\up{gal}(k)/\\plin(k). \\label{eq:scale} \\eneq For each cosmological model, fitting $\\wpp$ measurements determines HOD parameters and we can then compute a {\\it unique} prediction of $P_\\up{gal}(k)$, both shape and normalization (which is essentially pinned to the amplitude of $\\wpp$). Uncertainties in HOD parameters introduce uncertainty in $P_\\up{gal}(k)$ and $b^2(k)$, but these uncertainties can be accurately computed and marginalized over. Therefore, we can extend the wavenumber range over which $P_\\up{gal}(k)$ measurements can be used for cosmological parameter constraints, taking full advantage of precision measurements on quasi-linear scales. In practice, we are just using the measured $P_\\up{gal}(k)$ and $\\wpp$ to simultaneously constrain HOD parameters and the cosmological parameters and marginalizing over the former. Relative to the $Q$-model approach, our method adopts a more physically motivated computation of $P_\\up{gal}(k)$ and $b^2(k)$, requiring only the validity of the adopted HOD parametrization, and it brings in the additional information present in $\\wpp$ rather than using only the $P(k)$ shape itself to constrain the scale-dependence of bias.  In principle, the power spectrum $P(k)$ and correlation function $\\xi(r)$ contain the same information. However, they are in practice measured via different estimators and on different scales, where their signal-to-noise ratios are highest and systematic errors are relatively well understood. The information in $P(k)$ and $\\xi(r)$ measurements on these non-overlapping scales is therefore not identical, but complementary. Furthermore, the projected correlation function $\\wpp$ is measured to ease the difficulty in interpreting nonlinear redshift-space distortion of correlation function measurements on small scales. Therefore, the addition of $\\wpp$ measurements at $r_p\\leq30\\hmpc$ brings new information that is not present in $P(k)$ measurements at $k\\leq1\\hmpci$.  A different approach to this problem is to develop an analytic model for predicting the scale-dependence of galaxy bias by using higher-order perturbation theory (e.g., \\citealt{patMc,smith2006}). This approach is elegant and transparent in nature, since it is based on linear theory and its extension to higher-order, while our approach is less {\\it ab initio} in the sense of incorporating elements calibrated by numerical $N$-body simulations in our analytic model. However, the critical uncertainty for this approach based on higher-order perturbation theory is its applicability on quasi-linear scales ($\\gtrsim0.1\\hmpci$), where first-order linear theory is known to be inaccurate, but the measurement precision is highest in practice. In contrast, our approach is fully nonlinear, and phenomenological in nature, so it can be applied down to small scales, limited only by the point at which uncertainties in the HOD parameters introduce systematic uncertainty in the $P(k)$ recovery.  Analyses of galaxy redshift surveys typically estimate the angle-averaged power spectrum $P_0(k)$, i.e., the monopole of the redshift-space power spectrum (e.g., \\citealt{2dfn,will2}). Redshift-space distortions do not alter the shape in linear theory, but they do change the shape in the trans-linear regime (e.g., \\citealt{cole1}), and finger-of-god (FoG) effects have impact out to large scales (e.g., \\citealt{roman}). \\citet{nikhil} and \\citet{blake} deproject the angular clustering measurements of the SDSS LRG sample using photo-$z$ catalogs to estimate the real-space power spectrum, independent of redshift-space distortions. \\citet{max1,max2} use a linear combination of the redshift-space monopole, quadrupole, and hexadecapole that recovers the real-space power spectrum in the linear regime. We will denote this ``pseudo real-space'' power spectrum $\\pzr(k)$. The redshift-space power spectrum estimators can be applied directly to galaxy redshift data or applied after compressing FoG effects. We will investigate $P_\\up{gal}(k)$ and $b^2(k)$ for all of these cases.  To this end, we develop an analytic model in \\S~\\ref{sec:method} for calculating real-space and redshift-space galaxy power spectra given $\\plin(k)$ and a galaxy HOD, drawing on the \\citet{jeremy2} model for redshift-space distortion, which improves on previous work (e.g., \\citealt{uros3,martin,kang,coor}). \\citet{jeremy2} tests the model for computing the redshift-space correlation function against a series of populated $N$-body simulations. Here we extend the model and present additional tests of its applicability to modeling redshift-space power spectra in \\S~\\ref{sec:gal}.  In this paper, we use HOD parameters for volume-limited galaxy samples that have well defined classes of galaxies, focusing on SDSS main galaxy samples with absolute-magnitude limits $M_r\\leq-20$ and $M_r\\leq-21$ \\citep{idit3} in \\S~\\ref{sec:lin}, and SDSS LRG samples with absolute-magnitude limits $-23.2\\leq M_g\\leq-21.2$ and $-23.2\\leq M_g\\leq-21.8$ \\citep{daniel,idit4,zheng5} in \\S~\\ref{sec:lrg} for application of our method.\\footnote{For brevity, we quote the absolute magnitude thresholds $M_r-5\\log h$ and $M_g-5\\log h$ for $h\\equiv1$.} More complete modeling of the conditional luminosity function \\citep{yang} might allow use of flux-limited galaxy catalogs, though it requires more free parameters to provide complete descriptions of the galaxy samples. Here we only consider volume-limited galaxy samples, whose results can be combined to improve statistical precision. We summarize our main results in \\S~\\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum} We have developed an analytic model to predict observable galaxy power spectra $\\pobs(k)$ for specified cosmological and galaxy HOD parameters, and we have verified its accuracy using $N$-body simulations. As potentially observable power spectra $\\pobs(k)$, we have considered the real-space $\\pR(k)$, the redshift-space monopole $P_0(k)$, and the pseudo real-space $\\pzr(k)$, with varying levels of Finger-of-God (FoG) compression for the latter two. Once HOD parameters are determined by fitting the number density $\\bar n_g$ and projected correlation function $\\wpp$ of the observed SDSS galaxy samples, given a specified cosmological model, our analytic model can be used to predict $\\pobs(k)$ measurements.  The large-scale normalization of our predictions is also fixed in the process of fitting $\\wpp$, providing a unique prediction for each combination of cosmological and HOD parameters. In practice, one can simultaneously fit cosmological and HOD parameters using $\\pobs(k)$ and $\\wpp$ as constraints, then marginalize over the HOD in deriving cosmological parameters. By implementing a complete physical model of nonlinear galaxy bias and drawing on the additional information in $\\wpp$, our method allows one to take full advantage of precision measurements of $\\pobs(k)$ on quasi-linear scales ($k=0.1-0.4\\hmpci$), where linear theory or the phenomenological $Q$-model may be insufficiently accurate. Our main findings are as follows:  1. Our analytic model for calculating $\\wpp$ follows the method described in \\citet{jeremy}, with the improved treatment of the scale-dependent halo bias and ellipsoidal halo exclusion corrections. Drawing on the \\citet{jeremy2} model for redshift-space distortion, the analytic model is extended to incorporate calculating real-space and redshift-space power spectra. We have tested its predictions for $\\wpp$ and $\\pobs(k)$ against populated $N$-body simulations spanning cosmological parameter range $\\OM=0.1-0.63$ and $\\rms=0.6-0.95$, with HOD parameters matched to represent two SDSS galaxy samples with absolute-flux limits $M_r\\leq-20$ and $M_r\\leq-21$ \\citep{idit3}. The analytic model reproduces the numerical results of $\\wpp$ to 5\\% or better, and the predictions of $\\pobs(k)$ are consistent with the numerical results to 2\\% at $k=0.1-1\\hmpci$ and to 10\\% at $k=0.025-0.1\\hmpci$, though the finite box size of the simulations makes it difficult to assess the statistical significance of differences on large scales.  2. For the $M_r\\leq-20$ galaxy sample, the pseudo real-space power spectrum $\\pzr(k)$ recovers the true $\\pR(k)$ to 2\\% at $k\\leq0.2\\hmpci$, while the deviation between $\\pR(k)$ and the scaled monopole $P_0(k)$ is already 10\\% at $k=0.1\\hmpci$. However, the deviation of $\\pzr(k)$ from $\\pR(k)$ becomes substantial at $k\\geq0.3\\hmpci$. This deviation can be partly remedied by FoG compression, which suppresses nonlinear behavior of the redshift-space multipoles caused by the random motions of satellite galaxies within halos. With FoG compression threshold $\\sigma_\\up{h}=750~\\kms$, $\\pzrF{750}(k)$ can recover $\\pR(k)$ to 5\\% at $k\\leq0.45\\hmpci$, and at higher $k$ for $\\pzrF{400}(k)$. FoG compression also reduces nonlinearity of the monopole power spectrum, but $P_0^{400}(k)$ can only achieve 10\\% accuracy at $k\\leq0.3\\hmpci$. We conclude that the pseudo real-space method of \\citet{max0} is an effective tool for recovering the nonlinear real-space galaxy power spectrum from redshift-space measurements, especially if it is combined with accurate FoG compression.  3. The nonlinear $matter$ power spectrum describes the nonlinear real-space galaxy power spectra to 1\\% at $k\\leq0.2\\hmpci$ for the $M_r\\leq-20$ and $M_r\\leq-21$ galaxy samples, up to an overall bias factor $b^2_0$. The shape of the scale-dependent bias function $b^2(k)/b_0^2$ for $\\pzr(k)$ is qualitatively similar to $\\pR(k)$ at $k\\leq0.3\\hmpci$, but the shape for $P_0(k)$ is completely different over the entire range we consider here. FoG compression makes little difference to $b^2(k)/b^2_0$ for $\\pzr(k)$, but a large difference for $P_0(k)$. For these SDSS main galaxy samples, the $Q$-model prescription traces our calculation of $\\pR(k)$ relatively well at $k\\geq0.1\\hmpci$, but its shape on large scales differs, so it might induce some overall bias in cosmological parameters when fitted to $\\pobs(k)$ measurements that have large uncertainties at $k\\leq0.05\\hmpci$. Similar trends but with larger discrepancy are found in comparison to our $\\pzr(k)$ and $P_0(k)$ calculations.  4. Uncertainties in computing $\\pobs(k)$ in our method arise from observational uncertainties in the HOD parameters and from uncertainty in the adopted parametrization itself. We have examined these uncertainties by adopting a flexible HOD parametrization with freedom to explore a wider range of plausible halo occupation functions. For the $M_r\\leq-20$ sample with the \\citet{idit3} uncertainties in $\\wpp$, the uncertainty in the predicted $\\pobs(k)$ is 2\\% at $k=0.2\\hmpci$, becomes progressively smaller at lower $k$, and climbs up to 4\\% at $k=0.5\\hmpci$. The uncertainty is a factor of two smaller for the $M_r\\leq-21$ sample, roughly the ratio of the fractional $\\wpp$ measurement errors of the two samples. We have not investigated the uncertainties associated with possible environmental variations of the HOD \\citep{croton}. Based on work to date, we expect that such variations might lead to few percent uncertainties in the overall normalization predicted for $\\pobs(k)$ after fitting $\\wpp$, but that the impact on scale-dependence of $b^2(k)/b^2_0$ would be smaller.  5. Moving to the LRG regime, we have tested our analytic model predictions against the $z=0$ output of a large volume, 1024$^3$-particle $N$-body simulation \\citep{warren1}, populated based on Zheng et~al.'s (2008) HOD fits to $\\wpp$ for two volume limited SDSS LRG samples \\citep{idit4}. The analytic model predicts $\\xi_0(r)$ and $\\xiR(r)$ to 5\\% or better over the range $0.1\\hmpc\\leq r\\leq30\\hmpc$, and the predictions for $P_0(k)$ and $\\pR(k)$ have similar accuracy over the range $0.01\\hmpci\\leq k\\leq0.3\\hmpci$.  6. For the LRG samples, the linear (scale-independent) bias approximation remains accurate at the 5\\% level to $k=0.08\\hmpci$ for the $-23.2\\leq M_g\\leq-21.2$ sample and to $k=0.1\\hmpci$ for the $-23.2\\leq M_g\\leq-21.8$ sample. There is little variation among $\\pzr(k)$, $P_0(k)$, and $\\pR(k)$, because LRGs are mainly central galaxies in massive halos, so random motions of satellite galaxies have little impact. Similarly, FoG compression has only a small impact on $b^2(k)/b_0^2$ for these samples. Both samples show strong scale-dependence of bias at $k\\geq0.1\\hmpci$, much more than for main sample galaxies.\\footnote{This result, obtained by fitting HODs and computing $\\pobs(k)$ with our analytic model, confirms the result of T06 inferred by fitting $\\pobs(k)$ with $Q$-models. However, the actual scale-dependence we find for LRGs is somewhat weaker than that inferred by T06.} If we fit $b^2(k)/b_0^2$ from our HOD models with the best $Q$-model over the range $k=0.01-0.2\\hmpci$, the largest deviation is 7\\%.  7. We have presented a preliminary comparison of our analytic model predictions to the T06 measurements of the LRG $\\pobs(k)$, with no FoG compression. The difference between our volume-limited samples and the T06 flux-limited sample precludes a full quantitative assessment, but the qualitative agreement is remarkably good over the full range of the measurements, $k=0.01-0.2\\hmpci$ (Fig.~\\ref{fig:lrg-pow}). Fits with different cosmological parameters differ on large scales and, to a smaller degree, at $k\\geq0.2\\hmpci$, indicating that measurements to smaller scales would provide additional discriminatory power.  8. Looking to the future, we have generated synthetic $\\pobs(k)$ data from our analytic model with error bars half those of T06, then fit them to successively higher $k_\\up{max}$ with linear theory, the $Q$-model, and the HOD model. Cosmological parameters from linear theory fits are badly biased for $k_\\up{max}\\geq0.1\\hmpci$, while for $k_\\up{max}=0.09\\hmpci$ they are biased at less than 1-$\\sigma$. Parameters from the $Q$-model are minimally biased for $k_\\up{max}=0.09\\hmpci$, biased by 1.2-$\\sigma$ for $k_\\up{max}=0.2\\hmpci$, and biased by many-$\\sigma$ for $k_\\up{max}=0.4\\hmpci$. Since the synthetic data are generated from the HOD model, the HOD parameter estimates are unbiased, and the error bars in cosmological parameters shrink steadily as $k_\\up{max}$ is increased from $0.1\\hmpci$ to $0.2\\hmpci$ to $0.4\\hmpci$.  Results~7 and~8 are especially encouraging. Using only the HOD model and the information in $\\wpp$, our method predicts exactly the scale-dependent bias for LRGs that is required to transform the linear power spectrum from WMAP3 into the SDSS galaxy power spectrum measured by T06. This is in contrast to $Q$-model fitting, where a phenomenological parameter (motivated by simulation results but with no clear physical interpretation) is introduced specifically to account for the difference between the linear theory $P(k)$ and the observed power spectrum.  Despite the clear evidence that scale-dependent bias affects the LRG power spectrum beyond $k=0.1\\hmpci$, result~8 shows that one can gain substantial additional leverage on cosmological parameters with HOD modeling of power spectrum measurements up to $k=0.2-0.4\\hmpci$, and possibly beyond. Realizing this opportunity will require several investigations beyond those presented here. First, we will need more large volume simulations to test and, if necessary, refine our analytic model to the level of accuracy demanded by the final SDSS data set. Second, we must explore more thoroughly the uncertainties associated with the HOD fitting, including alternative parametrizations, the impact of velocity bias on redshift-space predictions, and the possible impact of environmental variations of the HOD. Given the growth of current and future galaxy surveys in depth and redshift, these investigations will be needed to go beyond linear theory. Precise measurement of the primordial matter power spectrum will play a crucial role in constraining cosmological parameters and testing dark energy theories."
    },
    "0808/0808.2493_arXiv.txt": {
        "abstract": "Studies of strong gravitational lensing in current and upcoming wide and deep photometric surveys, and of stellar kinematics from (integral-field) spectroscopy at increasing redshifts, promise to provide valuable constraints on galaxy density profiles and shapes. However, both methods are affected by various selection and modelling biases, whch we aim to investigate in a consistent way. In this first paper in a series we develop a flexible but efficient pipeline to simulate lensing by realistic galaxy models.  These galaxy models have separate stellar and dark matter components, each with a range of density profiles and shapes representative of early-type, central galaxies without significant contributions from other nearby galaxies. We use Fourier methods to calculate the lensing properties of galaxies with arbitrary surface density distributions, and Monte Carlo methods to compute lensing statistics such as point-source lensing cross-sections.  Incorporating a variety of magnification bias modes lets us examine different survey limitations in image resolution and flux. We rigorously test the numerical methods for systematic errors and sensitivity to basic assumptions.  We also determine the minimum number of viewing angles that must be sampled in order to recover accurate orientation-averaged lensing quantities. We find that for a range of non-isothermal stellar and dark matter density profiles typical of elliptical galaxies, the combined density profile and corresponding lensing properties are surprisingly close to isothermal around the Einstein radius.  The converse implication is that constraints from strong lensing and/or stellar kinematics, which are indeed consistent with isothermal models near the Einstein radius, cannot trivially be extrapolated to smaller and larger radii. ",
        "introduction": "\\label{S:motivation}  \\subsection{Learning from galaxy density profiles and shapes} \\label{SS:intro1}  Observational constraints on galaxy density profiles and shapes can be used to study a great variety of problems, from basic cosmology to the connections between dark matter and baryons.  For example, the spherically-averaged form of the dark matter halo density profile, and the distribution of halo shapes, both depend on cosmological parameters to the extent that those parameters affect the process of structure formation \\citep[e.g.,][]{2001MNRAS.321..559B, 2006MNRAS.367.1781A}, and on the physics of galaxy formation to the extent that baryons modify the dark matter distribution \\citep[e.g.,][]{1986ApJ...301...27B, 2004ApJ...616...16G, 2004ApJ...611L..73K, 2007ApJ...658..710N, 2008ApJ...672...19R}. While there is still some disagreement in the literature about the theoretical predictions, even among those papers cited above, one key result is that there is considerable scatter in both the density profiles and shapes of dark matter halos, which presumably reflects different formation histories \\citep[e.g.,][]{2002ApJ...568...52W}.  Several observational tools have been used to constrain the density profiles and shapes of galaxies, at different length and mass scales, and redshifts. The main tools we consider are gravitational lensing and stellar kinematics. X-ray data are also useful for understanding the density profiles and shapes of massive galaxies and galaxy clusters, but are beyond the scope of our investigation.  Gravitational lensing is the deflection of light from distant sources by the gravitational fields of intervening lens galaxies.  Strong lensing occurs in the central regions of galaxies (and clusters) where the light bending is extreme enough to produce multiple images of a background source.  It can be used to study the mass distributions of individual galaxies on projected scales of typically several kpc \\citep[for a review, see ][]{Saas-Fee}.  Weak lensing is a complementary phenomenon in which the deflection of light slightly distorts the shapes of background sources without creating multiple images \\citep[for a review, see][]{2001PhR...340..291B}.  Weak lensing probes galaxy mass distributions on scales from tens of kpc out to about 10 Mpc, but only yields constraints on ensemble averages, since currently weak lensing by galaxies can only be detected by stacking many lens galaxies. While galaxy clusters can by studied on an individual basis using weak lensing, our focus is on galaxies.  High-quality (two-dimensional) data on the kinematics of stars as well as gas in the inner parts (a few to ten kpc) of galaxies are now readily available, thanks in particular to strong progress in integral-field spectroscopy in the last decade. Kinematics in the outer parts of late-type galaxies can often be observed from the presence of neutral hydrogen \\citep[e.g.,][]{1981AJ.....86.1825B, 1985ApJ...295..305V, 1996MNRAS.281...27P, 2007MNRAS.376.1513N}. In the outer parts of early-type galaxies, however, cold gas is scarce \\citep[but see e.g.][]{1994ApJ...436..642F, 1997AJ....113..937M, 2008MNRAS.383.1343W}, so we are left with discrete kinematic tracers such as planetary nebulae and globular cluster \\citep[e.g.,][]{2003ApJ...591..850C, 2007ApJ...664..257D} out to tens of kpc.  Current and upcoming wide and deep photometric surveys will reveal a vast number of strong lensing events \\citep[e.g.,][]{ 2004AAS...20510827F, 2004NewAR..48.1085K, 2004ApJ...601..104K, 2005NewAR..49..387M}, and will also allow for extensive, complementary weak lensing analysis \\citep[e.g.][]{2002SPIE.4836...10T, 2004SPIE.5489...11K}. At the same time, there is a rapid increase in the availability of two-dimensional kinematics of galaxies nearby \\citep[e.g.,][]{2004MNRAS.352..721E, 2006MNRAS.373..906M}, and even at high(er) redshift \\citep[e.g.,][]{2007ApJ...668..738V, 2007ApJ...671..303B}. These data sets, individually as well as combined, may provide strong constraints on galaxy density profiles and shapes, but the analyses involved are subject to selection and modelling biases.  \\subsection{Selection and modelling biases} \\label{SS:intro2}  If we want to use strong lensing and kinematics to make robust tests of cosmological predictions, we need to answer two questions: Can strong lensing and kinematic analyses yield accurate constraints on galaxy density profiles?  Are the galaxies in which we can make the measurements (especially in the case of strong lensing) representative of all galaxies?  We refer to these two concerns as \\emph{modelling biases} and \\emph{selection biases}, respectively.  The question of selection bias is particularly important given the diversity in the galaxy population. It is well known that strong lensing favours early-type over late-type galaxies, and massive galaxies over dwarfs \\citep[e.g.,][]{1984ApJ...284....1T, 1991MNRAS.253...99F}.  But even within the population of massive early-type galaxies, to what extent does strong lensing favour galaxies whose dark matter halos have inner slopes that are steeper than average, or concentrations that are higher than average?  Also, to what extent does strong lensing favour galaxies with particular shapes and/or orientations with respect to the line-of-sight? To phrase these questions formally, consider some parameter $x$ describing the galaxy density profile or shape (e.g., the inner slope of the dark matter density profile). We need to understand how the distribution $p_\\mathrm{SL}(x)$ among strong lens galaxies compares to the underlying distribution $p(x)$ for all galaxies. Any analysis that includes strong lensing with some other technique used to study the same systems will suffer from strong lensing-related selection biases, since we never get to choose which systems will be a strong lens.  Selection biases are critical when we want to interpret constraints on lens galaxy density profiles and/or shapes from strong lensing, possibly in combination with kinematics \\citep[e.g.][]{2005ApJ...623..666R, 2006ApJ...649..599K} in comparison with predictions from cosmological models.  On the other hand, modelling biases may occur because of (often unavoidable) assumptions made when analysing gravitational lensing and/or kinematic data because of the finite number of constraints available from the data. Since galaxies are in general non-spherical, we can only correctly interpret the observations if we know the viewing direction, and even then the deprojection might not be unique \\citep[e.g.][]{1987IAUS..127..397R}. Moreover, strong lensing is subject to the so-called mass-sheet degeneracy: part of the deflection and magnification of the light from the background source can be due to mass along the line-of-sight that is not associated with the lens galaxy itself\\footnote{Part of this mass-sheet degeneracy may be overcome in future surveys through the measurements of time delays for an adopted Hubble constant.}. In addition, the constraints from strong lensing on the galaxy density are typically limited in radius to around the Einstein radius, and even then, without secure, unaffected measurements of the flux of the lensing images, the density profile is difficult to recover.  Kinematics can be obtained over a larger radial extent, but in particular stellar kinematics suffer from the so-called mass-anisotropy degeneracy: a change in the measured line-of-sight velocity dispersion can be due to a change in total mass, but may also be the result of velocity anisotropy.   In a series of papers, we are developing a pipeline for using realistic galaxy models to simulate strong lensing and kinematic data, and assess how both selection and modelling biases affect typical strong lensing and kinematic analyses. In this first paper, we present the simulation pipeline for point-source lensing. We focus on strong lensing of quasars by early-type, central galaxies at different mass scales, using two-component mass profiles (dark matter plus stellar component) that are consistent with existing photometry and stacked weak lensing data from SDSS \\citep{2003MNRAS.341...33K,2006MNRAS.368..715M}, $N$-body, and hydrodynamic simulations. In \\citeauthor{2008paperII} (\\citeyear{2008paperII}, hereafter Paper~II), we use this pipeline to study selection biases in strong lensing surveys. In future work we will address modelling biases in strong lensing and kinematics (both when studied separately and when combined). Our overall goal is to determine how strong lensing, kinematics, and possibly other probes (weak lensing, and X-ray data for cluster mass scales) can be used to obtain robust, unbiased constraints on galaxy density profiles and shapes.  To make our presentation coherent, and to clarify our notation and terminology, we first review the analytic treatment of ellipsoidal mass distributions with different profiles and shapes (Section~\\ref{S:profiles-shapes}), and the basic theory of strong lensing (Section~\\ref{S:stronglensing}).  We then describe our simulation pipeline for point-source lensing in depth.  In Section~\\ref{S:galmodels} we present our choices for the masses, profiles, and shapes of the galaxy models. In Section~\\ref{S:lensingcalc} we discuss the numerical methods we use for lensing calculations, including numerous tests. A summary of the simulation pipeline, and some implications for strong lensing analyses are given in Section~\\ref{S:conclusions}. ",
        "conclusions": "\\label{S:conclusions}     We have presented a flexible simulation pipeline for coherent investigations of selection and modelling biases in strong lensing surveys. We have focused on point-source lensing by two-component galaxy models meant to emulate realistic early-type, central galaxies at two different mass scales: a lower, $\\sim 2L_*$ galaxy mass scale and a higher, $\\sim 7 L_*$ group mass scale (with $L_*$ in the $r$-band). Below is a list of our main conclusions regarding the construction of the simulation pipeline for lensing by realistic galaxy models: \\begin{itemize} \\item We include both cusped and deprojected S\\'ersic density profiles, with observationally-motivated choices of masses and scale lengths, and a range in density profile parameters, separately for the stellar and dark matter component. We use seven different models for the galaxy shapes with the stellar component rounder than the dark matter component, but with the axes intrinsically aligned. [Sections~\\ref{SS:massscales}--\\ref{SS:sersicsim}] \\item When we change the density profile parameters away from the adopted fiducial values, we preserve the total and virial mass of the stellar and dark matter components (respectively) by changing the density amplitude. We keep the scale radius fixed to comply with observed size-luminosity relations for early-type galaxies and concentration-mass relations for dark matter halos. Similarly, for the non-spherical shapes we change the scale length to preserve the mass but keep the scale radius (and hence the concentration) the same as for the fiducial spherical case. [Sections~\\ref{SSS:normalization}, \\ref{SSS:scalelengthnonspherical}, \\ref{SSS:sersicnormalization}] \\item We choose fiducial lens and source redshifts of $0.3$ and $2.0$, resulting in Einstein radii $\\Rein$ at about one effective radius $R_e$. [Section~\\ref{SS:redshifts}] \\item We require multiple ways of handling magnification bias to sample different parts of the observed quasar luminosity function and to allow for application to different survey limitations in image resolution and flux. [Section~\\ref{SS:magbias}] \\item When the surface mass density is known analytically (as for our deprojected S\\'ersic models), we compute the corresponding lensing deflection and magnification analytically (for circular symmetry) or with numerical integrals (for elliptical symmetry).  When the surface mass density cannot be computed analytically, as for the cusped models, we construct a surface density map and use Fourier methods to compute the lensing deflection and magnification.  We validate and use Monte Carlo methods to calculate lensing statistics, including unbiased and biased cross-sections and image separations. [Sections~\\ref{SS:kapmap}--\\ref{SS:validation}] \\item For our map-based lensing calculations, tests for convergence at the 5 per cent level suggest that the required map resolution is $\\sim 25$ pixel lengths per $\\Rein$, corresponding to $0.025\\arcsec$ ($\\simeq 0.11$\\,kpc) and $0.075\\arcsec$ ($\\simeq 0.33$\\,kpc) for the lower and higher mass scale, respectively. The resolution needs to be fine enough to simultaneously resolve the steepness of the inner slope of the density and recover the smoothness of the shape of the inner critical curve. Similarly, convergence tests suggest that the required box size is $\\sim 16\\,\\Rein$, corresponding to $10\\arcsec$ ($\\simeq 43.6$\\,kpc) and $30\\arcsec$ ($\\simeq 131$\\,kpc) for the lower and higher mass scale, respectively. [Sections~\\ref{SS:resolution} and~\\ref{SS:boxsize}] \\item Lensing cross-sections for fairly general triaxial models vary in a smooth and monotonic way with viewing direction, which can be efficiently sampled through projected axis ratios (rather than viewing angles) of the surface density. Surprisingly few samplings are necessary: 10, 8, and 24 viewing directions can effectively determine the orientation-averaged cross-sections for oblate, prolate, and triaxial models (respectively) with 5 per cent precision. [Section~\\ref{SS:projangle}] \\end{itemize} We have implemented the construction of the realistic galaxy models and efficient computation of the corresponding surface mass density maps in \\idl\\ (see also Section~\\ref{SS:compsurfmaps}). All subsequent strong lensing calculations are done with an updated and extended version of \\gravlens. We use scripts to coordinate and analyse the extensive data flow to and from these two parts, and to make it into a coherent simulation pipeline.  Our investigations have revealed a number of useful, general points related to the interpretation of observational strong lensing results: \\begin{itemize} \\item In the vicinity of the Einstein radius, the intrinsic and projected logarithmic slopes of the total density are close to the values $\\gamma=2$ and $\\gamma'=1$ of an isothermal profile, for all of our galaxy models (see Figs.~\\ref{F:3dlogslope} and~\\ref{F:2dlogslope}, and the text of Sec.~\\ref{SSS:intrinsicprofiles}). Consequently, a measurement of the total density profile near the Einstein radius cannot (alone) be used to determine the dark matter inner slope for a two-component model. This is a non-trivial conclusion given the variety of intrinsic non-isothermal density profiles used for both the stellar and dark matter component, including profiles with a true central cusp and those that do not asymptote to a particular inner slope at any scale. \\item We also find that the lensing deflection curves are nearly flat around (and even beyond) $\\Rein$, similar to a flat deflection curve for an isothermal model (see Fig.~\\ref{F:alpha}), particularly for the lower-mass galaxy scale. This again implies that constraints from strong lensing cannot be extrapolated to radii much smaller or larger than the Einstein radius. We emphasize that the parameters for the galaxy models were \\emph{not} a priori chosen to mimic this effect, but that it truly seems to be an ``isothermal conspiracy'' \\citep[e.g.,][]{2003ApJ...595...29R}. \\item Because of our assumption of alignment between the intrinsic axes of the stellar and dark matter components, only the two shape models with a triaxial dark matter component and a rounder (triaxial or oblate) stellar component can have a significant misalignment between their projected axes.  Even then, significant misalignment occurs only for a limited number of viewing directions, several of which lead to a relatively round projected stellar component such that the misalignment is not very important in practice (the position angle of a round stellar component is on the sky, after all, difficult to establish). The small misalignments inferred for relatively isolated lens systems, as well as detailed dynamical modelling of early-type galaxies, support near intrinsic alignment between stars and dark matter (see Section~\\ref{SSS:misalignment}). \\end{itemize}  In Paper~II, we use the flexibility of the pipeline to study selection biases related to the galaxy mass, shape, orientation, and various parameters of the dark matter and stellar profiles. In subsequent work, we will investigate modelling biases by analysing the mock lens systems produced by the pipeline using the lens modelling tools within the \\gravlens\\ package. We have therefore created a lensing simulation pipeline that will be crucial for deriving inferences about galaxy density profiles and shapes, $H_0$, and other parameters from the thousands of lenses that will be discovered in the coming years in large photometric surveys such as Pan-STARRS, LSST, SNAP, and SKA.  Furthermore, following a similar approach as in \\citeauthor{2008MNRAS.385..614V} (2008a), we plan to extend the pipeline to produce projected kinematics of our galaxy models that mimic the two-dimensional kinematic observations of early-type galaxies. We will then investigate how well we can expect to recover the intrinsic density profile and shape of early-type galaxies, based solely on dynamical models fitted to these simulated kinematics, or on a combination of kinematics and strong lensing (using techniques similar to \\citeauthor{2008arXiv0807.4175V} 2008b). This work will be valuable for understanding the modelling biases in analyses of the kinematic data that are becoming available for hundreds of galaxies at increasing redshifts, and for understanding both modelling and selection biases in joint lensing+kinematics studies."
    },
    "0808/0808.1258_arXiv.txt": {
        "abstract": "{Modelling the emission properties of compact high energy sources such as X-ray binaries, AGN or $\\gamma$-ray bursts represents a complex problem. Contributions of numerous processes participate non linearly to produce the observed spectra: particle-particle, particle-photon and particle-wave interactions. Numerical simulations have been widely used to address the key properties of the high energy plasmas present in these sources.} {We present a code designed to investigate these questions. It includes most of the relevant processes required to simulate the emission of high energy sources. } {This code solves the time-dependent kinetic equations for homogeneous, isotropic distributions of photons, electrons, and positrons. We do not assume that the distribution has any particular shape. We consider the effects of synchrotron self-absorbed radiation, Compton scattering, pair production/annihilation, e-e and e-p Coulomb collisions, e-p bremsstrahlung radiation and some prescriptions for additional particle heating and acceleration.} {We illustrate the code's computational capabilities by presenting comparisons with earlier works and some examples.  Previous results are reproduced qualitatively but some differences are often found in the details of the particle distribution. As a first application of the code, we investigate acceleration by second order Fermi-like processes and find that the energy threshold for acceleration has a crucial influence on the particle distribution and the emitted spectrum.}{} ",
        "introduction": "High energy sources, such as X-ray binaries, active galactic nuclei (AGN hereafter), or $\\gamma$-ray bursts, exhibit spectra detectable to very high energy. This radiation must originate in a plasma for which a significant fraction of the particles have relativistic energies. Understanding the properties of these hot plasmas remains a challenge in the modelling of X- and $\\gamma$-ray sources.  Among the many processes at work, there are particle-particle interactions, such as Coulomb collisions, particle-photon interactions such as Compton scattering, synchrotron radiation, bremsstrahlung emission, or pair production/annihilation, and particle-wave interactions that lead to particle acceleration. However, the way they add or compete is highly non-linear and the cross sections involved are complex. Investigating a large parameter space is required, and in spite of important breakthroughs, these plasmas are still poorly understood. Analytical studies provided interesting qualitative results with approximations, but a more general approach based on numerical simulations is required to explain the details and complexity of contemporary observations.  The first detailed investigations were analytical attempts to model the Compton scattering in thermal plasmas of fixed temperature \\citep[][]{BZS71,Sunyaev80,Guilbert81,Zdziarski85,Guilbert86}. In parallel, some of these results were confirmed by Monte-Carlo simulations \\citep[e.g.][]{PSS83,Gorecki84}. The additional role of pair production and annihilation in thermal plasmas, whose temperatures were determined self-consistently, was then studied both analytically \\citep{Svensson82b,Svensson83,Guilbert85,Kusunose87} and numerically \\citep{Zdziarski84,Zdziarski85}. These works constituted significant advances because they explicitly accounted for the back reaction of the radiation field on the plasma temperature. However, they were limited to thermal distributions of particles, whereas significant evidence of strongly non-thermal populations was found in many sources. For instance, spectra of blazars or radio loud AGN were shown to be shaped at least by the synchrotron self-compton emission of purely non-thermal electrons \\citep[e.g.][]{Ghisellini98}. At these high energies, accelerated particles cool on very short timescales before they can be thermalized for instance by two body collisions. The balance between this cooling and acceleration typically produces non-thermal distributions. Acceleration processes are still poorly understood. A simple way to simulate the effect of particle acceleration is to inject particles at high energy. Although it does not reproduce exactly the physics involved, this prescription has been widely used and produced interesting results (as shown in most of the references cited here). Significant effort has also been taken in developing more precise modelling of acceleration mechanisms, but in such studies, the radiation field is treated crudely \\citep{LKL96,DML96,Li97,Katar06}.  With the increasing number of considered processes and the increasing precision of their description, numerical analysis has become a prime method of investigation. Even so, a full treatment of the problem accounting for the coupled evolution of inhomogeneous, anisotropic  distributions of leptons and photons, both in momentum and position spaces, appears to be still beyond the capabilities of present-day computers.  Numerical simulations of high energy plasmas have been performed mainly following three different approaches which all make trade-offs  between the various aspects of the problem.  First, the Monte Carlo technique \\citep{PSS77,Stern95} allows one to follow particles and photons in space, time, and energy as they undergo mutual interactions. It solves the full radiative transfer problem and accounts explicitly for geometrical effects. At present, the MC method is probably the mos effective way to model fully 3-dimensional problems. However, this detailed procedure is time consuming, particularly when modelling the rapid dynamics of the non-thermal electron population in momentum space \\citep{mj00}, and when synchrotron self-absorption effects are important \\citep[see discussion in][]{Stern95}. For this reason, the Monte-Carlo methods have been applied to date to pure Maxwellian plasmas and/or steady state problems with 3D geometry.   Another method that accounts correctly for the geometry, involves solving numerically the exact radiation transfer equation for given geometries and particle distributions \\citep{Poutanen96}. This method is far more efficient than Monte Carlo simulations which makes it easier to compare with data. It is, however, far less versatile than Monte Carlo methods and does not solve the kinetic equations for particles. The back reaction of the radiation field on the particle distribution is modelled only for the assumption of a Maxwellian plasma (in which case the plasma temperature may be adjusted according to energy balance). The method applicability is also limited to the resolution of steady state problems.  The third approach, which we adopt in this paper, abandons the detailed description of the geometry to concentrate on the kinetic effects. It consists of solving the local kinetic equations for the particle and photon distributions.To maximaze efficiency, radiative transfer is usually modelled with a simple photon escape probability formalism assuming isotropic photon and particle distributions. This method can be applied to different, possibly time-dependent, problems for which geometry does not play a crucial role\\footnote{For problems in which geometry is important, radiative transfer can in principle be accounted for by coupling this kinetic code with a radiation transfer solver or a Monte Carlo code  (as demonstrated by \\cite{BL01}), although  computing time may then become a serious issue.}. Within the limits of the one-zone approximation, it is more efficient than other methods and allows for fast data fitting.   The first detailed investigations of high energy plasmas with this technique concentrated on thermal pair plasmas \\citep{FBGPC86,Ghisellini87}. More precise modelling was then proposed in which the particle distributions were decomposed into the sum of a thermal low-energy pool and an arbitrary high energy tail \\citep{LZ87,Svensson87,Coppi92,ZLM93,GHF93,LKL96}. The latter models have been applied most to fitting and interpreting data. They do not however describe the possible deviation from a Maxwellian distribution at low energy, nor do they address explicitly any thermalization process.  Only recent numerical work considered fully arbitrary distributions of particles. \\citet{GHS98} concentrated on the role of synchrotron self-absorbed radiation in AGN. They confirmed previous analytical results \\citep{GGS88} by demonstrating that the exchange of energy between particles by means of synchrotron photons can be an efficient thermalization process in magnetized sources. These simulations focused however on this specific interaction, and other processes were only considered by crude approximations, particularly Compton radiation, or not considered at all. \\citet{NM98} investigated the thermalization of arbitrary distributions by two-body particle interactions and heating by high energy protons. The additional role of synchrotron radiation was not however considered. The most complete numerical treatment of high energy plasmas was probably one developed in the context of $\\gamma$-ray bursts by \\citet{PW05}. Our code is similar to this study but differs in that these authors considered neither particle stochastic acceleration nor the effect of Coulomb losses  The code presented here solves the time-dependent equations simultaneously for isotropic, arbitrary photon, electron, and positron distributions. The evolution of these populations is modelled in time while being affected by self-absorbed cyclo-synchrotron radiation, Compton scattering, pair production/annihilation, e-e and e-p Coulomb collisions, self-absorbed e-p bremsstrahlung radiation, and additional particle acceleration and heating. Each process is described with minimal approximations and by using in most of cases the exact cross sections. For instance, the formulae used for the synchrotron emission and absorption are valid from the sub-relativistic to the ultra-relativistic regime. This numerical strategy allows one to investigate many different astrophysical situations that occur in various high energy sources.  The structure of this paper is as follows. Section~\\ref{sec:rad} provides a description of the microphysics adapted into our code. Then, in Sec.~\\ref{sec:numsol}, we present the numerical techniques. Finally, in Sec.~\\ref{sec:testsandapps} the code is tested against previous published results, providing an overview of its capabilities. ",
        "conclusions": ""
    },
    "0808/0808.3294_arXiv.txt": {
        "abstract": "Results of a 3D MHD simulation of a sunspot with a photospheric size of about 20~Mm are presented. The simulation has been carried out with the MURaM code, which includes a realistic equation of state with partial ionization and radiative transfer along many ray directions. The largely relaxed state of the sunspot shows a division in a central dark umbral region with bright dots and a penumbra showing bright filaments of about $2$ to $3$ Mm length with central dark lanes. By a process similar to the formation of umbral dots, the penumbral filaments result from magneto-convection in the form of upflow plumes, which become elongated by the presence of an inclined magnetic field: the upflow is deflected in the outward direction while the magnetic field is weakened and becomes almost horizontal in the upper part of the plume near the level of optical depth unity. A dark lane forms owing to the piling up of matter near the cusp-shaped top of the rising plume that leads to an upward bulging of the surfaces of constant optical depth. The simulated penumbral structure corresponds well to the observationally inferred interlocking-comb structure of the magnetic field with Evershed outflows along dark-laned filaments with nearly horizontal magnetic field and overturning perpendicular (`twisting') motion, which are embedded in a background of stronger and less inclined field. Photospheric spectral lines are formed at the very top and somewhat above the upflow plumes, so that they do not fully sense the strong flow as well as the large field inclination and significant field strength reduction in the upper part of the plume structures. ",
        "introduction": "The physical understanding of sunspot structure has been hampered for decades by 1) insufficient resolution of the fine structure by observations, 2) lack of information about the layers below the visible surface, and 3) insufficient computational power to perform ab-initio 3D MHD simulation of a full sunspot including the surrounding granulation. Recently, we have seen considerable progress on all three of these fronts: 1) adaptive optics, image selection and reconstruction at ground-based telescopes and the advent of spectro-polarimetry in the visible from space with the {\\em Hinode} satellite have led to a wealth of new information about the fine structure of sunspot umbrae and penumbrae \\citep[e.g.][]{Bharti:etal:2007, Langhans:etal:2007, Ichimoto:etal:2007, Riethmueller:etal:2008, Rimmele:Marino:2006} 2) local helioseismology has started to probe the sub-surface structure of sunspots \\citep[e.g.,][]{Cameron:etal:2008}, and 3) the ever-increasing computational power of parallel computers have made ab-initio simulations of full sunspots come into reach.  While \\citet{Cameron:etal:2007b} simulated solar pores of up to about 3~Mm diameter and did not find indications for the development of a penumbral structure, the first attempt to simulate a sunspot together with the surrounding granulation is due to \\citet{Heinemann:etal:2007}. They considered a rectangular section of a (slab-like) small sunspot of about 4~Mm diameter. The main result of this simulation is the formation of filamentary structures in the outer part of the spot, various properties of which (such as dark cores, outflows, and strongly inclined magnetic field) are consistent with observational results.  However, the filaments found by \\citet{Heinemann:etal:2007} are much shorter than the typical lengths of real penumbral filaments and the overall extension of the simulated penumbra is very small.  Here we report about results of a sunspot simulation with the {\\em MURaM} code \\citep{Voegler:etal:2005}. The simulated sunspot has a total diameter of about 20~Mm, shared about equally by umbra and penumbra. ",
        "conclusions": "The properties of our simulated sunspot are consistent with the general picture of sunspot structure that has emerged from observational studies. This applies to the overall structure (e.g., distinction between umbra and penumbra, average `radial' profiles of the magnetic field components and inclination angle, average outflow in the penumbra) as well as to the detailed properties of the fine structure of umbra and penumbra. The penumbral filamentation results from magneto-convective energy transport in the form of hot rising plumes, very similar to the process giving rise to umbral dots \\citep{Schuessler:Voegler:2006}. The inclined magnetic field near the periphery of the spot causes a symmetry breaking which leads to elongated filaments with strong outflows along flow tubes of nearly horizontal field near optical depth unity. In addition to the flow along the filament, the upflow also turns over into a motion perpendicular to the filament axis. Dark lanes appear above the strongest upflows owing to the upward bulging of the surface of optical depth unity and the piling up of plasma in a cusp-shaped region at the top of the filament, above which the less inclined field outside the filament becomes laterally fairly homogeneous. The horizontal outflows are concentrated along the dark lanes. All these properties are consistent with recent observational results \\citep[e.g.,][]{Bellot-Robio:etal:2005, Rimmele:Marino:2006, Langhans:etal:2007, Ichimoto:etal:2007, Borrero:etal:2008,Noort:Rouppe:2008, Zakharov:etal:2008}.  Our results are also consistent with many properties of the short penumbral filaments found by \\citet{Heinemann:etal:2007}, who gave interpretations along very similar lines. The fact that simulations with two rather different numerical codes lead to basically the same picture for the formation of penumbral structure indicates that, in spite of all differences in detail, the simulations have indeed captured essential physical processes.  The explanation for the Evershed effect as a natural consequence of rising plumes in an inclined field \\citep[cf.][]{Scharmer:etal:2008} connects this flow directly to the basic magneto-convective structure of the penumbra, so that it should occur whenever penumbral structure is present.  The geometry of the flow pattern is such that the observable average outflow velocity need not to be connected with a net outflux of mass. This does not exclude the additional presence of siphon flows \\citep[e.g.,][]{Degenhardt:1993, Montesinos:Thomas:1997, Solanki:etal:1994}, but it is much less clear whether the pressure gradients required for a sustained outward siphon flow are maintained always and everywhere in all penumbrae.  What can be said on the basis of our simulation results about the various models that have been proposed to explain the penumbral structure? First of all, we do not see evidence for the `moving flux tube' model and interchange convection \\citep[e.g.][]{Schlichenmaier:etal:1998b, Schlichenmaier:etal:1998a}. In our simulations, the progression of the filament heads toward the umbra during their formation phase is not caused by the inward motion of a narrow flux tube, but rather due to the expansion of the sheet-like upflow plumes along the filament.  Furthermore, the simulations combine various aspects of the `embedded flux tube' model \\citep{Solanki:Montavon:1993, Borrero:etal:2005, Borrero:etal:2006} and the `gappy penumbra' model \\citep{Spruit:Scharmer:2006,Scharmer:Spruit:2006}. However, the simulated penumbral filaments are neither intrusions of field-free plasma from below nor are they confined to almost horizontal flux tubes disconnected from their environment, particularly in depth. The uppermost part of the plume structure forming the penumbral filament with its strong horizontal flow and almost horizontal field could possibly be represented by a kind of embedded flux tube. The main feature missing in the embedded flux tube models is the overturning convection within the filament and the deep-reaching upflow of plasma that provides the primary energy supply. In this respect, the underlying plume structure with its reduced (albeit non-vanishing) field strength has much similarity with the `gappy' configuration of the penumbra. This scenario captures the convective origin of the penumbral filamentation, even though the gaps form within the strong magnetic field. As a consequence, the gaps contain a horizontal field and, in most cases, are not connected to the almost field free convecting plasma below the penumbra.  There is no clear evidence in our simulations that the penumbral structure is affected or even caused by the fluting instability as suggested by \\citet{Weiss:etal:2004}. However, the periodic boundary condition in the horizontal ($x$) direction used in the simulation implicitly corresponds to the existence of identical sunspots just about 20~Mm from the penumbral boundaries, which certainly affects the field structure, particularly the inclination, in the outer penumbra. Therefore, the simulation possibly does not well represent the convective pumping effect suggested by \\citet{Weiss:etal:2004} as a mechanism for the downward dragging of magnetic flux in the outer penumbra. This might provide a possible explanation for the still rather small extension of the simulated penumbra. Observations in fact indicate that penumbrae are often suppressed on the side of a sunspot which faces a nearby spot of the same polarity.  Altogether, our results indicate a new level of realism in the theoretical modelling of sunspot structure. The properties of the simulated penumbral filaments are consistent with a variety of observational results and provide a basis for a physical understanding of umbral and penumbral structure in terms of magneto-convective processes. On the other hand, there are still clear discrepancies between the numerical results and real sunspots, so that there is some way ahead to be covered towards a completely satisfactory model. Our penumbral structure does not yet appear to be fully evolved and the overall extension of the penumbra is still somewhat small. The average intensity profile indicates that we have simulated the development of an inner penumbra, while the outer penumbra might be more strongly affected by convective pumping \\citep{Weiss:etal:2004}.  The lower boundary condition remains arbitrary since we still have no reliable observational constraints concerning the subsurface structure of sunspots.  Computational limits have forced us to use a rather coarse spatial resolution of 32~km and a fairly small computational box. Furthermore, we could only cover a relatively short overall evolution time. As a consequence, the effective diffusivities in the simulation are still much larger than the real values. Test calculations with different resolution show that 1) first indications for filamentary structure appear already at a horizontal resolution of 96~km and 2) the reduction of field strength in the plumes increases somewhat when we move to a resolution of 24~km.  On that basis, the fundamental physical process of sheet-like plume convection appears to be a robust feature. The results will certainly change in detail (and, hopefully, become even more similar to the observed penumbrae) as resolution increases, but we do not expect totally new processes replacing those that we have described here.  The rapid increase in available computational power and the foreseeable progress in local helioseismology will soon alleviate some of the limitations of the present approach and thus enable us to carry out even more realistic simulations."
    },
    "0808/0808.3777_arXiv.txt": {
        "abstract": "Growth of massive black holes (MBHs) in galactic centers comes mainly from gas accretion during their QSO/AGN phases. In this paper we apply an extended So{\\l}tan argument, connecting the local MBH mass function with the time-integral of the QSO luminosity function, to the demography of MBHs and QSOs from recent optical and X-ray surveys, and obtain robust constraints on the luminosity evolution (or mass growth history) of individual QSOs (or MBHs). We find that the luminosity evolution probably involves two phases: an initial exponentially increasing phase set by the Eddington limit and a following phase in which the luminosity declines with time as a power law (with a slope of $\\sim -1.2$---$-1.3$) set by a self-similar long-term evolution of disk accretion. Neither an evolution involving only the increasing phase with a single Eddington ratio nor an exponentially declining pattern in the second phase is likely. The period of a QSO radiating at a luminosity higher than 10\\% of its peak value is about 2--3$\\times10^8\\yr$, during which the MBH obtains $\\sim 80\\%$ of its mass. The mass-to-energy conversion efficiency is $\\simeq0.16\\pm0.04 ^{+0.05}_{-0}$, with the latter error accounting for the maximum uncertainty due to Compton-thick AGNs. The expected Eddington ratios in QSOs from the constrained luminosity evolution cluster around a single value close to 0.5--1 for high-luminosity QSOs and extend to a wide range of lower values for low-luminosity ones. The Eddington ratios for high luminosity QSOs appear to conflict with those estimated from observations ($\\sim 0.25$) by using some virial mass estimators for MBHs in QSOs unless the estimators systematically over-estimate MBH masses by a factor of 2--4. We also infer the fraction of optically obscured QSOs $\\sim 60-80\\%$. The constraints obtained above are not affected significantly by MBH mergers and multiple-times of nuclear activity (e.g., triggered by multiple times of galaxy wet major mergers) in the MBH growth history. We discuss further applications of the luminosity evolution of individual QSOs to obtaining the MBH mass function at high redshifts and the cosmic evolution of triggering rates of nuclear activity. ",
        "introduction": "\\label{sec:intro}  Massive black holes (MBHs), probably remnants of QSOs \\citep{LyndenBell}, have been detected in the nuclei of many nearby galaxies \\citep{KR95, Magorrian98, Richstone, KG01, FF05}. How do these local MBHs form and evolve, and what is the most important mechanism shaping the mass distribution of MBHs? The current consensus is that the local MBHs obtained their mass mainly through accretion during phases of nuclear activity when they appeared as QSOs/AGNs,\\footnote{Hereafter, we frequently use the term QSOs rather than QSOs/AGNs, if not otherwise specified, to represent QSOs and/or AGNs for convenience.} similar to the ones seen now in the distant universe \\citep[e.g.,][]{YT02,YL04a, Marconi04, Shankar04, Barger05, Hopkins06, Shankar07}.  The evolution of mass accretion onto a MBH is equivalent to the luminosity evolution, given the mass-to-energy conversion efficiency, and is recorded in the luminosity function (LF) of QSOs.  However, the QSO LF depends mainly on two functions: (1) ${\\cal G}(z;\\mbh)$, the rate of nuclear activity triggered at different redshifts $z$ for MBHs with present-day mass $\\mbh$; (2) ${\\cal L}(\\tau;\\mbh)$, the luminosity evolution history of a QSO, of which the remnant MBH has a present-day mass $\\mbh$, as a function of the age of its nuclear activity $\\tau$.  One cannot derive these two functions only from the knowledge of the QSO LF without additional assumptions.  In an extended version of the \\citet{S82} argument, the local MBH mass distribution function (BHMF) is related to QSOs found in the distant universe by the simple integral equation \\begin{eqnarray} \\int^{\\infty}_0 \\Psi_{L}(L,z)\\left|\\frac{dt}{dz}\\right|dz & = & \\int^{\\infty}_{0} n_{\\bh}(\\mbh,t_0) \\times \\nonumber \\\\ &  & \\tau\\life(\\mbh)P(L|\\mbh) d\\mbh, \\label{eq:relation} \\end{eqnarray} where $t_0$ is the present cosmic time, $n_{\\bh}(\\mbh,t_0)$ is the local BHMF, defined so that $n_{\\bh}(\\mbh,t_0)d\\mbh$ gives the number density of local MBHs with present-day mass in the range $\\mbh\\rightarrow \\mbh+d\\mbh$, $\\Psi_{L}(L,z)$ is the QSO LF, defined so that $\\Psi_{L}(L,z)dL$ gives the comoving number density of QSOs with nuclear luminosity in the range $L\\rightarrow L+dL$ at redshift $z$, \\be \\tau\\life(\\mbh)=\\int dL\\sum_{k} \\frac{1}{\\left|\\frac{d{\\cal L}(\\tau;\\mbh)}{d\\tau}|_{\\tau=\\tau_k(L,\\mbh)}\\right|} \\label{eq:taulife} \\ee is the time interval (or the QSO lifetime) in which that a MBH with present-day mass $\\mbh$ appeared as a QSO, and $\\tau_k(L,\\mbh)$ $(k=1,2,...)$ are the roots of the equation ${\\cal L}(\\tau;\\mbh)-L=0$ (see details of the derivation in \\citealt{YL04a}).  Here ${\\cal L}(\\tau;\\mbh)$ represents the luminosity of a QSO and its associated MBH with present-day mass $\\mbh$ at a time $\\tau$ after the triggering of nuclear activity.  The value of $\\tau\\life$ depends on the detailed definition of ``active nuclei'' or the lower threshold set to the nuclear luminosity. Finally, \\be P(L|\\mbh)=\\frac{1}{\\tau\\life(\\mbh)}\\sum_{k} \\frac{1}{\\left|\\frac{d{\\cal L}(\\tau;\\mbh)}{d\\tau}|_{\\tau=\\tau_k(L,\\mbh)}\\right|} \\label{eq:PLMbh} \\ee is the probability distribution function of the nuclear (bolometric) luminosity $L$ over the growth history of the MBH.  The right-hand-side of equation (\\ref{eq:relation}) gives the total time spent per unit $L$ at luminosity $L$ by the progenitors of all the local MBHs in a unit comoving volume, which should be the time integral of the QSO LF, i.e., the left-hand-side of the equation.  Multiplying equation (\\ref{eq:relation}) by the BH mass accretion rate $(1-\\epsilon)L/(\\epsilon c^2)=\\dot\\bh$ (see eqs.~\\ref{eq:mdotinf} and \\ref{eq:mdot} below), where $\\epsilon$ is the mass-to-energy conversion efficiency and $c$ is the speed of light, and then integrating it over cosmic time $t$ reduces to the So{\\l}tan (1982) argument \\citep{YL04a}.  Provided that two basic quantities, i.e., the local BHMF and the QSO LF, can be observationally determined with sufficient accuracy, the kernel $\\tau\\life(\\mbh)P(L|\\mbh)$, containing information on the luminosity evolution history of individual QSOs/MBHs,  may be solved from the integral equation~(\\ref{eq:relation}). Therefore, the extended So{\\l}tan argument is expected to give robust but more detailed constraints on the growth of MBHs than the simple energetic argument due to \\citet{S82}.  As an alternative approach to the theoretical models based on the hierarchical co-evolution of MBHs and galaxies/galactic halos studied intensively in the literature \\citep[e.g.,][]{ER88, HR93, HNR98, KH00, Granato01, WL03, Volonteri03, Croton06, Bower06,Malbon07}, in this paper we use the integral equation (\\ref{eq:relation}) to statistically constrain the growth history of individual MBHs or ${\\cal L}$. The advantages of this approach are: (1) the accretion history of individual QSOs, ${\\cal L}(\\tau;\\mbh)$, is isolated from the triggering rate of nuclear activity, ${\\cal G}(z;\\mbh)$, which is presumably associated with mergers of galaxies or instabilities of galactic disks; and (2) it is free of the many adjustable parameters introduced in the co-evolution models and probably also avoids uncertain assumptions on seed BHs. Note that these two functions,  ${\\cal G}(z;\\mbh)$ and ${\\cal L}(\\tau;\\mbh)$, are mixed in the differential continuity equation for BHMF evolution presented in \\citet[][see also \\citealt{Cavaliere71}, \\citealt{CP89}, and \\citealt{CP90}]{SB92}, which is widely used in studying the growth of MBHs \\citep[e.g.,][]{Marconi04, Shankar07}. Using the luminosity evolution curves, i.e., ${\\cal L}(\\tau;\\mbh)$, obtained from numerical simulations of colliding galaxies, \\citet{Hopkins06} elaborated a unified model for the origin of QSOs and MBHs (see also their other papers listed therein). A possible concern with that approach is that simulations of colliding galaxies have a spatial resolution much larger than the scale of accretion disks around MBHs and therefore may not reflect the real luminosity evolution, as the disk accretion is probably self-regulated in the vicinity of MBHs rather than being directly determined by the material infall rate from a much larger scale or the Bondi-accretion rate (see discussions in \\S~\\ref{sec:models}). (For another model of the possible light curve, see \\citealt{CO07}.)  Estimating the local BHMF can be done with recent advances in observations \\citep[e.g.,][]{Salucci99,AR02,YL04a,Marconi04,Shankar04,Lauer07a,Tundo07}. First, MBHs are believed to exist in the nuclei of most, if not all, nearby galaxies \\citep{KR95,Magorrian98,Richstone,KG01,FF05}. Second, it has been well established that tight correlations exist between the MBH mass and various galactic properties, such as mass, luminosity, stellar velocity-dispersion, light concentration and binding energy of the hot components of galaxies \\citep[here hot components mean either ellipticals or spiral bulges;][]{KR95, Magorrian98, FM00, Gebhardt00, Tremaine02, HR04, MH04, Graham01, AR07}.  Third, the luminosity or velocity-dispersion functions of nearby galaxies have been well determined by large surveys such as the Sloan Digital Sky Survey \\citep[SDSS;][]{Blanton03, Bernardi03, Sheth03}.  Combining the correlation between the MBH mass and galaxy velocity dispersion (or luminosity) with the velocity-dispersion (or luminosity) distribution of nearby galaxies, we estimate the local BHMF in \\S~\\ref{sec:BHMF}.  In the past several years, the QSO LF has been determined over unprecedentedly large luminosity and redshift ranges both from optical surveys such as the Two Degree Field QSO Redshift Survey (2Qz) and SDSS, and from X-ray surveys by ASCA, Chandra and XMM-Newton. For example, the optical QSO LF has been obtained over the redshift range $0.4<z<2.1$ and the magnitude range $M_{\\rm b_J}<-22.5$ using a sample of more than 15,000 QSOs from 2Qz \\citep{Croom04}; \\citet{Richards06a} estimated the QSO LF over a larger redshift range ($0.3<z<5$), but only for bright QSOs, using a homogeneous statistical sample of 15,343 QSOs drawn from SDSS Data Release 3; using the COMBO-17 data, \\citet{Wolf03} estimated the LF for faint QSOs over the range $1.2<z<4.8$; and \\citet{Jiang06} estimated the QSO LF over the range $0.5<z<3.6$ by using a deep survey of faint QSOs in the SDSS.  Obscured (or type 2) QSOs may be missed in the optical surveys but can be detected in hard X-ray surveys. \\citet{LaFranca05} use 508 AGNs to estimate the hard X-ray LF (HXLF; $2-10$~keV) over the range $0<z<2.5$ by combining data from XMM-Newton (Lockman hole) and the Chandra Deep Field (CDF).  \\citet{Barger05} use a spectroscopically complete deep and wide-area Chandra survey to estimate the HXLF ($2-8$~keV) over the range $0<z<5$.  \\citet{Silverman08} measure the HXLF ($2-8$~keV) up to $z\\sim 5$ with fewer uncertainties by combining the observations from the CDF and the Chandra Multiwavelength Project. Combining all these observations, the time integrals of the QSO LF are estimated in \\S~\\ref{sec:QSOLF}.  In \\S~\\ref{sec:models}, we assume several models for the luminosity evolution history of individual QSOs, i.e., ${\\cal L}(\\tau;\\mbh)$, and then apply the models and the observational BHMF and QSO LF to equation (\\ref{eq:relation}) to give constraints on the growth of individual MBHs and the associated parameters, specifically, the efficiency (mainly determined by the spin of a MBH), the lifetime of nuclear activity, and the long-term evolution of disk accretion etc.  We find that a reference model for the luminosity evolution history of individual QSOs, i.e., an initial rapid accretion phase with a rate close to the Eddington limit and then a following power-law declining phase set by the self-similar long-term evolution of disk accretion ($\\dot{\\bh}\\propto \\tau^{-\\gamma}$, and $\\gamma\\sim 1.2-1.3$), can satisfy the extended So{\\l}tan argument (eq.~\\ref{eq:relation}) well. Using the reference model for ${\\cal L}(\\mbh,\\tau)$, we discuss the role of obscuration in the BH growth history in \\S~\\ref{sec:obscuration} and find that obscuration is unlikely to be solely an evolutionary effect. The luminosity (or accretion-rate) evolution constrained by the extended So{\\l}tan argument also implies a distribution of Eddington ratios (i.e., the accretion rate in units of the Eddington limit) in QSOs. In \\S~\\ref{sec:accrete}, we particularly discuss its distribution expected from the models and compare them with observations.  In \\S~\\ref{sec:BHmerger}, by using toy models, we discuss the effects of BH mergers on our results, which are shown to be insignificant. In \\S~\\ref{sec:discon}, we discuss further implications of the luminosity evolution obtained from the extended So{\\l}tan argument. Together with the QSO LF, ${\\cal L}(\\tau;\\mbh)$ can be used to further derive the BHMF at redshift $z$ and the triggering rate of nuclear activity ${\\cal G}(z; \\mbh)$.  Given ${\\cal L}(\\tau;\\mbh)$ and ${\\cal G}(z;\\mbh)$, many statistical properties of QSOs can be inferred and comparison of them  with observations may further deepen our understanding of the growth of MBHs.  Conclusions are given in \\S~\\ref{sec:conclusion}.  In this paper we set the Hubble constant as $H_0=100h\\kms\\Mpc^{-1}$; and if not otherwise specified, the cosmological model used is $(\\Omega_m, \\Omega_{\\Lambda},h)=(0.3,0.7,0.7)$. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we have studied the observational constraints on the growth of MBHs using the extended So{\\l}tan argument. In this approach, the local BHMF is directly connected with the time-integral of the QSO LF through only the luminosity evolution of individual QSOs (and correspondingly the accretion-rate evolution, given the mass-to-energy conversion efficiency), and the luminosity evolution of individual QSOs is isolated from the cosmic evolution of the triggering rate of nuclear activity.  The luminosity (or accretion-rate) evolution of individual QSOs has an unambiguous physical definition, and it is different from the `mean accretion rate' as a function of mass and/or redshift widely used in the literature \\citep[e.g.,][]{Marconi04, Shankar04, Shankar07} in that the `mean accretion rate' is a combined property depending on both the luminosity evolution of individual QSOs and the cosmic evolution of the triggering rate of nuclear activity.  With recent knowledge of the relationships between MBH mass and host galaxy properties (either the $\\mbh-\\sigma$ relation or the $\\mbh-L_{\\rm bulge}$ relation) and the distribution of galaxy properties (either $\\sigma$ or $L_{\\rm bulge}$), we estimate the local BHMF. We obtain the time-integral of the QSO LF from recent estimates of QSO LFs in both optical and X-ray bands. Using the local BHMF and the time-integral of the QSO LF, we obtain robust constraints on the luminosity (or accretion rate) evolution of individual QSOs and important characteristic parameters describing the growth of individual MBHs, such as the mass-to-energy conversion efficiency and lifetime, which are summarized below.  \\begin{itemize}  \\item The luminosity (or accretion rate) evolution of individual QSOs probably involves two phases: an initially exponentially increasing phase set by the Eddington limit (i.e., ${\\cal L}\\simeq L\\Edd$) when the infall material to feed the central MBHs is over-supplied; and then followed by a phase with power-law declining  set by a self-similar long-term evolution of disk accretion (i.e., ${\\cal L}\\propto \\tau^{-\\gamma}$ and $\\gamma\\sim 1.2-1.3$).  With this type of luminosity evolution, the time-integral of QSO LF can be well matched by that inferred from the local BHMF. Other simple luminosity evolution models, such as a single Eddington ratio for all MBHs/QSOs or an initially exponentially increasing phase followed by an exponentially decay phase, cannot satisfy the extended So{\\l}tan argument simultaneously at both the high-luminosity end and low-luminosity ends, and thus are ruled out.  \\item The mass-to-energy conversion efficiency $\\epsilon$ is $\\simeq 0.16\\pm 0.04^{+0.05}_{-0}$ (correspondingly the  spin parameter $a$ is in the range from $0.8$ to $0.99$) if adopting the local BHMF estimated from the $\\mbh-\\sigma$ relation, which is fully consistent with the theoretical expectations of $\\sim 0.10-0.20$, i.e., the spin of MBHs in QSOs may stay at an equilibrium of $\\sim 0.7-0.9$ for most of the QSO lifetime \\citep[e.g.,][]{Gammie04, Shapiro05, HBK07}. However, the efficiency $\\epsilon$ is reduced to $\\simeq 0.08 \\pm 0.02^{+0.03}_{-0}$ (and correspondingly the spin parameter $a$ is in the range from $0.1$ to $0.8$) if adopting the local BHMF estimated from the $\\mbh-L_{\\rm bulge}$ relation, which is lower than but may be still marginally consistent with theoretical expectations.  \\item The lifetime of QSOs/AGNs, which depends on detailed definition of the nuclear activity or the lower threshold set to the active nuclear luminosity, can be as long as a few $10^9$~yr, and the characteristic timescale in the luminosity increasing phase and transition timescale to the declining phase do not necessarily depend on the mass of their central MBHs. The period that a QSO or MBH radiating at a luminosity larger than $10\\%$ of its peak luminosity is only about $2-3\\times 10^8$~yr, and during this period the MBH obtained most of its mass. If adopting the local BHMF estimated from the $\\mbh-L_{\\rm bulge}$ relation, the above values related to the QSO lifetime decrease by a factor $\\sim 2$.  \\item For individual QSOs, the characteristic timescale for the luminosity (or accretion rate) to decline from its peak $L_{P}$ to $0.1L_{P}$ in their second phase should be relatively short compared to the Salpeter timescale, which suggests that the material infalling from a large galactic scale and deposited in the vicinity of MBHs be consumed by rapid accretion onto the central MBH and at the mean time further deposit of material can be efficiently suppressed by some mechanisms, probably the AGN feedback mechanism \\citep[e.g.,][]{SR98, King, Murray05, DSH05}, on a timescale $\\la \\taus$.  \\item The majority of high-luminosity ($L_{\\rm bol}\\ga 10^{45.75}\\ergs$) QSOs accrete material via an almost single Eddington ratio, close to 1 and not smaller than half of the Eddington limit, which suggests that the disk accretion onto MBHs should be indeed self-regulated by the Eddington limit when the accretion material is over-supplied in the initial phase.  \\item Low-luminosity QSOs ($L_{\\rm bol}\\la 10^{45.25}\\ergs$) accrete material via a much wider range of Eddington ratios, and a significant fraction of them accrete material via low Eddington ratio ($\\dot{m}\\la 0.1$), which corresponds to the self-similar long-term evolution of disk accretion around MBHs ($\\bh\\sim \\tau^{-\\gamma}$ and $\\gamma\\sim 1.2-1.3$) when the accretion material is under-supplied.  \\item The Eddington ratio distribution among QSOs/AGNs inferred from the extended So{\\l}tan argument concentrates toward high Eddington ratios (close to 1), especially for high luminosity QSOs, which appears to conflict with that estimated directly from observations (with a mean value of $\\sim 10^{-0.6}-10^{-1.1}$) by using the virial mass estimator(s). To make these two distributions consistent with each other, an offset of $0.3-0.6$~dex in the MBH mass estimated from the virial mass estimator(s) is required, at least for high-luminosity QSOs, which suggests that MBHs masses obtained from the virial mass estimator(s), have been systematically over-estimated by a factor of $2-4$.  \\item The fraction of optically obscured QSOs/AGNs inferred from the extended So{\\l}tan argument can be as high as $80\\%$ at $M_{\\rm B}\\sim -20$---$-23$ and slightly decreases to $60\\%$ at $M_{\\rm B}=-24$---$-27$, and these numbers are consistent with recent observations by \\citet{Reyes08}. The dependence of the fraction of type 2 AGNs on the X-ray luminosity cannot be solely an evolutionary effect arising from their individual evolution after nuclear activity is triggered (i.e., QSOs are more likely to be obscured in the early stage of the MBH growth), and some other effects (e.g., those introduced by the receding torus model; \\citealt{Lawrence91}) should be responsible for the larger fraction of type 2 AGNs at low luminosities ($L_{X}\\la 10^{43.5}\\ergs$).  \\end{itemize}  We estimate possible effects due to MBH mergers (which may re-shape the local BHMF) and multiple times of nuclear activity and accretion (e.g., triggered by multiple times of galaxy `wet' major mergers) in the growth history of a MBH, and we find that these effects on our conclusions are insignificant, which again supports that the constraints obtained above are robust.  The constraints on the luminosity evolution of individual QSOs obtained from the extended So{\\l}tan argument in this paper, together with the QSO LF, can be further used to derive the BHMF at high redshifts and the cosmic evolution of the triggering rate of nuclear activity. These constraints and those recent estimates on the Eddington ratio distribution in QSOs ask for serious theoretical modeling of the long-term evolution of disk accretion around MBHs. More detailed modeling of accretion and radiation transfer physics in the vicinity of MBHs may have to be involved to determine the self-regulation of the disk accretion (rather than the simple Eddington limit argument) at the initial phase with sufficient deposited accretion material. It should be one of the important long-term goals for theoretical studies on the growth of MBHs to answer questions like what determines the transition from the initial rapid accretion phase with Eddington ratio close to 1 to the rapid declining phase, what shuts off the efficient accretion process around MBHs (probably jointly determined by an efficient feedback mechanism and the accretion disk viscosity), and what determines the evolution of the spin of MBHs."
    },
    "0808/0808.4157_arXiv.txt": {
        "abstract": "We examine variability of absorption line strength of intervening systems along the line of sight to GRB 060206 at $z = 4.05$, by the low-resolution optical spectra obtained by the Subaru telescope from six to ten hours after the burst. Strong variabilities of Fe~\\emissiontype{II} and Mg~\\emissiontype{II} lines at $z=1.48$ during $t=$5--8 hours have been reported for this GRB \\citep{Hao07}, and this has been used to support the idea of clumpy Mg~\\emissiontype{II} cloudlets that was originally proposed to explain the anomalously high incidence of Mg~\\emissiontype{II} absorbers in GRB spectra compared with quasars. However, our spectra with higher signal-to-noise ratio do not show any evidence for variability in $t=$6-10 hours. There is a clear discrepancy between our data and Hao et al. data in the overlapping time interval.  Furthermore, the line strengths in our data are in good agreement with those observed at $t \\sim$ 2 hours by \\citet{Tho08}.  We also detected Fe~\\emissiontype{II} and Mg~\\emissiontype{II} absorption lines for a system at $z = 2.26$, and these lines do not show evidence for variability either. Therefore we conclude that there is no strong evidence for variability of intervening absorption lines toward GRB 060206, significantly weakening the support to the Mg~\\emissiontype{II} cloudlet hypothesis by the GRB~060206 data. ",
        "introduction": "The gamma-ray burst (GRB) has opened up a new window to probe high redshift universe, now reaching comparable redshifts with the studies by galaxies and quasars, and starting to probe the era of the cosmic reionization \\citep{Kawai06}.  The optical GRB afterglow spectra give us a lot of information about GRB host galaxies \\citep{Mir02, Sav03, Prochaska07, Chen07}, intervening absorption systems along the line of sight \\citep{Proch06, Sud07, Tej07} and the reionization of the universe \\citep{Tot06}.  GRB~060206 is the 10th highest redshift GRB as of 2008 November.  GRB~060206 was triggered by the $Swift$ Burst Alert Telescope on 2006 February 6 at 04:46:53 UT.  The redshift of the burst was determined to be 4.048 \\citep{Fyn06}.  The intervening absorption lines are reported at $z=1.48$ \\citep{Fyn06,Hao07,Tho08} and $z=$2.26 (Aoki et al. 2006; this work).  Surprisingly, \\citet{Hao07} claimed strong variation of the equivalent widths (EWs) of Fe~\\emissiontype{II} $\\lambda 2600.2$ and Mg~\\emissiontype{II} $\\lambda2803.5$ at $z$=1.48; the EW of the Fe~\\emissiontype{II} $\\lambda 2600.2$ system at $z=1.48$ showed a more than 80\\% decrease within 1.5 hours duration. It is difficult to interpret this system to be located in the GRB host galaxy, because a large redshift difference requires relativistic motion and hence a large velocity dispersion of lines is expected, while the absorption line widths are only $14 - 20$ km s$^{-1}$ \\citep{Tho08}.  Then, almost unique possible interpretation is that the intervening absorption system is smaller than the image size of the GRB afterglow, and the line strength changes with the size evolution of the GRB afterglow. This requires the size of the absorption system to be less than $\\sim 10^{16}$ cm \\citep{Hao07}.  This variability, if real, would change the standard picture of intervening absorption systems found in bright cosmological objects such as quasars and GRBs. Furthermore, this result has been used to support the Mg~\\emissiontype{II} cloudlet hypothesis that has been proposed to explain the anomalous Mg~\\emissiontype{II} incidence along the GRB sight lines \\citep{Frank07}. \\citet{Proch06} reported that the incidence of Mg~\\emissiontype{II} absorption lines along GRB sight lines is about four times higher than that for quasars, at a statistical significance greater than 99.9\\% confidence limit, and this cannot easily be explained by known selection effects.  If the intervening line variability of GRB 060206 is in fact due to the small size of the Mg~\\emissiontype{II} cloudlets, the anomalous incidence of Mg~\\emissiontype{II} systems may also be explained by the larger image size of quasars than GRB afterglows.  However, GRB 060206 is currently the only one case for the intervening line variability, and it should be examined carefully.  Here we report the analysis of the optical afterglow spectra of GRB 060206 taken by the Subaru telescope during $t = $ 6.01--9.73 hours after the burst, a part of which is overlapping with the Hao et al.'s observation ($t = $ 4.13--7.63 hours).  We will examine especially the line variability of the intervening Mg~\\emissiontype{II} and Fe~\\emissiontype{II} systems. In addition to the $z=1.48$ system, we will give detailed analysis of the $z=2.26$ system as well and examine the variability.  When we completed most of this paper, the revised version of Th\\\"one et al. (2008, arXiv:0708.3448v2) appeared on the preprint server, in which they also analyzed the same data from the Subaru archive and reached a similar conclusion of no evidence for variability. However, details of the Subaru observation and data analysis are not given in their paper. Here we give more detailed descriptions about the observation logs and the processing of the Subaru data taken by ourselves, as well as the implications for the Mg~\\emissiontype{II} incidence problem.  The $z=2.26$ system was originally reported by Aoki et al. (2006) and it is mentioned also in Fynbo et al. (2006), but Th\\\"one et al. (2008) concluded that this is not real. However, we will show that this system is real, based on the detection of seven Fe~\\emissiontype{II} lines and Mg II doublet. ",
        "conclusions": "In addition to the Subaru data, the equivalent width measurements of the $z = 1.48$ system by \\citet{Hao07} and the WHT/ISIS data of \\citet{Tho08} are shown in Figure \\ref{fig4}.  Our early four spectra are overlapping in time with the last three observations by \\citet{Hao07}.  Our EW measurement of $z=1.48$ Fe~\\emissiontype{II} $\\lambda2600.2$ at $t = $ 6.011 hours is clearly in disagreement with the value measured by \\citet{Hao07} at $t= $ 6.15 hours.  It should be noted that Fe~\\emissiontype{II} $\\lambda2586.7$ at $z=1.48$ is clearly detected at 6415 \\AA~ in our spectra (Figure \\ref{fig2}) with an equivalent width of $\\sim 0.55$ \\AA.  On the other hand, it is not mentioned by \\citet{Hao07}, and in fact, this line cannot be clearly seen in the Hao et al.'s spectra.  This indicates that an EW of $\\lesssim 0.5$ \\AA~ cannot reliably be measured by the Hao et al.'s observation, and hence the claim of the variability of Fe~\\emissiontype{II} $\\lambda2600.2$ line should suffer from a large uncertainty.  The Hao et al.'s spectra show a strong variability of Mg~\\emissiontype{II} lines around $t = 5 - 5.5$ hours, and we cannot examine this variability because this epoch is not covered by our data. However, the early measurements of EWs of these lines at $z = 1.48$ by \\citet{Tho08} at $t \\sim 2$ hours are in good agreement with our measurements. It seems rather unlikely that the lines show a variability only in the time duration of the Hao et al.'s observation and no variability at all in the earlier or later data with higher signal-to-noise ratio.  Furthermore, there is no evidence for time evolution of EWs in another absorption system at $z = 2.26$.  From these results, we conclude that the variability reported by \\citet{Hao07} is likely due to some artificial effects or statistical fluke in low signal-to-noise data.  Our result implies that the data of GRB 060206 no longer provide a strong support to the Mg~\\emissiontype{II} cloudlet hypothesis proposed by \\citet{Frank07} for the anomalously high incidence of Mg~\\emissiontype{II} systems in GRB spectra.  Furthermore, several difficulties of this hypothesis have been pointed out.  One such argument is that apparently unsaturated Mg~\\emissiontype{II} absorption lines with the doublet ratio of 1:1 are never found in intervening absorption lines, although such lines are expected if optically thick Mg~\\emissiontype{II} cloudlets partially cover quasar beams \\citep{Proch06, Porc07, Sud07}.  Another argument is that there is no apparent change in Mg~\\emissiontype{II} incidence over the continuum and broad emission lines of quasar spectra, although the broad line regions of quasars are much larger than the continuum emitting region \\citep{Pont07}.  It should also be noted that the suggested size of the Mg~\\emissiontype{II} cloudlets is several orders of magnitude smaller than those derived from the observations of gravitationally lensed quasars \\citep{Rau02, Ell04}. \\par Therefore, we consider that the Mg~\\emissiontype{II} incidence anomaly is likely to be explained by some other effects.  Although the anomaly has not yet been solved, it could be explained, or at least the statistical significance of the anomaly is reduced, by a combination of some selection effects such as dust extinction and gravitational lensing \\citep{Porc07, Sud07}.   \\bigskip We are grateful to the Subaru Telescope staff for their assistance during our observations.  This work was partly supported by the Grant-in-Aid for Scientific Research on Priority Areas (19047003) from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan.  \\appendix"
    },
    "0808/0808.1634_arXiv.txt": {
        "abstract": "{} {The first-order ordinary differential equation (ODE) that describes the mid-plane gravitational potential in flat finite size discs of surface density $\\Sigma(R) \\propto R^{s}$ (Hur\\'e \\& Hersant 2007) is solved exactly in terms of infinite series.} {The formal solution of the ODE is derived and then converted into a series representation by expanding the elliptic integral of the first kind over its modulus before analytical integration.} {Inside the disc, the gravitational potential consists of three terms: a power law of radius $R$ with index $1+s$, and two infinite series of the variables $R$ and $1/R$. The convergence of the series can be accelerated, enabling the construction of reliable approximations. At the lowest-order, the potential inside large astrophysical discs ($s \\sim -1.5 \\pm 1$) is described by a very simple formula % whose accuracy (a few percent typically) is easily increased by considering successive orders through a recurrence. A basic algorithm is given. } {Applications concern all theoretical models and numerical simulations where the influence of disc gravity must be checked and/or reliably taken into account.} {} ",
        "introduction": "Gaseous discs in which the main physical quantities (density, pressure, temperature, thickness, velocity) scale with cylindrical radius as power laws, i.e. ``power-law discs'', represent an important class of theoretical systems. These are used customary to model accretion in evolved binaries \\citep{ss73,pringle81}, circumstellar matter \\citep{dubrulle92,edgar07}, the environment of massive black holes inside active galactic nuclei \\citep{cd90,hure98,semerak04} or even the stellar component of some galaxies \\citep{evans94,zhao99}. In most applications however, power-law discs are truncated either to avoid diverging values at the disc centre (such as density, mass) or in attempting to reproduce the properties of observed discs of finite extension and mass. Although self-similarity is not compatible with the presence of edges, it is generally considered that power laws offer a good description of disc properties in some regions (far from the edges).  Note that the presence of sharp edges can be misleading when interpreting observational data \\citep[e.g.][]{hughes08}. In general, the surface density $\\Sigma$ in the outer parts of discs is a decreasing function of the cylindrical radius $R$. Depending on the models, hypotheses, and objects, we have, for instance, $\\Sigma \\propto R^{-3/5}$ in binaries  \\citep{ss73}, $\\Sigma \\propto R^{\\pm9/20}$ in active galactic nuclei \\citep{cd90}, $\\Sigma \\propto R^{-1}$ for a Mestel disc \\citep{mestel63}, or $\\Sigma \\propto R^{-3/2}$ in circumstellar discs \\citep{pietu07}. In the context of stationary viscous $\\alpha$-discs, a wide range of power-law exponents is allowed since the temperature $T$, $\\Sigma$, and $R$ satisfy the condition \\citep[e.g.][]{pringle81}: $$\\Sigma T R^{3/2} =cst,$$ while it is $\\Sigma \\propto R^{-1/2}$ in $\\beta$-discs \\citep{hrz01}.\\\\  The calculus of the gravitation potential of finite-size, power-law discs has received little attention yet. Several reasons can be put forward. Solving the Poisson equation or computing the integral of the potential is not a trivial procedure, especially in the presence of edges. It is generally believed that gravity due to low mass discs is unimportant compared with that of a central proto-star or black hole, and cannot be probed \\citep[see however][]{baruteaumasset08}. Many studies employ the multi-pole expansion which is known to converge too slowly inside sources to be efficient for the numerical applications \\citep[e.g.][]{clement74,stonenorman92}. \\cite{hh07} demonstrated that the mid-plane potential of flat power-law discs obeys an inhomogeneous first-order Ordinary Differential Equation (ODE). In this second paper, we discuss the exact solutions of this ODE for the entire physical range (outside and inside the disc) in terms of infinite series. In particular, it is shown that the mid-plane potential is a combination of a power law for the radius $R$ and two series of the variables $R$ and $1/R$. Since these series converge rapidly inside large discs, it is possible to derive reliable approximations by truncating the series at low orders.\\\\  This paper is organised as follows. The ODE for the potential is briefly recalled in Sect. \\ref{sec:unified} and its formal solution is derived in Sect. \\ref{sec:Formal solution of the ODE}. In Sect. \\ref{sec:pate}, we express the potential at the two disc edges and consider a few special cases. The inside and outside solutions in the form of series are presented in Sect. \\ref{sec:Solution of the ODE}. In Sect. \\ref{sec:Potential inside the disc}, we analyse the potential in the disc inside in detail, and in particular, the power-law contribution. Since all series involved converge rapidly, we are able to derive reliable approximations for the potential; this is done in Sect. \\ref{sec:Approximate formulae}. We discuss in Sect. \\ref{sec:Discs with no inner/outer edges} the case of discs with no inner and/or outer edge. The paper ends with a few concluding remarks.  \\begin{figure}[h] \\includegraphics[width=9.0cm]{fig1hure09682.eps} \\caption{Configuration for a finite-size flat disc.} \\label{fig:scheme.eps} \\end{figure} ",
        "conclusions": "In this paper, we have derived an exact expression for the gravitational potential in the plane of flat power-law discs as a solution of the ODE reported in Paper I. This expression is valid over the entire spatial domain and takes into account finite size effects. Inside the disc (the most difficult case to treat in general), it consists of three terms of comparable magnitude: a power law of the cylindrical radius $R$ with index $1+s$ (where $s$ is the exponent of the surface density) and two series of $R$ and $1/R$. In terms of convergence, our expression is by far superior to the multi-pole expansion method. Reliable approximations for the potential can be produced by performing fully controlled truncations. We have shown that the potential can be expressed by means of a simple function of $R$ and $s$, which is valid to within a few percents in the range of exponents $-3 \\lesssim s \\lesssim 0$. This formula should be sufficiently accurate for most astrophysical applications. If necessary, more accurate formulae can be developped by including successive terms. These results should help in investigating various phenomena where disc gravity plays a significant role.  An interesting point concerns the case of discs for which the surface density is not a power-law function. As shown in Paper I, it is easy to reproduce numerically the potential when the profile $\\Sigma(R)$ is a mixture of power laws. From an analytical point of view however, the construction of a reliable formula for $\\psi$ as compact as the one obtained here is not guaranteed at all. For instance, for an expansion of the form: \\begin{equation} \\Sigma(R)=\\sum_{s=0}^M{\\alpha_s R^s}, \\end{equation} where $s$ is an integer, each of the $M+1$ series $\\{a_n,b_n\\}$  should be truncated at a rank $N \\gtrsim n_\\Delta \\sim (s+1)/2$ (see Eqs. (\\ref{eq:psifinal}) and (\\ref{eq:psiapprox})), which corresponds to an approximate formula for $\\psi$ containing about $2M(M+1)$ terms. The number of terms to consider can become prohibitively large when several power laws are required."
    },
    "0808/0808.3407_arXiv.txt": {
        "abstract": "We investigate the close environment of 203 {\\em Spitzer} \\tfm-selected sources at $0.6<z<1.0$ using zCOSMOS-bright redshifts and spectra of $I<22.5$ AB mag galaxies, over 1.5 deg$^2$ of the COSMOS field.   We quantify the degree of passivity of the LIRG and ULIRG environments by analysing the fraction of close neighbours with $\\rm D_n \\,(4000)>1.4$. We find that LIRGs at $0.6<z<0.8$ live in more passive environments than those of other optical galaxies that have the same stellar mass distribution. Instead, ULIRGs inhabit more active regions (e.g. LIRGs and ULIRGs at $0.6<z<0.8$ have, respectively, $(42.0 \\pm 4.9)\\%$ and $(24.5 \\pm 5.9)\\%$ of  neighbours with $\\rm D_n \\,(4000)>1.4$ within 1~Mpc and $\\pm \\, 500 \\, \\rm km \\, s^{-1}$). The contrast between the activities of the close environments of LIRGs and ULIRGs appears  especially enhanced in the COSMOS field density peak at $z\\sim0.67$, because LIRGs on this peak  have a larger fraction of passive neighbours, while ULIRGs have as active close environments  as those outside the large-scale structure.  The differential environmental activity is related to the differences in the distributions of stellar mass ratios between LIRGs/ULIRGs and their close neighbours, as well as in the general local density fields. At $0.8<z<1.0$, instead, we find no differences in the environment densities of ULIRGs and other similarly massive galaxies, in spite of the differential activities. We discuss a possible scenario to explain these findings. ",
        "introduction": "\\label{sec-intro}  Our knowledge of the nature of the sources composing the extragalactic IR background (Puget et al.~1996; Dole et al.~2006) has been much clarified over the last years thanks to studies conducted with the {\\em Spitzer Space Telescope} (Werner et al.~2004) and now also the {\\em Akari Telescope} (Matsuhara et al.~2006). Many properties of these galaxies are now quite well known, including their redshift distribution (e.g. Rowan-Robinson et al.~2005; Caputi et al.~2006a; Desai et al.~2008); luminosity evolution (e.g. Le Floc'h et al.~2005; P\\'erez-Gonz\\'alez et al.~2005; Babbedge et al.~2006; Caputi et al.~2007);   assembled stellar masses (Daddi et al.~2005; Caputi et al.~2006a,b; Papovich et al.~2006; Noeske et al.~2007), and multiwavelength spectral characteristics (e.g. Houck et al.~2005; Yan et al.~2005; Choi et al.~2006; Weedman et al.~2006; Berta et al.~2007; Lacy et al.~2007; Sajina et al.~2007; Caputi et al.~2008; Rigby et al.~2008).   The properties of the environment in which IR-selected galaxies reside are less well-determined. This is  mainly because environmental studies require the combination of wide area surveys with accurate redshift determinations  that, in general, are only provided  by large spectroscopic samples. Recent works have explored the link between environment and star-formation activity in IR-selected galaxies. Elbaz et al.~(2007) determined that the star-formation rate (SFR)-density relation observed locally was reversed at $z\\sim1$. Serjeant et al.~(2008) suggested that this reversed relation could also hold for sub-millimetre galaxies at $1.0<z<1.5$.  Marcillac et al. (2008) found that luminous and ultra-luminous IR galaxies (LIRGs and ULIRGs, respectively) at $0.7<z<1.0$ are present, on average, in denser environments than those of other optical galaxies, but they argued that this difference is mainly driven by stellar mass.    The Cosmic Evolution Survey (COSMOS; Scoville et al.~2007)  has been designed to probe galaxy evolution and the effects of environment up to high redshifts over 2 deg$^2$ of the sky.  COSMOS observations include the coverage of the field with the Multiband Photometer for {\\em Spitzer} (MIPS; Rieke et al.~2004) as part of two  {\\em Spitzer} Legacy Programs (S-COSMOS; Sanders et al. 2007). COSMOS also comprises a large spectroscopic follow-up (zCOSMOS; Lilly et al.~2007) with the  Visible Multiobject Spectrograph (VIMOS; Le F\\`evre et al.~2003) on the {\\em Very Large Telescope (VLT)}.  The combination of the S-COSMOS and zCOSMOS datasets are ideal to characterise the environment of IR-selected galaxies and investigate the properties of their close neighbours.  In this paper, we make use of zCOSMOS-bright spectra to  investigate the characteristics of the environment of the most luminous \\tfm-selected galaxies with $I<22.5$ AB mag at redshifts $0.6<z<1.0$.  Firstly, in \\S\\ref{sec_closeenv}, we analyse the activity and stellar masses of galaxies which are close to \\tfm galaxies of different IR luminosities. Secondly, in \\S\\ref{sec_denspeak},  we particularly study the activity of close neighbours of IR galaxies lying on the COSMOS field density peak at $z\\sim0.67$.  Finally, in \\S\\ref{sec_densf}, we characterise the environment of \\tfm sources using the general density fields produced with zCOSMOS (Kova\\v{c} et al.~2008). We adopt throughout a cosmology with $\\rm H_0=70 \\,{\\rm km \\, s^{-1} Mpc^{-1}}$, $\\rm \\Omega_M=0.3$ and $\\rm \\Omega_\\Lambda=0.7$. ",
        "conclusions": "\\label{sec_concl}   We conclude that LIRGs and ULIRGs at $0.6<z<0.8$ have close (within $2 \\, \\rm Mpc$ and $\\pm \\, 500 \\, \\rm km \\, s^{-1}$) environments characterised by significantly different degrees of activity: the fraction of `passive' galaxies around LIRGs is substantially higher than the fraction around ULIRGs. Remarkably, the differential activity is especially enhanced in the COSMOS density field peak at $z\\sim0.67$. By analysing control samples with the same stellar mass distributions as the two respective populations of IR galaxies, we show that the differences in environment activity are not an effect of the central galaxy stellar mass.  Instead, this is probably a consequence of the IR-galaxy density fields. LIRGs and ULIRGs at $0.6<z<0.8$ are surrounded by close neighbours of different stellar masses, and their general local over-densities are also different. ULIRGs inhabit under-dense environments, while LIRGs  prefer over-denser regions. Thus, the differential activities observed in the  environments of the two populations are related to their density fields.  At $0.8<z<1.0$, we also find that ULIRGs reside in more active close ($< 1 \\, \\rm Mpc$) environments than those of other optical galaxies with similar stellar masses. However, we do not find any difference between the stellar mass distributions of their close neighbours. This is in contrast with the environmental properties of ULIRGs at $0.6<z<0.8$.  At $0.8<z<1.0$, the similarities of the environments in which ULIRGs and other equally massive galaxies live suggest that these other galaxies might have also passed by a ULIRG phase at some moment in the past (consistently with the high incidence of ULIRGs among massive galaxies at $z \\gsim 1.5$ found by Daddi et al.~2005 and Caputi et al.~2006b). The quenching of the IR phase in these galaxies could have happened along with the cessation of star formation in some of their close neighbours, raising the fraction of inactive sources in the surroundings.  The number density of ULIRGs has a rapid decline at $z<1$ and, in particular, it reduces to around a half in the 1 Gyr of elapsed time between redshifts $z\\sim0.9$ and $z\\sim0.7$ (Le Floc'h et al.~2005; Caputi et al.~2007). This evolution occurs together with a change of environments, suggesting that the conditions under which ULIRGs occur are not the same at $z\\sim0.9$ and $z\\sim0.7$.  ULIRGs at $z\\sim0.9$ are found in similar kinds of sites as other equally massive galaxies,  while those at $z\\sim0.7$ appear mostly restricted to regions where they are the dominant galaxy of their close neighbourhood. It is plausible that the ULIRGs we observe at $z\\sim0.7$ have been produced after the merging of close galaxies, leaving the final ULIRG close neighbourhood empty of massive neighbours. At $z\\sim0.9$, even when mergers could also have been a possible mechanism for triggering the ULIRG phase, it seems that there has been no need to consume most of the close environment for these ULIRGs to be produced."
    },
    "0808/0808.0068_arXiv.txt": {
        "abstract": "An overview is given of the chemical processes that occur in primordial systems under the influence of radiation, metal abundances and dust surface reactions. It is found that radiative feedback effects differ for UV and X-ray photons at any metallicity, with molecules surviving quite well under irradiation by X-rays. Starburst and AGN will therefore enjoy quite different cooling abilities for their dense molecular gas. The presence of a cool molecular phase is strongly dependent on metallicity. Strong irradiation by cosmic rays ($>200\\times$ the Milky Way value) forces a large fraction of the CO gas into neutral carbon. Dust is important for H$_2$ and HD formation, already at metallicities of $10^{-4}-10^{-3}$ solar, for electron abundances below $10^{-3}$. ",
        "introduction": "In the study of primordial chemistry, and subsequently the formation of the first stars, it is crucial to understand the ability of interstellar gas to cool through atomic and molecular emissions and to collapse. Furthermore, atomic and molecular species allow for probes of the ambient conditions, like density and temperature, under which stars form. The basic questions of this contribution are: What sets the abundances of atoms and molecules that cool gas? What is the role of radiation, metallicity and dust in molecule formation?  A number of different chemical processes are relevant to this effect:  \\smallskip Ion-molecule reactions: A$^+$ + BC$\\rightarrow$AB$^+$ + C;  Neutral-neutral reactions: A + BC$\\rightarrow$AB + C;  Dissociative recombination: AB$^+$ + e$^-$$\\rightarrow$A + B;  Radiative recombination: A$^+$ + e$^-$$\\rightarrow$A + $h\\nu$;  Radiative association: A + B$\\rightarrow$AB + $h\\nu$;  Ionization: A + CR/UV/X-ray$\\rightarrow$A$^+$ + e$^-$;  Dissociation: AB + UV$\\rightarrow$A + B;  Charge transfer: A$^+$ + B$\\rightarrow$A + B+;  Grain surface reactions: Grain + A + B$\\rightarrow$Grain + AB. \\smallskip  In any chemical network, the above reactions play an important role. For example, the charge transfer between H$^+$ and O, followed by reactions with H$_2$ to H$_3$O$^+$, and dissociative recombination with e$^-$, leads to species like OH and H$_2$O (following certain branching ratios). Similarly complex routes exist for CO. In any case, many species are typically joined through different chemical routes. Thus, it is not trivial to construct concise chemical networks if one wants to include important molecules like CO and H$_2$O\\footnote{It is important to realize that water can be quite an important heating agent in the presence of a warm infrared background, like $T>50$ K dust or a $z>15$ CMB (Spaans \\& Silk 2000).}. Of course, in the limit of low metallicity, chemistry simplifies. Basically, no metals implies no molecules except for H$_2$ and HD (and a few minor species). Still, even small amounts of metals and dust ($\\sim 10^{-4}$ solar) can be crucial to the efficient formation of species like H$_2$, HD, CO, H$_2$O and many others, which is the purpose of this contribution. ",
        "conclusions": "Radiative feedback effects differ for UV and X-ray photons at any metallicity, with molecules surviving quite well under irradiation by X-rays. Starburst and AGN will therefore enjoy quite different cooling abilities for their dense molecular gas. The presence of a cool molecular phase is strongly dependent on metallicity. Strong irradiation by cosmic rays ($>200\\times$ the Milky Way value) forces a large fraction of the CO gas into neutral carbon. Dust is important for H$_2$ and HD formation, already at metallicities of $10^{-4}-10^{-3}$ solar.  Finally, one should always solve the equations of statistical equilibrium to distinguish properly between the excitation, radiation and kinetic temperature of a system. I.e., the thermodynamic floor set by the CMB is only a hard one if the density is high enough (larger than the critical density of a particular transition) to drive collisional de-excitation. \\bigskip\\bigskip\\bigskip"
    },
    "0808/0808.0542_arXiv.txt": {
        "abstract": "A period of slow contraction with equation of state $w > 1$, known as an ekpyrotic phase, has been shown to flatten and smooth  the universe if it begins the phase with small perturbations. In this paper, we explore how robust and powerful the ekpyrotic smoothing mechanism is by beginning with highly inhomogeneous and anisotropic initial conditions and numerically solving for the subsequent evolution of the universe. Our studies,  based on a universe with gravity plus  a scalar field with a negative exponential potential, show  that some regions become homogeneous and isotropic while others exhibit inhomogeneous and anisotropic behavior in which the scalar field behaves like a fluid with $w=1$.  We find that the ekpyrotic smoothing mechanism is robust in the sense that the ratio of the proper volume of the smooth to non-smooth region grows  exponentially  fast along time slices of constant mean curvature. ",
        "introduction": "For over two decades, the only known mechanism for homogenizing, isotropizing and flattening the universe was inflation, a period of accelerated expansion with an equation of state $w$ (ratio of  pressure to  energy density)  near -1. Its success in resolving the horizon and flatness problems is a principal reason why inflation became an essential part of the standard model of cosmology.  In recent years, an alternative mechanism has been discovered in which smoothing and flattening occurs before the big bang as the universe undergoes a period of slow contraction with $w >1$. This alternative, known as the ekpyrotic mechanism \\cite{Khoury:2001wf,Erickson:2003zm}, has been incorporated in alternatives to standard big bang inflationary cosmology including the ``ekpyrotic\" \\cite{Khoury:2001wf},  ``new ekpyrotic\" \\cite{Buchbinder:2007ad} and cyclic models \\cite{Steinhardt:2002ih}.  We note that both the inflationary and ekpyrotic mechanisms can produce nearly scale-invariant spectra for density perturbations in addition to smoothing and flattening the universe.  Until now, the ekpyrotic mechanism has only been shown to work in cases where the deviations from smoothness and flatness are small and perturbative when the ekpyrotic phase begins. The purpose of this paper is to show that the ekpyrotic mechanism is powerful and robust enough to smooth the universe even when the initial perturbations are large and non-linear.  Let us first review how the inflation and the ekpyrotic mechanism work in the perturbative regime where the cosmic evolution is well approximated  by the Friedmann equation with an anisotropy term: \\begin{equation} \\label{eq1} H^2 = \\frac{8 \\pi G}{3} \\left(\\frac{\\rho_m^0}{a^3} + \\frac{\\rho_r^0}{a^4}+ \\frac{\\rho_w^0}{a^{3(1+w)}} \\right) - \\frac{k}{a^2} + \\frac{\\sigma^2}{a^6}, \\end{equation} where $H \\equiv \\dot{a}/{a}$ is the Hubble parameter; $a(t)$ is the Friedmann-Robertson-Walker scale factor normalized so that the  value today, $t_0$, is $a(t_0)=1$; $\\rho_i^0$ represents the present value of the energy density for component $i$, where $m$ represents non-relativistic matter, $r$ represents radiation and $w$ represents an energy component with equation of state $w$, such as a scalar field and its potential. The last two terms on the right-hand side represent spatial curvature and anisotropy.  For an expanding universe, the term that dominates Eq.~(\\ref{eq1}) after a long period of expansion is the one with the smallest power of $a$ in the denominator.  With only radiation and matter, the dominant term would be the spatial curvature, leading to a universe that is unacceptably open or closed by the present epoch.  However, introducing an energy component with $w \\approx -1$ totally changes the outcome  because this component ($\\rho_w$) then has the smallest exponent and dominates the Einstein equation, while the curvature and anisotropy (and other energy components) become negligible. This is the essence of how inflation works if the initial conditions are perturbative.  For inflation, there are several ``Cosmic No-hair'' theorems in addition to numerical results supporting the claim that homogeneity, isotropy and flatness develop even when the initial conditions are non-linear and non-perturbative~\\cite{kolb_turner_86,Goldwirth}.   Now consider the analogous arguments for a contracting universe. With $a(t)$ shrinking, the dominant term in Eq.~(1) will be the one with the largest exponent of $a$ in the denominator.  For a universe with matter and radiation only, this would be the anisotropy term, which famously overtakes the evolution and drives the universe into chaotic mixmaster behavior.  On the other hand, if there is an energy component with $w >+1$, then this energy component dominates instead  of anisotropy or spatial curvature, and chaotic mixmaster behavior never begins \\cite{Erickson:2003zm}.  For a scalar field $\\phi$ with potential energy density $V(\\phi)$, the ratio of pressure to energy density is \\begin{equation}\\label{eq2} w \\equiv \\frac{ \\frac{1}{2} \\dot{\\phi}^2 -V}{\\frac{1}{2} \\dot{\\phi}^2 +V}, \\end{equation} which can be significantly greater than unity when $V$ is less than zero and non-negligible and which approaches $w=1$ if the scalar field kinetic energy dominates. The ekpyrotic phase in ekpyrotic and cyclic models includes an effective scalar field component of this type.  For the cyclic model, the ekpyrotic phase is preceded by a period of dark energy domination and accelerated expansion which, if sustained long enough, would make the universe uniform and flat before the ekpyrotic contraction phase begins.  In this case, the perturbative argument above should be reliable and sufficient to conclude that the universe is smooth and flat as it approaches the big crunch.  However, in the ekpyrotic or new ekpyrotic models, generally, or in the cyclic model with a very short dark energy phase, the conditions at the beginning of the ekpyrotic phase are under less control.  This paper investigates the robustness of the ekpyrotic smoothing and flattening mechanism when the initial conditions are non-linear and non- perturbative to determine how the situation compares with the perturbative case and with inflation. For this study we are not concerned with any particular form for the initial conditions that may be motivated by some specific model; rather, we would like to understand how a generic, highly inhomogeneous and anisotropic space-time evolves under the influence of the proposed smoothing mechanism, modeled here by a scalar field with a negative exponential potential. Such negative exponential potentials arise naturally in supergravity and in string theory. Investigation of the proposed smoothing mechanism requires numerical solution of the coupled Einstein-scalar system of equations.  For simplicity, we restrict our studies to deviations from smoothness along a single spatial dimension.  We use an orthonormal frame representation of the equations written in terms of Hubble-normalized, scale invariant variables~\\cite{Uggla:2003fp} similar to that described in \\cite{Curtis:2005va}, though here coupling to a scalar field instead of a fluid, and using constant-mean-curvature (CMC) time slices. We discretize the equations using second-order accurate finite difference techniques, and solve them with a variant of the Berger and Oliger~\\cite{BO} adaptive mesh refinement (AMR) algorithm for coupled elliptic-hyperbolic equations~\\cite{Pretorius:2005ua}. We find smooth regions that are scalar field dominated in which the scalar field (kinetic plus potential energy density) component behaves like a fluid with $w \\gg 1$, and also regions where the scalar field kinetic energy dominates over the potential energy and the scalar field behaves like a fluid with $w=1$. These latter regions remain inhomogeneous and anisotropic, and throughout this paper we will refer to these parts of the universe as the ``anisotropic regions''. Note however that the anisotropic regions are neither anisotropy nor matter dominated because both the scalar field and the anisotropy of the metric play important roles in the dynamics. Futhermore, note that despite the fact that matter in the smooth regions behaves effectively like a fluid with $w \\gg 1$ there is no issue of superluminal propagation as might arise from an actual fluid with such an equation of state.  This is because we are always solving the scalar wave equation with potential where disturbances always propagate within the light cone. In the anisotropic regions ``spikes'' also form, which are places where the fields change on very small spatial scales, and are similar to regions with this property that have been observerd in numerical simulations of singularities in vacuum spacetimes\\cite{dgspike}. Despite the presence of the scalar field, the anisotropic regions exhibit dynamical behavior similar to chaotic mixmaster vacuum solutions, where there are a series of relatively quick transitions between longer epochs where the solution can be described by a $w=1$ Bianchi type I spacetime. A difference here though is that there are only a {\\em finite} number of transistions, so the mixmaster behavior terminates after several transitions. These dynamics are also known to occur in spacetimes where the matter {\\em is} a fluid with $w=1$~\\cite{Curtis:2005va, Coley05}. AMR is necessary to resolve the spiky features that form both in the anisotropic regions and, in some instances, briefly in what will eventually become smooth scalar field dominated regions, and to resolve the almost domain wall-like transitions that develop between the smooth and anisotropic regions. %  The outline for the rest of the paper is as follows. In Sec. \\ref{sec_methods} we describe the equations, initial conditions, and numerical methods used to solve them. We present the results in Sec. \\ref{sec_results}. The primary conclusion is that a scalar field with a potential inspired by cyclic models is a remarkably powerful and robust smoothing mechanism during a contracting phase of the universe, able to drive the spacetime to homogeneity and isostropy even starting with highly non-linear deviations from an FRW spacetime. Concluding remarks and a discussion of future work is given in Sec. \\ref{sec_conclusion}. ",
        "conclusions": "Our computations provide evidence that the ekpyrotic mechanism for smoothing and flattening the universe is robust and powerful, comparable qualitatively and quantitatively to the inflationary mechanism incorporated in the conventional big bang model. This evidence is the behavior as the singularity is approached of a class of spacetimes that, while not completely general, contain several degrees of freedom and begin far from FRW. Both the inflationary and the ekpyrotic mechanism require the addition of an energy component that is commonly mocked up as a scalar field with potential energy. For inflation, the important feature is that, for some initial conditions, the scalar field can act like a fluid with $w \\approx -1$.  It has been shown numerically that, beginning from highly non-linear and highly irregular initial conditions, regions dominated by the scalar field and with $w \\approx -1$ come to dominate the volume of the universe\\cite{Goldwirth}. Similarly, we have found the regions in which the scalar field acts like a fluid with $w >1$ come to dominate exponentially the volume of the universe during a contracting phase.  This result addresses one of the key criticisms raised when the ekpyrotic model of the universe was first introduced; namely, it was suggested that the model required smooth initial conditions \\cite{Kallosh:2001ai}.   One of the motivations for extending the ekpyrotic picture into a cyclic model was to include a period of dark energy domination before the ekpyrotic phase began in order to prepare smooth conditions  \\cite{Steinhardt:2002ih}.  Now, from the results here, it is clear that the dark energy epoch is not required for this purpose.  Not only does this allow the possibility that the dark energy phase lasts only a few e-folds in the cyclic picture, as suggested in \\cite{Erickson:2006wc}, but it also opens the way for more general bouncing cosmologies that incorporate the ekpyrotic mechanism but do not cycle.  With these results in hand, we are now prepared to tackle the bounce itself in the case that it is non-singular ($a(t)$ shrinks to a non-zero value and then begins to increase). For the non-singular bounce, the equation of state must decrease from $w>1$ to $w< -1$ for a finite period during which anisotropy and inhomogeneity grows.  Our goal is to determine if their growth can be kept at a level consistent with observations, establishing the viability of these bouncing cosmological models."
    },
    "0808/0808.2292_arXiv.txt": {
        "abstract": "{\\psr is one of only 9 known double neutron star systems. These systems are highly valuable for measuring the masses of neutron stars, measuring the effects of gravity, and testing gravitational theories.  } {We determine an improved timing solution for a mildly relativistic double neutron star system, combining data from multiple telescopes. We set better constraints on relativistic parameters and the separate masses of the system, and discuss the evolution of \\psr in the context of other double neutron star systems.} {\\psr has been regularly observed for more than 10 years by the European Pulsar Timing Array (EPTA) network using the Westerbork, Jodrell Bank, Effelsberg and Nan\\c cay radio telescopes. The data were analysed using the updated timing software \\tempo.} {We have improved the timing solution for this double neutron star system. The periastron advance has been refined and a significant detection of proper motion is presented. It is not likely that more post-Keplerian parameters, with which the individual neutron star masses and the inclination angle of the system can be determined separately, can be measured in the near future. } {Using a combination of the high-quality data sets present in the EPTA collaboration, extended with the original GBT data, we have constrained the masses in the system to $m_\\mathrm{p}<1.17$~\\msun and $m_\\mathrm{c}>1.55$~\\msun ($95.4\\% $ confidence), and the inclination angle of the orbit to be less than $47$ degrees (99\\%). From this we derive that the pulsar in this system possibly has one of the lowest neutron star masses measured to date. From evolutionary considerations it seems likely that the companion star, despite its high mass, was formed in an electron-capture supernova. } ",
        "introduction": "\\psr was discovered in the Green Bank Northern Sky Survey \\citep{snt97}. It is one of only 9 double neutron star systems (DNSs) known.  The pulsar orbits its companion neutron star in 8.63~days, and its 40~ms spin period, together with the derived low surface magnetic field, is typical of a mildly recycled pulsar.  \\cite{nst96} first described the system and have already shown that the space velocity is probably quite low. They also measured the periastron advance from which the total system mass was estimated.  \\cite{tc99}, using a Bayesian analysis with no constraints on the orientation of the orbit, found the masses of the pulsar and its companion to be $m_\\mathrm{p}= 1.56^{+0.20}_{-1.20}$ \\msun and $m_\\mathrm{c}= 1.05^{+1.21}_{-0.14}$ \\msun ($95\\% $\\,confidence).  An improved timing solution was presented by \\cite{hlk+04}, although not discussed in detail.  DNSs are presently the best available tool for testing strong-field gravity effects.  The \\psr\\ orbit is only mildly relativistic, making it of limited use for tests of gravitational theories.  However, any constraints on its post-Keplerian (PK) parameters, the inclination of its orbit, or the masses of the pulsar and its companion are highly valuable for studies of the evolution of these systems (e.g. \\citealt{plp+04,kbk+07}).  Like PSR~J1811$-$1736 \\citep{cks+07}, \\psr\\ has a rather wide orbit. This could indicate that the evolution into a DNS system has been slightly different than most other systems, which generally have tight orbits with periods of several hours. However, \\psr\\ does follow the recently discovered spin period-eccentricity relation \\citep{fkl+05,dpp05}, suggesting that the evolution cannot be too different from the tight DNSs.  \\psr has been observed regularly using the four 100~m class radio telescopes in Europe. This has allowed us to select high-quality data resulting in the best possible analysis of the system parameters to date. For completeness, we have also included data from telescopes at Green Bank as presented in \\cite{nst96, nts99}.  We present the new timing solution and discuss the limits our solution puts on the inclination and the masses of the neutron stars in the system.  We discuss the future prospects of detecting multiple PK parameters, and we related our results to the evolution of DNSs. ",
        "conclusions": "\\subsection{Masses}\\label{section:masses}  \\begin{figure} \\centering \\includegraphics[width=6.4cm, angle=270]{square.pdf} \\caption{Mass-mass diagram for \\psr and its companion. The hatched region is excluded for $i > 90$ degrees, the dash-dotted lines are indicating constraints resulting from the \\xdot\\ measurement, $i < 69$ degrees and shapiro delay limit, $i < 47$ degrees. The diagonal line constrains $\\dot\\omega$, with errors. The dotted lines for constant pulsar and companion mass indicate the range of neutron star masses measured in other DNSs \\citep{tc99,lbk+04,fkl+05}. \\label{fig:massplot}} \\end{figure}  For any given theory of gravity, measuring PK parameters will put constraints on the masses and the inclination of the system, depending only on the individual masses and the Keplerian parameters (e.g. \\citealt{sta03}). Figure~\\ref{fig:massplot} shows all the current restrictions we were able to derive from our timing solution (Table~\\ref{tab:jwnge-2}). The improved, highly significant \\omdot\\ measurement gives us a very accurate determination of the total mass. From the non-detection of Shapiro delay parameters it is clear that the inclination of the orbit is quite low.  Furthermore, the $\\dot{x}$ measurement presented in \\S 3.3, taken at face value, confirms that the system is at low inclination angles.  The relativistic time-dilation/gravitational redshift parameter, $\\gamma$, is only measurable when the periastron advance is large enough to decouple its effect on the timing residuals from the measurement of the semimajor axis and its time derivative \\citep{bt75, mt77, dt92}. Changing $\\omega$ by only a few degrees will take several hundred years, and a detection of $\\gamma$ can therefore not be expected in the near future. Moreover, \\xdot\\ turns out to be highly covariant with $\\gamma$ in the timing analysis.  This makes an independent measurement of $\\gamma$ impossible and obscures the interpretation of the apparent measurement of \\xdot.  To clarify the influence of all the PK perturbations on our timing data, and to calculate accurate, refined values of the stellar masses, we performed a comprehensive, self-consistent timing analysis which simultaneously incorporated all relativistic and kinematic phenomena in the timing solution.  Our primary goal was to measure or constrain the values of $m_\\mathrm{p}$ and $m_\\mathrm{c}$.  This procedure also finds constraints on the two angles which describe the orientation of the orbit, $i$ and $\\Omega$.  For convenience, we always represent the latter as an offset from the proper motion position angle, i.e. $\\Theta_\\mathrm{\\mu}-\\Omega$.  This angle runs between $-180^\\circ$ and 180$^\\circ$, while inclination runs between $0^\\circ$ and $180^\\circ$.  The four degrees of freedom in this problem ($m_\\mathrm{p}$, $m_\\mathrm{c}$, $i$, and $\\Theta_\\mathrm{\\mu}-\\Omega$) are reduced to three by noting that the two masses and the inclination are related via the Keplerian mass function equation: \\begin{equation}\\label{eqn:kepler} f\\equiv\\frac{(m_\\mathrm{c}\\sin i)^3}{(M_\\mathrm{T})^2} = \\frac{x^3}{T_\\mathrm{\\odot}}\\left(\\frac{2\\pi}{P_\\mathrm{b}}\\right)^2, \\end{equation} where all quantities on the right side of the equation are measured to high precision.  We analysed a grid of timing solutions in a three-dimensional parameter space following the approach as explained in detail in \\cite{sna+02, sns+05}. For the three variables, we used $M_\\mathrm{T}$, $\\cos i$, and $\\Theta_\\mathrm{\\mu}-\\Omega$, and we assumed a uniform prior in each of these.  For randomly oriented binary systems, $\\cos i$ and $\\Theta_\\mathrm{\\mu}-\\Omega$ each follow uniform distributions, and since \\omdot\\ is well known, only a small range of $M_\\mathrm{T}$ values need be considered, so a uniform prior is acceptable for this variable as well. For each point in our three dimensional grid, we used Eq.~\\ref{eqn:kepler} to calculate the value of $m_\\mathrm{c}$ and the corresponding $m_\\mathrm{p}$. We then used the masses and orientation angles to calculate the kinematic perturbation parameters (Eqs.~\\ref{eq:xdot},~\\ref{eq:omdot}) and relativistic timing parameters according to GR (e.g. \\S\\,4.1 of \\cite{sta03}).  The \\omdot\\ used in the analysis was the sum of the kinematic and relativistic terms held fixed while all other pulsar timing parameters (astrometric, rotational, and Keplerian orbital parameters) were free to vary.  We recorded the $\\chi^2$ of each timing solution, assigning a probability to each grid point based on the difference between its $\\chi^2$ and the global minimum. We used this ensemble of probabilities to calculate confidence regions for the parameters of interest and to place limits on the stellar masses.  The results are given in Figs. \\ref{fig:abcd}, \\ref{fig:chi7} and \\ref{fig:paramplot}.  Figure \\ref{fig:abcd} shows the orientations of the orbit allowed by the timing solution.  The figure is a projection of the 95.4\\% confidence volume of the three dimensional analysis grid onto the two dimensional space shown.  There are four regions allowed by the timing data.  The Keplerian orbital parameters, combined with the lack of Shapiro delay restrict $i$, but because the Shapiro delay depends on $\\sin~i$, the resulting constraint is degenerate: if $i$\\, is allowed, so is $180^\\circ-i$.  Since there is a detectable \\xdot, Eq. \\ref{eq:xdot} restricts $\\Theta_\\mathrm{\\mu}-\\Omega$ to two possibilities (one positive and one negative value) for any given value of $i$.  Thus there are four regions of allowed solutions, labeled $A$, $B$, $C$ and $D$ in the figure.  Figure \\ref{fig:chi7} shows 68.3\\%, 95.4\\%, and 99.7\\% confidence limits on $\\cos~i$ and $M_\\mathrm{T}$ space.  For reference, the figure shows lines of constant mass difference $m_\\mathrm{c}-m_\\mathrm{p}$, at intervals of 0.2\\msun, as calculated from $i$ and $M_\\mathrm{T}$, and it indicates the region at low inclination angle that is excluded because $m_\\mathrm{p}>0$ is not satisfied in that region. The confidence limits were calculated by marginalizing over the probability values for values of $\\Theta_\\mathrm{\\mu}-\\Omega$ for a given combination of $\\cos~i$ and $M_\\mathrm{T}$.  The best-fit solutions have relatively high $m_\\mathrm{c}-m_\\mathrm{p}$, i.e. high companion masses and low pulsar masses, although the confidence contours stretch to much lower values of $m_\\mathrm{c}-m_\\mathrm{p}$.  The determination of $M_\\mathrm{T}$ is dominated by the relativistic \\omdot, but the obtained values are perturbed by the kinematic $\\delta$\\omdot, given by Eq. \\ref{eq:omdot}. This perturbation splits the allowed values for the total mass into two regions.  This can be understood by plugging the values of $\\Theta_\\mathrm{\\mu}-\\Omega$\\, from each of the four regions of Fig.~\\ref{fig:abcd} into Eq.~2, and noting that $\\csc~i$ is always positive, so that $\\delta$\\omdot\\ is positive for solutions $A$ and $B$ and negative for solutions $C$ and $D$.  This means that the observed $\\delta$\\omdot\\ has been biased towards higher values for solutions $A$ and $B$, so that the true relativistic $\\delta$\\omdot\\ is lower than that calculated without the kinematic correction.  Since $M_\\mathrm{T}$ is proportional to $\\delta$\\omdot$^{3/2}$, this means that the true total mass is lower for solutions $A$ and $B$. Similarly, the true mass is higher for solutions $C$ and $D$.  Figure \\ref{fig:paramplot} shows the 95.4\\% confidence volume projected into several two-dimensional parameter spaces.  For this figure, the values of $i$ and $\\Theta_\\mathrm{\\mu}-\\Omega$ correspond to solution $D$, but the distributions of all other quantities are essentially identical for all four solution regions. We note that, for completeness, all possible values for the parameters are shown in Figs. \\ref{fig:abcd}, \\ref{fig:chi7} and \\ref{fig:paramplot} however we consider it unlikely that the pulsar mass will be lower than $1$\\msun.  We used the probabilities from the grid analysis to constrain $m_\\mathrm{p}$ and $m_\\mathrm{c}$.  The results are essentially identical for all four solution regions. The central 95.4\\% confidence intervals are $m_\\mathrm{p}=0.72^{+0.51}_{-0.58}$\\,\\msun and $m_\\mathrm{c}=2.00^{+0.58}_{-0.51}$\\,\\msun. If, instead of central confidence intervals, we use 95.4\\% confidence upper and lower limits for the individual masses, these convert to $m_\\mathrm{p}<1.17$\\,\\msun and $m_\\mathrm{c}>1.55$\\,\\msun. Note that the apparent discrepancy between these numbers originates from different areas covered in the probability distribution. For the total mass, the 95.4\\% confidence intervals are $2.7188\\pm0.0011$\\,\\msun for solutions $A$ and $B$, and $2.7217\\pm0.0018$\\,\\msun for solutions $C$ and $D$.  However, allowing any of the four solutions yields the range $2.720_{-0.002}^{+0.003}$\\msun. This last number is the most accurate value for $M_\\mathrm{T}$, as it reflects the uncertainty as to which orientation of the orbit is correct.      \\begin{figure} \\centering \\includegraphics[width=7.0cm, angle=270]{fig6new.pdf} \\caption{Allowed values for orbital orientation $i$ and $\\theta_\\mathrm{\\mu}-\\Omega$.  The figure shows a projection of the 95.4\\% confidence volume of the three dimensional grid analysis of timing solutions onto this two dimensional space.  See the text for discussion. \\label{fig:abcd}} \\end{figure}    \\begin{figure} \\centering \\includegraphics[width=7.0cm, angle=270]{fig7new.pdf} \\caption{Allowed values for total mass and $\\cos~i$.  The two sets of contours are 68.3\\%, 95.4\\%, and 99.7\\% confidence limits on these quantities for solutions $A$ and $B$ (lower, white contours) and $C$ and $D$ (upper, gray contours).  Dotted lines indicate values of $m_\\mathrm{c}-m_\\mathrm{p}$, at intervals of 0.2\\msun. The dark region on the right of the plot is excluded as it does not satisfy $m_\\mathrm{p}>0$. \\label{fig:chi7}} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=10.0cm]{fig8new.pdf} \\caption{Values of relativistic $\\gamma$, kinematic $\\delta\\dot{\\omega}$, kinematic $\\dot{x}$, orbital orientation $\\theta_\\mathrm{\\mu}-\\Omega$, and masses $m_\\mathrm{p}$ and $m_\\mathrm{c}$ allowed by the timing solution, as a function of inclination angle.  The figure shows a projection of the 95.4\\% confidence volume of the three dimensional grid analysis of timing solutions onto each of the two dimensional spaces shown.  The regions corresponds to solution $D$, but other solutions give essentially identical results. \\label{fig:paramplot}} \\end{figure}   \\subsection{Evolution}  The system is valuable for studies of DNS evolutionary scenarios.  It has been believed for a long time that supernova explosions result in high kick velocities on the remaining neutron star \\citep{bai89}. Recently, it has been argued that different evolutionary scenarios can be possible in binary systems with particular mass and chemical properties \\citep{plp+04,heu07}. The so-called electron-capture collapse of an O-Ne-Mg core of a helium star leads to a {\\it fast} supernova explosion which is believed to be symmetric and therefore does not result in a kick of the second-born neutron star.  For most DNSs in which proper motion measurements have been possible, a low space velocity has been derived. Our measurement of the proper motion of \\psr, implying a velocity of about 25 km\\,s$^{-1}$, is consistent with the evolutionary scenario mentioned above, and another piece of evidence that not all pulsars receive a large kick at birth.  Van den Heuvel (2007) argues that the low velocity of DNS systems may be correlated with the masses of the second-formed neutron stars in DNSs being somewhat lower than normal: $1.25(6)$\\msun.  Our 95.4\\% limit\\footnote{The 99.7\\% limit is $m_\\mathrm{c} > 1.39$\\msun} of $m_\\mathrm{c} > 1.55$\\msun appears to contradict this correlation.  In case of a symmetric supernova the mass lost during the explosion can be calculated by the expression $e=\\Delta M_\\mathrm{SN}/M_\\mathrm{T}$ \\citep{bv91}, where $e$ is the eccentricity and $M_\\mathrm{T}$ the total mass after the supernova explosion. For \\psr\\ this results in $\\Delta M_\\mathrm{SN}=0.68$\\msun.  If we use our 95.4\\% lower limit on the companion of 1.55~\\msun this leads to a core-progenitor mass of 2.23\\msun, which is considered to be in the lower range for helium cores \\citep{dp03b}, and the progenitor star must have had a mass of about $9$\\msun. As these stars are believed to produce degenerate O-Ne-Mg cores \\citep{plp+04,phlh08}, this is another indication that the companion neutron star must have formed in an electron-capture core collapse. The somewhat higher companion mass can possibly be explained by assuming extra fall-back during the supernova event of $\\sim0.2$\\msun (e.g. \\citealt{fk01,fry07}). To allow for this amount of fallback the supernova explosion has probably been very weak. An electron-capture core collapse is the most probable explanation for the parameters of this system. However, constraining the masses even better would be very interesting.    \\subsection{Search for the companion}  Apart from the original discovery papers \\cite{nst96,snt97}, there is no report of searches for pulsations of the companion of \\psr. The Green Bank Northern Sky Survey, optimized to find millisecond pulsars, had a flux-density limit at 370 MHz of 8 mJy for slow pulsars (with periods above 20 ms).  The very high characteristic age and the low magnetic field imply that \\psr is the first-born and recycled neutron star in the system. The (unseen) second neutron star will be the young object, with a long spin period, which has probably already slowed down significantly and possibly passed the death-line to become undetectable as a radio pulsar. Moreover, if the second neutron star is still active as a radio pulsar, it has a reasonable chance of being beamed away from us. In this case, taking into account that the precession timescale of the spin axis is very long, it is not expected that the companion, if active as a pulsar, will rotate into view on a short timescale.  However to be certain that no weak pulsed signal from the second neutron star has been missed, we have searched several WSRT observations of 30 minutes at 840, 1380 and 2300 MHz for pulsations of the companion of \\psr. We performed an acceleration search on the data although, because of the wide orbit, the expected smearing due to the orbital motion of the companion is likely to be negligible.  In the only double pulsar system known so far, PSR~J0737$-$3039A/B, the second pulsar is only visible for a small part of each orbit \\citep{lbk+04}.  To take this possibility into account, and assuming the second pulsar would be visible in an equal fraction of the orbit as PSR~J0737$-$3039B, we have searched another set of observations spread equally across the full orbital phase range. In all searches no evidence of pulsations of a companion were found to the limits of 1.1, 0.25 and 0.5 mJy at 840, 1380 and 2300 MHz respectively, which means that if the companion is active as a pulsar, and beamed towards us, it must be less luminous than 0.098 mJy~kpc$^2$ at 1380 MHz."
    },
    "0808/0808.3227_arXiv.txt": {
        "abstract": "Previous solar observations have shown that coronal loops near 1\\,MK are difficult to reconcile with simple heating models. These loops have lifetimes that are long relative to a radiative cooling time, suggesting quasi-steady heating. The electron densities in these loops, however, are too high to be consistent with thermodynamic equilibrium. Models proposed to explain these properties generally rely on the existence of smaller scale filaments within the loop that are in various stages of heating and cooling. Such a framework implies that there should be a distribution of temperatures within a coronal loop.  In this paper we analyze new observations from the EUV Imaging Spectrometer (EIS) on \\textit{Hinode}. EIS is capable of observing active regions over a wide range of temperatures (\\ion{Fe}{8}--\\ion{Fe}{17}) at relatively high spatial resolution (1\\arcsec).  We find that most isolated coronal loops that are bright in \\ion{Fe}{12} generally have very narrow temperature distributions ($\\sigma_T \\lesssim 3\\times10^5$\\,K), but are not isothermal. We also derive volumetric filling factors in these loops of approximately 10\\%. Both results lend support to the filament models. ",
        "introduction": "High spatial resolution solar observations have shown that coronal loops with temperatures near 1\\,MK have properties that are difficult to reconcile with physical models. Loops at these temperatures persist for much longer than a radiative cooling time, suggesting quasi-steady heating. The densities inferred from the observations, however, are much higher than can be reproduced by steady, uniform heating models. The temperature gradients along the loops are also much smaller than predicted by the simple models.  (e.g., \\citealt{lenz1999,aschwanden2000b,winebarger2003}).  Several models have been proposed to explain the properties of coronal loops at these temperatures. \\cite{aschwanden2000b}, for example, suggested that the observed loops were actually composed of smaller scale threads that were steadily heated at their footpoints. Footpoint heating leads to somewhat higher densities and flatter temperature gradients relative to steady heating models. At high densities, however, loops at these temperatures can become thermodynamically unstable (e.g., \\citealt{mok2005,muller2004,winebarger2003}), leading to catastrophic cooling. Multi-thread, impulsive heating models have also been suggested (e.g., \\citealt{warren2003}). In these models the fact that loops cool much more rapidly than they drain accounts for the high densities. Multiple threads in various stages of heating and cooling are needed to explain the observed lifetimes and temperature gradients. In both cases, these multi-thread models indicate the need for emission formed over a range of temperatures to reproduce the observed intensities (see also \\citealt{reale2000}).  One limitation of the observational results from \\textit{TRACE} is that they are derived from narrowband filtergrams with somewhat limited diagnostic capabilities.  The launch of the EUV Imaging Spectrometer (EIS) on the \\textit{Hinode} mission provides us with an opportunity to revisit some of these observational results using spectroscopic data. EIS is a high spatial and spectral resolution spectrometer that covers much of the same wavelength range as \\textit{TRACE}. EIS has a very broad temperature coverage and can image the solar corona in individual emission lines from the lower transition region to the hottest flares.  In this paper we focus on measuring the emission measure distribution in coronal loops near 1\\,MK. We have selected 20 relatively isolated loop segments from several different active region observations and computed differential emission measure distributions from the background subtracted loop intensities. For this work we focus on loops that are bright in \\ion{Fe}{12} and find that for these loops the distribution of temperatures is almost always narrow, with a dispersion of several times $10^5$\\,K.  We also find volumetric filling factors of approximately 10\\%. These results support the idea that coronal loops are composed of smaller scale filaments that are below the spatial resolution of current solar instruments. ",
        "conclusions": "Some previous work has suggested that coronal loops, as currently observed, are isothermal.  For example, \\cite{aschwanden2005b} find that the majority of narrowest loops observed with TRACE are consistent with an isothermal DEM. Since TRACE is limited to observations in only three channels (\\ion{Fe}{9}, \\ion{Fe}{12}, and \\ion{Fe}{15}) it is difficult to distinguish between an isothermal distribution and the narrow distributions that we measure spectroscopically. The general absence of \\ion{Fe}{15} emission in the loops that we have studied is consistent with \\cite{aschwanden2005b}. \\cite{delzanna2003} also found examples of relatively cool ($\\sim0.9$\\,MK) nearly isothermal loops observed with low resolution spectroscopic data. These results also suggested filling factors near 1. Our filling factor results are smaller than this, but we also find that the filling factor to be inversely proportional to the loop pressure (also see \\citealt{warren2008}). We do see some loops with a relatively broad emission measure distribution ($\\log\\sigma_T\\sim5.7$), which is consistent with the results of \\cite{schmelz2007} and \\cite{patsourakos2007}. Our sample, which is small, suggests that such loops are rare, however.  These new observational results lend support to the non-equilibrium, multi-thread models of these ``warm'' coronal loops. It remains to be seen if hydrodynamic models can reproduce the observed loop properties. The combination of high densities and narrow temperature ranges will be difficult to reconcile with nanoflare models (e.g., \\citealt{patsourakos2006}). The narrow temperature distributions suggest that these filaments are evolving coherently."
    },
    "0808/0808.3820_arXiv.txt": {
        "abstract": "Precession in an accretion-powered pulsar is expected to produce characteristic variations in the pulse properties. Assuming surface intensity maps with one and two hotspots, we compute theoretically the periodic modulation of the mean flux, pulse-phase residuals and fractional amplitudes of the first and second harmonic of the pulse profiles. These quantities are characterised in terms of their relative precession phase offsets. We then search for these signatures in 37 days of X-ray timing data from the accreting millisecond pulsar XTE J1814$-$338. We analyse a 12.2-d modulation observed previously and show that it is consistent with a freely precessing neutron star only if the inclination angle is $< 0.1^\\circ$, an a priori unlikely orientation. We conclude that if the observed flux variations are due to precession, our model incompletely describes the relative precession phase offsets (e.g. the surface intensity map is over-simplified). We are still able to place an upper limit on $\\epsilon$ of $3.0 \\times 10^{-9}$ independently of our model, and estimate the phase-independent tilt angle $\\theta$ to lie roughly between $5^\\circ$ and $10^\\circ$. On the other hand, if the observed flux variations are not due to precession, the detected signal serves as a firm upper limit for any underlying precession signal. We then place an upper limit on the product $\\epsilon \\cos \\theta$ of $\\leq 9.9 \\times 10^{-10}$. The first scenario translates into a maximum gravitational wave strain of $10^{-27}$ from XTE J1814$-$338 (assuming a distance of 8 kpc), and a corresponding signal-to-noise ratio of $\\leq 10^{-3}$ (for a 120 day integration time) for the advanced LIGO ground-based gravitational wave detector. ",
        "introduction": "\\label{intro} Accreting millisecond pulsars (AMSPs) are a subset of neutron stars in low-mass X-ray binaries (LMXBs) that exhibit persistent X-ray pulsations with periods below 10 ms. In the standard recycling scenario, AMSPs are the evolutionary link between LMXBs and nonaccreting, radio millisecond pulsars \\citep{alpar82, radha82}. Eight AMSPs have been discovered at the time of writing \\citep{wij04, morgan05, galloway07, krimm07}.  Most AMSPs are X-ray transients. Once every few years, they emerge from quiescence and become detectable during an outburst lasting several weeks. The outburst is attributed to enhanced accretion [e.g. \\citet{lasota01}], funnelled onto a small number of hotspots on the star. Little is known about the shape, position, or number of these hotspots \\citep{roma04, kulk05}, but they do give rise to detectable X-ray pulsations, from which the spin period and orbital parameters can be determined. During an outburst, surface thermonuclear burning also causes type I X-ray bursts, which last a few minutes and occur on average once every few days. Type I X-ray bursts have been observed in three AMSPs to date: SAX J1808.4$-$3658, XTE J1814$-$338 \\citep{wijnands05}, and HETE J1900.1$-$2455 \\citep{vanderspek05}. In the burst tails, a small component of the X-ray flux ($\\sim 15$\\% for XTE J1814$-$338) oscillates at the spin frequency.  AMSPs are expected to be relatively powerful gravitational wave sources \\citep{wattskrishnan08}. The fastest, IGR J00291+5934 \\citep{eckert04}, spins at $\\Omega_\\ast/2\\pi = 599$ Hz, well below the theoretical breakup frequency for most nuclear equations of state ($\\sim$ 1.5 kHz) \\citep{cook94, bildsten98}. Similarly, the fastest radio millisecond pulsar, PSR J1748$-$2446ad \\citep{hessels06}, and the fastest nonpulsating LMXB, 4U 1608$-$52 \\citep{hartman03}, spin at frequencies of 716 Hz and 619 Hz respectively. The gap below the breakup frequency is explained if the star is deformed by one part in $\\sim 10^{8}$, such that gravitational radiation balances the accretion torque at hectohertz frequencies \\citep{bildsten98}. Several physical mechanisms can produce the requisite deformation: magnetic mountains \\citep{pamel04, melpa05, pamel06, vigmel08}, thermocompositional mountains caused by electron capture gradients \\citep{usho00}, toroidal internal magnetic fields \\citep{cutler02}, and r-modes \\citep{anders98, owen98, nayyar06}. AMSPs are therefore promising targets for ground-based, long-baseline interferometers like the Laser Interferometer Gravitational-Wave Observatory (LIGO). An AMSP at a distance of 1 kpc, spinning at 0.4 kHz with ellipticity $\\epsilon = 10^{-8}$, generates a wave strain $h \\sim 10^{-27}$. By comparison, initial LIGO's sensitivity threshold in the 0.1--0.4 kHz band is $\\sim 10^{-26}$ during the S4 run \\citep{abbott07}. Advanced LIGO will get down to $h \\sim 10^{-27}$ in the same band, and narrowband tunability will increase its sensitivity to AMSPs further, as $\\Omega_\\ast$ is known a priori from X-ray timing.  An AMSP with ellipticity $\\epsilon \\sim 10^{-8}$ is expected to precess with a period of hours to days.  Magnetic mountains, for example, are built around the magnetic axis, which is misaligned in general with the rotation axis in objects which pulsate \\citep{pamel06}. More generally, a mass quadrupole of any provenance should be kicked out of alignment continuously by stochastic accretion torques \\citep{joneand02}. Hence AMSPs are promising observational candidates for observing short-period precession. Until now, however, precession has been difficult to detect in neutron stars. Only one source, the radio pulsar PSR B1828$-$11, precesses unambiguously, with period $P_{\\rm{p}} =$ 250 d \\citep{stalyshem00}. Oscillatory trends in pulse arrival times, with periods of several days, have also been reported tentatively in a few other objects \\citep{melatos00, hobbs06, pamel06}, but the physical cause is unclear.  Free precession consists of a fast wobble about the angular momentum vector $\\mathbf{J}$, at approximately the pulsar spin period $P_\\ast = 2\\pi/\\Omega_{\\ast}$ and a slow retrograde rotation about the symmetry axis, with period $P_{\\rm{p}} = 2 \\pi/\\Omega_{\\rm{p}}$, which modulates the pulse shape and arrival times \\citep{zimsz79, alpines85, joneand01, joneand02, link03}. The precession frequency $\\Omega_{\\rm{p}}$ depends on the ellipticity, $\\epsilon$, and the tilt angle $\\theta$ (between the symmetry axis and \\textbf{J}), with \\begin{equation} \\label{gw} \\epsilon \\cos \\theta \\approx \\Omega_p/\\Omega_{\\star}. \\end{equation}  The amplitude ratio of the gravitational wave signal at the spin frequency and its second harmonic \\citep{zimsz79, jara98}, and in the + and $\\times$ polarizations, provides \\textit{independent} information on $\\epsilon, \\theta$, the orientation of \\textbf{J}, and the emission pattern on the surface of the star. Narrowband tunability facilitates extraction of this information.  In this paper, we compute theoretically the X-ray signal from a precessing pulsar for a range of orientations and compare three quantities from each pulse profile to the data: the mean flux of the profile, the zero-to-peak pulse amplitude, and the pulse-phase residuals.  We search for the signature of precession in X-ray timing data from one particular AMSP, XTE J1814$-$338. An analogous search was carried out by \\citet{akgun06} for the radio pulsar PSR B1828-11, who modelled the period residuals and pulse shapes analytically taking into account precession effects (biaxial and triaxial) as well as the contribution from the magnetic spin-down torque. The authors performed searches over a range of beam locations, degrees of triaxiality, tilt angles and angle-dependent spin-down torques, finding a wide range of parameters which match the data. Thus, they were unable to constrain the shape of the star but did find that the angle-dependent spin-down torque contributes to the period residuals. Their method differs from ours in that, instead of fitting the shape of the residuals and comparing for each set of parameters, they determined the validity of a configuration by calculating Bayesian probability distribution functions for the parameters under certain constraints.  The paper is structured as follows. Section \\ref{model} describes the precession model and its implementation. Sections \\ref{single} and \\ref{double} characterize the predicted X-ray signal for a biaxial, precessing pulsar with one and two hotspots respectively, specifically the relative precession phases between the flux, pulse amplitude, and pulse-phase. Section \\ref{triaxial} repeats the predictions for a triaxial, precessing pulsar. Section \\ref{datareduction} describes the data reduction and timing analysis of XTE J1814$-$338. We compare the measurements with the theory in Section \\ref{compare} and derive upper limits on $\\epsilon$, $\\theta$, and the associated gravitational wave strain in Section \\ref{conclusion}. The limit on $\\theta$ constrains the relative strength of the driving and damping forces in the system. ",
        "conclusions": "either the star is precessing but our surface intensity map is too simplistic, or the source is not precessing. If we attribute the 12.2-d periodicity to precession, this implies an ellipticity of $\\epsilon \\leq 3 \\times 10^{-9}$, a gravitational wave strain $h_0 \\leq 10^{-27}$, and hence a signal-to-noise ratio of $10^{-3}$ for initial LIGO and $10^{-2}$ for advanced LIGO (for a coherent 120-day search). On the other hand, if the precession is damped by internal dissipation ($\\theta$ is small), or the precession period is much longer than the 37-day data span ($\\epsilon$ is small), some other mechanism must cause the observed modulation. In this scenario, we find $\\epsilon \\cos \\theta \\leq 9.9 \\times 10^{-10}$ and $h_0 \\leq 10 \\times 10^{-27} \\cos(\\theta)^{-1}$.  Although we face a negative result for this particular source, this paper establishes a framework for analyzing modulations in X-ray flux from AMSPs for a range of geometrical configurations and surface intensity maps. We anticipate that the framework will be applied to other AMSPs in the future. Given the values of $\\epsilon$ inferred from the gravitational-wave stalling hypothesis \\citep{bildsten98} and the theoretical models, e.g. of magnetic mountains \\citep{pamel06, vigmel08}, it is clear that long-term X-ray monitoring of AMSPs (over years) is essential for predicting, and then searching for, their gravitational wave signal."
    },
    "0808/0808.0018.txt": {
        "abstract": "We have obtained near-infrared spectra covering the \\ion{Ca}{2} triplet lines for a large number of stars associated with 16 SMC clusters using the VLT + FORS2. These data compose the largest available sample of SMC clusters with spectroscopically derived abundances and velocities.  Our clusters span a wide range of ages  and provide good areal coverage of the galaxy.  Cluster members are selected using a combination of their positions relative to the cluster center as well as their location in the CMD, abundances and radial velocities.  We determine mean cluster velocities to typically 2.7 km s$^{-1}$ and metallicities to 0.05 dex (random errors), from an average of 6.4 members per cluster.  By combining our clusters with previously published results, we compile a sample of 25 clusters on a homogenous metallicity scale and with relatively small metalliciy errors, and thereby investigate the metallicity distribution, metallicity gradient and age-metallicity relation (AMR) of the SMC cluster system.  For all 25 clusters in our expanded sample, the mean metallicity [Fe/H] = $-$0.96 with $\\sigma$ = 0.19. The metallicity distribution may possibly be bimodal, with peaks at $\\sim -$0.9 dex and $-$1.15 dex.  Similar to the LMC, the SMC cluster system gives no indication of a radial metallicity gradient. However, intermediate-age SMC clusters are both significantly more metal-poor and have a larger metallicity spread than their LMC counterparts. Our AMR shows evidence for 3 phases: a very early ($>11$ Gyr) phase in which the metallicity reached $\\sim -$1.2 dex, a long intermediate phase from $\\sim 10-3$ Gyr in which the metallicity only slightly increased, and a final phase from 3$-$1 Gyr ago in which the rate of enrichment was substantially faster.  We find good overall agreement with the model of \\citet{pag98}, which assumes a burst of star formation at 4 Gyr. Finally, we find that the mean radial velocity of the cluster system is 148 km s$^{-1}$, with a velocity dispersion of 23.6 km s$^{-1}$ and no obvious signs of rotation amongst the clusters.  Our result is similar to what has been found from a wide variety of kinematic tracers in the SMC, and shows that the SMC is best represented as a pressure supported system. ",
        "introduction": "The Small Magellanic Cloud (SMC) has long been recognized as being of fundamental importance for a wide variety of astrophysical studies. First, the current paradigm of galaxy formation suggests that spiral galaxy spheroids, such as the Milky Way (MW) halo, are formed by the accretion/merger of smaller, satellite galaxies (e.g., \\citealt{sea78,zen03}). As one of the nearest low mass galaxies known, the SMC is thus a possible prototype of pre-Galactic fragments. However, current $\\Lambda$CDM models suggest that the majority of the MW\u00b4s building blocks were assimilated very early in its history and that existing low mass galaxies like the SMC are survivors of this process and thus underwent a different chemical evolution (e.g., \\citealt{rob05}). At the least, many dynamical simulations (e.g., \\citealt{bek04}) suggest that the MW, Large Magellanic Cloud (LMC) and SMC compose a longterm interacting system, with the SMC and LMC likely to be eventually consumed by the MW.  In addition, close encounters in the last few Gyr may well have stimulated star formation on a global scale in both the SMC and LMC (\\citealt{mur80,gar94,yos03,bek04}). These early simulations predict bursts of star formation $\\sim 0.2$Gyr ago that led to the formation of the eastern wing of the SMC and the Magellanic Bridge, and  a similar close encounter event some 4 Gyr ago. Note however that the accurate proper motions crucial for properly modeling the orbits of the SMC and LMC are still problematic, and that the most recent values all suggest the Magellanic Cloud may be unbound to each other and only beginning their first close encounter with the MW (\\citealt{kal06,pia08}).  The SMC, with its low global metallicity, is the best local counterpart to the host of distant dwarf irregular and blue compact dwarf galaxies, making it an attractive target for exploring the importance of metallicity in a number of contexts, including star formation, initial mass function, stellar and galaxy evolution, etc.  In particular, the star clusters of the SMC, because they are (at least to first order) simple stellar populations (SSPs), are an invaluable resource with which we can explore the structure, kinematics, star formation and chemical evolution history of the SMC.  As a majority of stars may have formed in clusters (see e.g., \\citealt{lad03} for the solar neighborhood and \\citealt{cha06} for the SMC, but see \\citealt{bas09} for an opposing view), their study has become even more relevant.  On a cosmological scale, star clusters in the SMC are of utmost importance for the study and understanding of stellar populations in distant galaxies, since they cover age and abundance space that is not occupied by their Galactic counterparts. In addition, while the LMC cluster system suffers from the well known age gap, where only one cluster is known to have formed between $\\sim$3 and $\\sim$13 Gyr ago, (i.e.~during 3/4 of its life; \\citealt{dac91}, \\citealt{gei97}), the SMC posesses clusters that have been forming more or less continuously over the past $\\sim$11 Gyr (e.g., \\citealt{gl08b}).  The SMC is the only dwarf galaxy in the Local Group that has formed and preserved populous star clusters, without a significant age gap, across (most of) the age of the Universe, and the production has been prolific; \\citet{hod86} estimates the SMC has some 2000 clusters.  Despite their utility and many clear advantages, SMC clusters have been surprisingly underexploited. The number of clusters with well-determined ages from deep main-sequence photometry is minimal.  Only two clusters have had their detailed chemical abundances derived via high resolution spectroscopy (NGC\\,330 - \\citealt{gon99}, \\citealt{hil99}, NGC\\,121 - \\citealt{joh04}), and an additional six clusters have metallicities that were derived from \\ion{Ca}{2} triplet (CaT) spectroscopy of individual stars (\\citealt{dch98}, hereafter DH98).  Aside from these eight clusters, all other existing abundance determinations are based only on photometry or integrated spectroscopy.  Thus, the really detailed information, viz. accurate ages and abundances, necessary to fully utilize SMC clusters, both   as tracers of the SMC's formation and chemical evolution history, and as templates for studying stellar populations in more distant galaxies, is sorely lacking.  Our group has been working to ameliorate this situation over the past several years.  Using Washington photometry, we have derived ages and metallicities for almost 50 previously unstudied or poorly studied clusters \\citep{pia01,pi05a,pi07a,pi07b,pi07c}.  These results have greatly improved our understanding of the global properties of the SMC cluster system, provided constraints on the age-metallicity relation (AMR), explored possible gradients, cluster formation scenarios, etc.  However, these photometric cluster ages and abundances suffer from two problems.  First, our photometry comes from data obtained with a 1m class telescope, which is barely adequate to reach the main sequence turn off for clusters older than several Gyr at the distance of the SMC.  Second, while the photometry for the red giant stars that are used to derive cluster abundances is generally good, the Washington technique is known to require a significant age correction to metallicities derived for intermediate-age ($\\lesssim$ 5 Gyr) objects \\citep{gei03}, an age range that includes the majority of the SMC clusters so far observed. This is of course illustrative of the infamous age-metallicity degeneracy and is not a problem restricted to the Washington system; while some age and metallicity estimates from isochrone fitting to CMDs constructed in other photometric systems exist, the degeneracy between age and metallicity also makes these estimates inherently uncertain in the absence of more solid metallicity measurements based on spectroscopic data.  As mentioned above, spectroscopic based abundances only exist for a handful of SMC clusters, with most of the sample coming from the work by DH98.  They combined their CaT based metallicities with ages from the literature to create the first accurate  AMR for the SMC, and found that it was consistent with a simple closed box model and did not require the significant gas infall or strong galactic winds that were needed to explain previous SMC AMR's \\citep[e.g.,][]{dop91}.  In addition, DH98 found that their cluster velocities, like other kinematic tracers in the SMC, show no evidence for any systematic rotation of the SMC.  However, their small sample size and large age errors severely limit how well we can constrain the kinematics and chemical evolution of the SMC.  While the CaT method does not measure Fe abundances directly, previous authors have shown that the strength of the CaT lines are an excellent proxy for [Fe/H], and can be used in stellar populations covering a wide range of ages and abundances \\citep[e.g.,][]{col04,car07}. An added benefit of using the CaT is that it is very efficient; not only are RGB stars near their brightest in the infrared, but the multiplexing capabilities of many moderate-resolution spectrographs allows the observation of dozens of stars simultaneously, greatly increasing the probability of identifying cluster members.  In a previous paper \\citep[hereafter G06]{gro06}, we have applied this method to the LMC with excellent results, following up on the pioneering  work by \\citet{ols91}. Using FORS2 on the VLT, we were able to identify more than 200 member stars in 28 populous LMC clusters and determine accurate mean cluster velocities ($\\sigma$ = 1.6 km s$^{-1}$) and metallicities ($\\sigma$ = 0.04 dex) and use these to explore the global cluster metallicity distribution, kinematics, etc.  To further refine our knowledge of the velocities and abundances of clusters in the SMC, we here apply this powerful  technique to the SMC, again using FORS2 on the VLT to obtain medium-resolution near-infrared spectra of the CaT lines in 270 individual RGB stars in and around 16 SMC clusters.  Herein we present our efforts to identify cluster members and derive mean abundances and velocities for these  clusters.  In a future paper we will analyze the several hundred field stars that were also observed as part of this project.  This paper is organized as follows: In \\S 2, we describe our target selection process, and our spectroscopic observations and reduction procedures are detailed in \\S 3.  In \\S 4 and \\S 5, we present the radial velocities and equivalent width measurements and the metallicity derivation, respectively.  The membership selection process is described in \\S 6 and \\S 7 compares our metallicity values with previous determinations. In \\S 7 we also discuss our metallicity results and in \\S 8 the kinematics. Finally, in \\S 9 we summarize our results. ",
        "conclusions": "Magellanic Cloud clusters are an excellent laboratory for helping to unlock the secrets of cluster and galaxy formation. They are also crucial testbeds for stellar and chemical evolution models and interpreting the integrated light of distant galaxies. However, despite their utility and proximity, Small Magellanic Cloud clusters in particular have been overlooked in this regard. In order to help remedy this situation, we have carried out a large-scale investigation of the kinematics and metallicities for a number of SMC star clusters. Building on our experience with deriving these parameters for LMC clusters using the CaT technique (G06), we have used the same technique to explore SMC clusters. We obtained CaT spectra for a number of stars associated with 16 SMC clusters using the VLT + FORS2 MXU instrument. This provides the largest sample of SMC clusters with velocities and well-determined spectroscopic metallicities currently available, more than doubling the only previous study (DH98). Target clusters were selected from our Washington photometric system studies, which provide a rough estimate of age and abundance, to be old enough to possess red giant branch stars and cover a wide area across the galaxy.  We used the same reduction and analysis techniques as we did in G06 to measure stellar radial velocities  and metallicity. Typical radial velocities errors are 7.5 km s$^{-1}$ and metallicity errors are 0.17 dex per star. Cluster members are selected using an analysis combining their location in the CMD, position in the cluster, and radial velocities and metallicity with respect to other stars in the same field. We determine mean cluster velocities to typically 2.7 km s$^{-1}$ and metallicities to 0.05 dex (random error), from a mean of 6.4 members per cluster.  A comparison of our mean cluster metallicities with those derived from Washington photometry for 11 clusters in common shows very good overall agreement, with the CaT metallicities being much more precise. This indicates that the Washington metallicities,  and especially their procedure used to correct the age dependence of their metallicities for their age dependence, is appropriate.  The metallicity distribution (MD), metallicity gradient and age-metallicity relation (AMR) are investigated, combining our clusters with those observed by DH98 and Glatt et al. (also with CaT) and the one cluster with a detailed, high resolution metallicity to compile a sample of 25 clusters on a homogenous metallicity scale with relatively small errors, although the ages are somewhat heterogeneous. The mean metallicity is [Fe/H] $=-0.96$, with $\\sigma = 0.19$. Dividing the sample into two age groups at 3 Gyr, the 12 older clusters have a mean metallicity of $-1.08, \\sigma = 0.17$, while the 13 younger have $-0.85, 0.15$. Most clusters lie between [Fe/H] = $-$0.75 and $-$1.25. There is a suggestion for bimodality in the MD, with peaks at [Fe/H] $\\sim -0.9$ and $-$1.15. No clear gradient is seen with distance from the center of the SMC. However, intermediate-age SMC clusters are both significantly more metal-poor and have a larger metallicity spread than their LMC counterparts. The AMR shows evidence for 3 phases: a very early ($>11$ Gyr) phase in which the metallicity reached $\\sim -1.2$, a long intermediate phase from $\\sim 10-3$ Gyr in which the metallicity only slightly increased although a number of clusters formed, and a final phase from 3-1 Gyr ago in which the rate of enrichment was substantially faster. These salient features agree with those found by most other AMR studies using other tracers and techniques. We find good overall agreement with the model of PT98 which assumes a burst of star formation at 4 Gyr. A hybrid infall $+$ outflow model of \\citet{car05} also  fits the data reasonably well. The simple closed box model of DH98 yields a much poorer fit, and the AMRs derived by \\citet{har04}  and \\citet{idi07}  are significantly offset to higher metallicities for intermediate-age clusters. A number of different lines of evidence point to the likelihood of a burst in the SMC star and cluster formation about 3 Gyr ago. The cause of such a burst is currently a source of much speculation. The suggestion by \\citet{bek04} that it is due to a close passage of the SMC and LMC is intriguing but requires better knowledge of their orbits, especially proper motions, to be definitively tested.  We finally examine the kinematics of our CaT clusters. Their mean radial velocity is = 148 km s$^{-1}$, with a velocity dispersion of 23.6 km s$^{-1}$. These values are in very good agreement with those found by DH98 from their smaller sample. Combining the 2 cluster samples, we find the kinematics are dominated by the velocity dispersion, as found in virtually all other kinematic studies of a wide variety of SMC populations.\\\\"
    },
    "0808/0808.0189_arXiv.txt": {
        "abstract": "Exploring the diversity of dark energy dynamics, we discover a calibration relation, a uniform stretching of the amplitude of the equation of state time variation with scale factor.  This defines homogeneous families of dark energy physics.  The calibration factor has a close relation to the standard time variation parameter $w_a$, and we show that the new, calibrated $w_a$ describes observables, i.e.\\ distance and Hubble parameter as a function of redshift, typically to an accuracy level of $10^{-3}$.  We discuss implications for figures of merit for dark energy science programs. ",
        "introduction": "}  Understanding the nature of the dark energy accelerating the cosmic expansion is one of the premier questions in physics.  The answer offers the possibility of deep insights into the nature of spacetime and gravity, extra dimensions, the quantum vacuum, and possibly the unification of gravitation and quantum physics.  Precision mapping of the expansion history provides one path to characterizing the dark energy, in particular its equation of state and time variation.  Guidance from theory is useful to predict observable signatures for cosmological probes such as distance and Hubble parameter measurements, in particular what level of accuracy is required to distinguish between models.  From a model one can predict distance-redshift relations etc.\\ but the number of models is vast; one would like to identify model independent or at least generic characteristics of the dark energy. Indeed, such properties exist, as discussed in detail recently by \\cite{cahndl}, for classes of behavior in the early time evolution of dark energy, valid for $z\\gtrsim2$ when the dark energy does not strongly affect the background expansion.  In this article we seek to extend characterization of the dark energy properties in terms of the equation of state to the entire observable history.  This requires a different approach, calibrating the evolution through a ``stretch'' relation between the amplitude of the time variation and the time variable or scale factor of the expansion. The calibration then provides a physical basis for a compact and highly accurate parametrization of the dark energy influence on observables.  In \\S\\ref{sec:dyn} we examine several diverse models, looking for similarities and distinctions.  We introduce the calibration in \\S\\ref{sec:stretch} and discuss its relation to a standard parametrization of the equation of state.  \\S\\ref{sec:obs} examines the utility of the description and shows that it achieves robustness and accuracy at the $10^{-3}$ level, sufficient for next generation data.  We discuss some implications for figures of merit of dark energy science programs in \\S\\ref{sec:fom}. Those readers wanting to get right to the results could start in the middle of \\S\\ref{sec:stretch}. ",
        "conclusions": "}  Having investigated a diverse group of dark energy models to explain the acceleration of the cosmic expansion, we find a homogeneous ``stretch'' relation that calibrates the time variation behavior into tight families.  This stretch factor is closely related to the standard time variation measure $w_a$, and we verify that the equation of state form $w(a)=w_0+w_a(1-a)$, with $w_a$ now treated as a fit parameter to observables, delivers fractional accuracy at the $10^{-3}$ level.  Such accuracy is sufficient for next generation data and the $w_0$-$w_a$ form can be viewed as an appropriate compression of the expansion history information that can be extracted from such observations.  That is, this form neither overcompresses (loses important information) nor undercompresses (lacks additional leverage).  This indicates there is no need nor generic benefit for going to a third parameter. Note that \\cite{barnardchi} saw similar compression and tight relations within a principal component analysis relying on many modes.  To gain insight into the nature of dark energy, particular combinations of $w_0$-$w_a$ may have enhanced leverage and hence merit, separating the cosmological constant from the thawing class, each from the freezing class, and possibly zeroing in on specific models within a class.  The calibration, and its robustness and accuracy in accounting for the observable relations, offers a well-defined method for assessing the next generation dark energy science program. Interpretation of those observations should offer promising insights into the physics of the accelerating universe."
    },
    "0808/0808.2529_arXiv.txt": {
        "abstract": "We present $H$-band (1.65$\\mu$m) surface photometry of 57 galaxies drawn from the Local Sphere of Influence (LSI) with distances of less than 10\\,Mpc from the Milky Way. The images with a typical surface brightness limit 4 mag fainter than 2MASS ($24.5$\\,mag arcsec${}^{-2}< \\mu_{lim}<26$\\,mag arcsec${}^{-2}$ ) have been obtained with IRIS2 on the 3.9~m Anglo-Australian Telescope. A total of 22 galaxies that remained previously undetected in the near-IR and potentially could have been genuinely young galaxies were found to have an old stellar population with a star density $1-2$ magnitudes below the 2MASS detection threshold. The cleaned near-IR images reveal the morphology and extent of many of the galaxies for the first time. For all program galaxies, we derive radial luminosity profiles, ellipticities, and position angles, together with global parameters such as  total magnitude, mean effective surface brightness and half-light radius. Our results show that 2MASS underestimates the total magnitude of galaxies with $\\langle\\mu_H\\rangle_{eff}$ between $18-21$\\,mag arcsec${}^{-2}$ by up to 2.5\\,mag. The S\\'ersic parameters best describing the observed surface brightness profiles are also presented. Adopting accurate galaxy distances and a $H$-band mass-to-light ratio of  $\\Upsilon_{\\ast}^H=1.0\\pm 0.4$, the LSI galaxies are found to cover a stellar mass range of $5.6<\\log_{10}(\\mathcal{M}_{\\rm stars})<11.1$. The results are discussed along with previously obtained optical data. Our sample of low luminosity galaxies is found to follow closely the optical-infrared $B$ versus $H$ luminosity relation defined by brighter galaxies with a slope of $1.14 \\pm 0.02$ and scatter of $0.3$\\,magnitudes. Finally we analyse the luminosity -- surface brightness relation to determine an empirical mass-to-light ratio of $\\Upsilon_{\\ast}^H=0.78\\pm0.08$ for late-type galaxies in the $H$-band. ",
        "introduction": "The observational properties of nearby galaxies such as fluxes, colours, morphologies and sizes reflect their underlying physical properties (stellar/baryonic and dark matter content, star formation rates, formation history and angular momenta). Exactly how these observational and physical properties are related is still poorly understood. By technical necessity, the observational quantities are mainly based on the optical $B$-band ($390 - 480$\\,nm). However galaxies evolving in low density environments with little external stimulation for star formation often contain significant quantities of dust (eg.  see \\citealt{driver07}) which can attenuate and distort their optical light profiles.  In contrast, dust attenuation is vastly reduced at near-IR wavelengths and hence  the near-IR provides a spectral regime where a more accurate, unaltered representation of a galaxy's underlying stellar distribution can be obtained \\citep{gavazzi96}.  Furthermore, the stellar mass of most galaxies is dominated by the quiescent old stellar component whose energy output peaks at near-IR wavelengths.  Even in the extreme case of Blue Compact Dwarf (BCD) galaxies, previously thought to be primeval galaxies forming their first stars at the present epoch \\citep{thuan97}, the analysis of their resolved stellar populations has revealed the presence of stars at least a few Gyrs of age (\\citealt[eg. the BCD galaxies: VII Zw 403, Mrk 178 and I Zw 36 as discussed by][respectively]{schulteladbeck98,schulteladbeck00, schulteladbeck01}; SBS 1415+437 discussed by \\citealt{aloisi05}; I Zw 18 by \\citealt{aloisi07}; and CGCG 269-049 by \\citealt{corbin08}).  In order to obtain a deeper understanding of the connection between the light and matter distribution in galaxies, a representative sample of nearby stellar systems needs to be studied in detail. The Local Sphere of Influence (LSI, $D< 10\\,$Mpc) contains large numbers of early (dE) and late-type (dIrr) dwarf galaxies that make up about 85\\% of the local galaxy population~\\citep{kraan79, schmidt92, karachentsev04}. Dwarf galaxies contribute about 4\\% to the local luminosity density and about 10-15\\% to the local H\\,{\\sc i} mass density \\citep{karachentsev04}. Due to their proximity to the Milky Way, LSI galaxies are ideal for a near-IR study which includes significant numbers of dwarf systems.  Previous near-IR surveys include the Two Micron All Sky Survey \\citep[2MASS,][]{skrutskie06} as well as deeper targeted galaxy surveys \\citep{gavazzi96b, gavazzi96a, gavazzi00, boselli00}.  2MASS photometry for galaxies suffers from  a number of important drawbacks that are becoming more evident as the samples of independently investigated galaxies become larger. The short integration time of 2MASS observations resulted in most of the low surface brightness (LSB) dwarfs in the LSI remaining undetected, and if they were detected, 2MASS underestimated the fluxes by as much as 70\\% \\citep{andreon02}. The targeted $H$-band observations of \\cite{gavazzi96b, gavazzi96a, gavazzi00} and  \\cite{boselli00} were inherently deeper however the samples included few LSB dwarfs. This serious limitation demands a deeper and higher resolution study to investigate those galaxies that were beyond the reach of photometric near-IR studies to date.  A reference atlas of images needs to have the necessary spatial resolution to probe the morphological fine-structure of these nearby galaxies and contain a significant number of dwarf galaxies that are generally overlooked. A LSI sample has the additional advantage that an increasingly large number of nearby galaxies have accurately known distances. \\cite{karachentsev06} report that 214 out of 451 LSI galaxies have  distance estimates (with less than 10\\% uncertainty) by means of the tip magnitude of the red giant branch (TRGB), the Tully-Fisher relation, and the surface brightness fluctuations (SBF) method (see  for example \\citealt{jerjen98, jerjen01} and \\citealt{karachentsev04}). The remaining galaxies have rough distance estimates from the luminosity of their brightest stars, radial velocities or their suspected membership to a known galaxy group.  The purpose of this paper is to present a near-IR $H$-band (1.65$\\mu$m) atlas of 57 LSI galaxies, probing to flux levels approximately 4 mag arcsec${}^{-2}$ or 40 times fainter than 2MASS. The majority of the galaxies presented here are much fainter than those in previous targeted surveys. We derive photometric parameters for  each object such as the total magnitude, the effective radius and effective surface brightness, S\\'ersic fitting parameters, etc. Using the best distances currently available in the literature allows us to derive physical parameters such as their  luminosities and  stellar masses.  The paper is organised as follows: we describe the sample selection  in \\S \\ref{s:sample}. In \\S\\ref{s:obs} and \\S\\ref{s:photometry} we discuss the observing strategies,  the data reduction, and the photometric calibration of the images.   The 11 galaxies in the sample which remained undetected at our faint detection limit or had images which could not be usefully analysed are discussed in \\S\\ref{s:nondetect}. The new data is compared to  2MASS photometry and optical ($B$-band) data in \\S\\ref{s:results} and the luminosity - surface brightness relation discussed. Interesting properties of individual galaxies are described in \\S\\ref{s:interesting}. Finally,  the results are summarised in \\S\\ref{s:summary}.  ~\\linebreak ",
        "conclusions": "\\label{s:summary}  We have presented the deepest $H$-band images available to date for 57 galaxies in the Local Sphere of Influence ($D<10$\\,Mpc), obtained using the near-IR camera IRIS2 at the 3.9m Anglo-Australian Telescope. Of the 68 targets, 11 remained undetected or could not be usefully analysed due to contamination by foreground stars. The surface brightness limit reaches down to $\\mu_{lim}<26$\\,mag\\,arcsec, 4 magnitudes fainter than 2MASS.  The images, cleaned from Galactic foreground contamination, reveal the morphology and extent of many of the galaxies for the first time. For 56 galaxies, we derive radial luminosity profiles, ellipticities, and position angles, together with global parameters such as  total magnitude, mean effective surface brightness, half-light radius, S\\'ersic parameters, and stellar mass.  No genuine young galaxies have been found in this survey. Some sample galaxies were previously identified on $B$-band photographic plates but remain undetected in the near-IR. In each case there is a plausible alternative explanation for the non-detection: \\begin{itemize} \\item AM0717-571: DSS $B_J$-band morphology resembles that of a Galactic nebula, but true nature still remains unclear. \\item HIZOAJ1616-55 and SJK98 J1616-55: possibly one or two high velocity clouds. \\item KK2000-03: Superimposed star hampers analysis however the marginal detection in the H-band suggests an unusual blue galaxy. \\item KK2000-04: Originally assumed to be a companion of NGC1313 however possibly a photographic plate flaw. \\item KK2000-06: Originally assumed to be a companion of NGC1313. More likely a background galaxy at $\\approx$2250\\,km\\,s$^{-1}$. \\item NGC2784 DW1: intrinsic extreme low surface brightness dwarf satellite of NGC2784. \\end{itemize}  We also detected a double nucleus in KKS2000-09 and propose to reclassify this system as a peculiar galaxy. KKS2000-25 was shown to have distinct spiral arms in the $H$-band and thus should be classified as ``Sb\". Morphology and angular size strongly suggest that this is a background galaxy beyond 10\\,Mpc.  We found compelling evidence that the short integration time of 2MASS resulted in serious underestimation of a galaxy's luminosity. The magnitudes of galaxies, with  $H$-band surface brightnesses fainter than 18\\,mag\\,arcsec${}^{-2}$, obtained in our study are up to 2.5\\,mag brighter than those obtained by 2MASS. As the mean effective surface brightness correlates with the luminosity of a galaxy, we expect serious selection biases for a 2MASS-based $H$-band galaxy luminosity function fainter than M$_{H}=-20$\\,mag.    There is a tight correlation (correlation coefficient = 0.97) between the $B$- and $H$-band magnitudes of a galaxy and this correlation has been demonstrated over a range of 15 magnitudes. The linear transformation between the $B$- and $H$-bands has a small scatter (0.3 mag) for bright galaxies. In the dwarf regime, there is a marginal  increase in scatter and possibly a slight trend for galaxies to be redder (by approximately 1 magnitude) than indicated by the transformation found for bright galaxies.  The galaxy luminosity -- mean effective surface brightness  relation has been analysed to derive a semi-empirical stellar mass-to-light ratio of $\\Upsilon_{\\ast}^H=0.78\\pm0.08$ in the $H$-band.  All raw and reduced $H$-band images of the 57 program galaxies in this near-IR survey will be made publicly available and can be obtained via email request."
    },
    "0808/0808.3005_arXiv.txt": {
        "abstract": "This paper presents an analytic perturbation approach to the dynamics of a classical spinning particle, according to the Mathisson-Papapetrou-Dixon (MPD) equations of motion, with a direct application to circular motion around a Kerr black hole. The formalism is established in terms of a power series expansion with respect to the particle's spin magnitude, where the particle's kinematic and dynamical degrees are expressed in a completely general form that can be constructed to infinite order in the expansion parameter. It is further shown that the particle's squared mass and spin magnitude can shift due to a classical analogue of radiative corrections that arise from spin-curvature coupling. Explicit expressions are determined for the case of circular motion near the event horizon a Kerr black hole, where the mass and spin shift contributions are dependent on the initial conditions of the particle's spin orientation. A preliminary analysis of the stability properties of the orbital motion in the Kerr background due to spin-curvature interactions is explored and briefly discussed. ",
        "introduction": "\\label{sec:1}  One of the earliest and on-going research interests in general relativity concerns the dynamics of extended bodies in the presence of strong gravitational backgrounds. Considering that virtually all astrophysical objects in the Universe, such as black holes, neutron stars, and other isolated massive bodies, have at least some spin angular momentum in their formation, it is not difficult to surmise that an in-depth study of moving relativistic systems with spin is a useful endeavour. A relevant example concerns the motion of rapidly rotating neutron stars in circular orbit around supermassive black holes like ones believed to exist in the centre of galaxies, which serve as candidate sources for emitting low-frequency gravitational wave radiation that may be detected by the space-based LISA gravitational wave observatory \\cite{LISA}.  A first attempt to understand the dynamics of extended bodies in curved space-time was put forward by Mathisson \\cite{Mathisson}, who showed the existence of an interaction term involving the direct coupling of particle spin to the Riemann curvature tensor generated by a background source. Steady progress was made since this first attempt, with a notable contribution made several years afterwards by Papapetrou \\cite{Papapetrou}, who proposed that the spinning particle exists within a space-time world tube containing its centre-of-mass worldline, where its associated matter field has compact support. In addition, multipole moment contributions, i.e. beyond the mass monopole and spin dipole, to the extended objects full equations of motion were considered by Tulczyjew \\cite{Tulczyjew} and others, ultimately leading to the expressions obtained by Dixon \\cite{Dixon1,Dixon2}, with a self-consistent description for all multipole moment contributions to infinite order. While the various theories of extended body motion in curved space-time differ with respect to the higher-order multipole moments, all of them recover the ``pole-dipole approximation'' identified initially by Mathisson and Papapetrou, which are satisfactory for most practical calculations, so long as the dimensions of the spinning body are small when compared to the background space-time's local radius of curvature. These truncated expressions of the full equations of motion are commonly known as the Mathisson-Papapetrou-Dixon (MPD) equations.  There has been widespread interest in applying the MPD equations to the dynamics of classical spinning particles in orbit around rotating black holes, as described by the Kerr metric \\cite{Mashhoon1,Wald,Tod,Semerak,Suzuki1}. In many ways, the Kerr background is an ideal testing ground for the MPD equations, since both mass sources are spinning, which introduce interesting spin-curvature effects that impact upon the orbiting particle's overall evolution. Furthermore, it lends itself well to numerical simulations of deterministic chaos under extreme conditions \\cite{Suzuki1,Suzuki2,Hartl1,Hartl2}, as well as studies of gravitational wave generation \\cite{Mino,Tanaka} arising from spin-induced deviations away from geodesic motion.  More formal study of the MPD equations have also occurred in various forms \\cite{Ehlers,Bailey,Noonan}, including a recent perturbative approach developed by Chicone, Mashhoon, and Punsly (CMP) \\cite{Chicone}, with application to the study of rotating plasma clumps propagating in astrophysical jets directed along a Kerr black hole's axis of symmetry. Another application of the CMP approximation by Mashhoon and Singh \\cite{Mashhoon2} determined analytic expressions for leading-order spin-curvature perturbations of a spinning particle's circular orbit around a Kerr black hole. This analysis is successful in reproducing the spinning particle's kinematic behaviour compared to numerical simulations of the full MPD equations for situations where, for spin magnitude $s$ and mass $m$, the M{\\o}ller radius \\cite{Mashhoon2,Moller} for the spinning particle is $s/m \\lesssim 10^{-3} \\, r$, and $r$ is the particle's radial distance away from the background mass source. However, this approximation starts to break down when $s/(m r) \\sim 10^{-2}-10^{-1}$ for $r = 10 \\, M$, where $M$ is the Kerr black hole mass, suggesting that higher-order spin-curvature coupling terms are required to more completely describe the orbital motion.  It was for this initial purpose that a generalization of the CMP approximation was very recently introduced by Singh \\cite{Singh0} to incorporate higher-order analytic contributions to the perturbation approach for the MPD equations. This generalization has several nice features. For example, as a power series expansion with respect to the particle's spin magnitude, it can be extended to formally {\\em infinite order} in the expansion. In addition, it leads to expressions that are background independent, and is fully applicable to {\\em arbitrary motion} of the particle, without recourse to any space-time symmetries within the metric. As a result, this generalization is very robust, with applicability for many distinct scenarios in theoretical astrophysics, such as the modelling of globular clusters and other many-body dynamical systems in curved space-time, and also spinning particle interactions with gravitational waves, the results of which can be compared with existing treatments \\cite{Nieto,Mohseni,Kessari}. Furthermore, this approach identifies the existence of a classical analogue for ``radiative corrections'' that shift the particle's overall squared mass and spin magnitude due to higher-order spin-curvature contributions, a feature not thought about before. It would, therefore, be very useful to investigate the computational capacity of this generalization when applied to circular motion in the Kerr background. This is especially so in extreme conditions where a transition from stable to chaotic motion may be analytically identified, for comparison with existing approaches \\cite{Suzuki1,Suzuki2,Hartl1,Hartl2} which use primarily numerical methods.  The purpose of this paper is to present the generalized form of the CMP approximation for the MPD equations within the context of circular motion around a Kerr black hole, and explore the derived physical consequences. It begins with Sec.~\\ref{sec:2}, which displays the full MPD equations, followed by a presentation of the formalism behind the generalized CMP approximation in Sec.~\\ref{sec:3}. Afterwards, Sec.~\\ref{sec:4} presents the formal application of the generalized CMP approximation to the case of circular motion around a Kerr black hole, up to second-order in the perturbation expansion parameter. This is followed, in Sec.~\\ref{sec:5}, by analysis of the predicted kinematic and dynamical properties of the perturbed system, including the predicted effective squared mass and spin magnitude of the spinning particle. A general discussion of the main results obtained in this paper is found in Sec.~\\ref{sec:6}, with a brief conclusion thereafter. The metric convention adopted is $+2$ signature with Riemann and Ricci tensor definitions following MTW \\cite{MTW}, and geometric units of $G = c =1$ are assumed throughout. ",
        "conclusions": "\\label{sec:6}  This paper outlines the generalization of an analytic perturbation approach to the Mathisson-Papapetrou-Dixon equations for a spinning point particle, first introduced by Chicone, Mashhoon, and Punsly, with an application to circular motion around a Kerr black hole. The formalism shows the existence of ``radiative corrections'' to the particle's squared mass and spin magnitudes due to spin-curvature interactions, represented in power series expansion form. In performing the analysis, it is possible to semi-analytically identify the emergence of instabilities during the particle's orbital motion, which serves as a basis for a more precise treatment in the future.  One of the underlying goals of the formalism presented in this paper is to determine the perturbed orbit of the spinning particle according to the generalized CMP approximation, following the approach taken earlier \\cite{Mashhoon2}. However, to do this properly requires a modification of the equations of motion to incorporate dissipative effects due to gravitational radiation, which have not yet been taken into account. Such a modification would most certainly require evaluation of the Teukolsky equations for determining the radiation effects corresponding to an adiabatic inspiral for the spinning particle's orbit. This is a non-trivial exercise with both conceptual and technical challenges to still overcome. Once this is better understood, a determination of the perturbed orbit due to spin-curvature interactions will be considered in a future publication. For now, a second paper on the generalized CMP approximation in the Vaidya background is forthcoming \\cite{Singh2} as a companion piece to accompany and compare with this paper."
    },
    "0808/0808.1236_arXiv.txt": {
        "abstract": "We discuss the anisotropic arrival directions of the ultra high energy cosmic rays detected by Auger which I consider one of the biggest discoverie in astrophysics during the last year. ",
        "introduction": "As you may conclude from the different concluding remarks at this meeting there are various approaches to doing them. Most of my esteemed colleagues spoke on several topic that they consider important now and in the future. I decided to concentrate on one single topic that excited not only me but a large number of astrophysicists that were never interested in ultra high energy cosmic rays (UHECR) - the observed correlation of the highest energy events detected by the Pierre Auger Observatory in Argentina with active galactic nuclei (AGN)~\\cite{Auger1,Auger2}. This discovery marked the beginning of a new type of astronomy - cosmic ray astronomy - that may become as important in the future as the TeV gamma ray astronomy has recently become.  TeV gamma ray astronomy does not only reveal the type of the gamma ray sources, it also provides an indirect measurement of the infrared/optical background whose estimates are based on theoretical calculations involving the light emission of different types of stars and galaxies and the cosmological evolution of these objects and the matter density around them in the Universe. Cosmic ray astronomy also involves some general features of the nearby Universe that are otherwise extremely difficult to measure such as the local (within 200 Mpc) distribution of matter and different powerful astrophysical objects, the average strength of the intergalactic magnetic fields (that could even be mapped when the sources are confirmed) and the galactic ones.  I will discuss these results and report on the small contribution that I have made to this topic. It is not connected to the type of the UHECR sources, but includes a better than 3$\\sigma$ confirmation of the anisotropy detected by Auger~\\cite{Stanev08}. ",
        "conclusions": ""
    },
    "0808/0808.3887_arXiv.txt": {
        "abstract": "We review the properties of the very young ($\\sim 2$\\,Myr) open cluster NGC\\,6383. The cluster is dominated by the massive binary HD\\,159176 (O7\\,V + O7\\,V). The distance to NGC\\,6383 is consistently found to be $1.3 \\pm 0.1$\\,kpc and the average reddening is determined to be $E(\\bv) = 0.32 \\pm 0.02$. Several pre-main sequence candidates have been identified using different criteria relying on the detection of emission lines, infrared excesses, photometric variability and X-ray emission. ",
        "introduction": "NGC\\,6383 ($\\alpha_{2000} = 17^h34^m48^s$, $\\delta_{2000} = -32^{\\circ}34\\farcm0$; $l_{II} = 355.69^{\\circ}$, $b_{II} = +0.04^{\\circ}$) is a rather compact open cluster which could be part of the Sgr\\,OB1 association together with NGC\\,6530 and NGC\\,6531. The cluster was originally discovered by John Herschel in 1834, and listed as h\\,3689 in Herschel's Cape Catalogue published in 1847. Most probably due to a clerical error, it also appears in Dreyer's New General Catalogue as NGC\\,6374, an identifier that is however not commonly used.  Trumpler \\cite{Trumpler} pointed out that NGC\\,6383 belongs to a category of clusters in which a few dozen faint stars are closely grouped around a bright hot central star that frequently happens to be a binary system. In the case of NGC\\,6383, the cluster is centered on the O-type binary HD\\,159176 (m$_V$ = 5.7, see Fig.\\,\\ref{optical}) that dominates the emission from the cluster over a broad range of energies from the near-infrared to the X-ray domain. Apart from HD\\,159176, the cluster currently harbours no stars earlier than B1, although de Wit et al.\\ \\cite{deWit} suggested that the eclipsing binary candidate HD\\,158186 (O9.5\\,V, Marchenko et al.\\ 1998) might have been ejected from NGC\\,6383 through dynamical interactions in the cluster core.  NGC\\,6383 and more specifically HD\\,159176 are likely responsible for the ionization of the H\\,{\\sc ii} region RCW\\,132 (Rodgers, Campbell \\& Whiteoak 1960) also known as S\\,11 (Sharpless 1953)\\footnote{Note that the S\\,11 identifier is not to be confused with the number of the nebula in the revised catalogue of H\\,{\\sc ii} regions published by Sharpless \\cite{Sharp2}. In the latter catalogue, the nebula corresponds to entry number 12 (in the SIMBAD database, this identifier is referred to as Sh 2-012).} or Stromlo\\,67 (Gum 1955). Rodgers et al.\\ \\cite{RCW} described RCW\\,132 as a 110\\,arcmin $\\times$ 80\\,arcmin medium brightness crescent-shaped region. Images of this emission nebula can be found for instance on plate 131 of Lyng\\aa\\ \\& Hansson \\cite{LH}.  During a survey of the galactic ridge at 1390\\,MHz, Westerhout \\cite{Westerhout} detected a large ($1.5^{\\circ} \\times 1.5^{\\circ}$) ring-like radio structure roughly centered on NGC\\,6383. However, Westerhout cautioned that this feature was difficult to separate from the background emission. Part of this source is likely due to RCW\\,132.  \\begin{figure}[htb] \\centering \\includegraphics[draft=False,width=\\textwidth]{ngc6383b.eps} \\caption{The region around the core of NGC\\,6383 as seen on the UKSTU survey red plate. The bright star in the middle of the field of view is the massive binary system HD\\,159176. The field size is 10\\,arcmin $\\times$ 10\\,arcmin.\\label{optical}} \\end{figure} ",
        "conclusions": ""
    },
    "0808/0808.4053_arXiv.txt": {
        "abstract": "{} {We report the $\\gamma$-ray activity from the intermediate BL Lac S5 0716+714 during observations acquired by the AGILE satellite in September and October 2007. These detections of activity were contemporaneous with a period of intense optical activity, which was monitored by GASP--WEBT. This simultaneous optical and $\\gamma$-ray coverage allows us to study in detail the light curves, time lags, $\\gamma$-ray photon spectrum, and Spectral Energy Distributions (SEDs) during different states of activity.} {AGILE observed the source with its two co-aligned imagers, the Gamma-Ray Imaging Detector (GRID) and the hard X-ray imager (Super-AGILE), which are sensitive to the 30~MeV--50~GeV and 18--60 keV energy ranges, respectively. Observations were completed in two different periods, the first between 2007 September 4 -- 23, and the second between 2007 October 24 -- November 1.} {Over the period 2007 September 7 -- 12, AGILE detected $\\gamma$-ray emission from the source at a significance level of 9.6-$\\sigma$ with an average flux (E$>$100~MeV) of $(97 \\pm 15) \\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$, which increased by a factor of at least four within three days. No emission was detected by Super-AGILE for the energy range 18--60~keV to a 3-$\\sigma$ upper limit of 10 mCrab in 335 ksec. In October 2007, AGILE repointed toward S5 0716+714 following an intense optical flare, measuring an average flux of $(47 \\pm 11) \\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$ at a significance level of 6.0-$\\sigma$. } { The $\\gamma$-ray flux of S5 0716+714 detected by AGILE is the highest ever detected for this blazar and one of the most intense $\\gamma$-ray fluxes detected from a BL Lac object. The SED of mid-September appears to be consistent with the synchrotron self-Compton (SSC) emission model, but only by including two SSC components of different variabilities.} ",
        "introduction": "\\label{introduction}  The source S5 0716+714 was discovered in 1979 as the optical counterpart to an extragalactic radio source (K\\\"uhr et al. 1981).  Two years later, it was classified as a BL Lac (Biermann et al. 1981) because of its featureless optical spectrum and high linear polarization. The optical continuum was so featureless that every attempt to determine its spectroscopic redshift has failed. However, by optical imaging of the underlying galaxy, Nilsson et al. (2008) derived a redshift of z = 0.31 $\\pm$ 0.08.  According to its spectral energy distribution the source belongs to the intermediate BL Lac class, observations by BeppoSAX (Tagliaferri et al. 2003, Giommi et al. 1999) and XMM-$Newton$ (Foschini et al. 2006, Ferrero et al. 2007) provide evidence for a concave X-ray spectrum in the 0.1--10~keV band, which is a signature of the presence of both the steep tail of the synchrotron emission and the flat part of the Inverse Compton spectrum. The detection in the X-ray band of fast variability only in the soft X-ray component can be interpreted as a slowly variable Compton component and a fast, erratic variable tail of the synchrotron component.  In general, the variability of this blazar is strong in every band on both long and short (intraday) timescales. The optical and radio historical behaviour was analyzed by Raiteri et al. (2003), while the EGRET telescope on the Compton Gamma-Ray Observatory (Hartman et al. 1999) detected S5 0716+714 several times in the $\\gamma$-rays (Lin et al. 1995; von Montigny et al. 1995). The integrated flux above 100 MeV varied between (13 $\\pm$ 5) and (53 $\\pm$ 13) x 10$^{-8}$ photons cm$^{-2}$ s$^{-1}$.  In this Letter, we present the analysis of the AGILE data obtained during the S5 0716+714 observations in the period 2007 September -- October, in particular two flaring episodes: the first in mid-September, the other on 2007 October 22-23. Preliminary results were communicated in Giuliani et al. (2007).  The intense $\\gamma$-ray flare detected by AGILE in mid-September triggered observations by the GLAST-AGILE Support Program (GASP) of the WEBT\\footnote[1]{\\texttt{http://www.oato.inaf.it/blazars/webt/}\\\\ see e.g. Villata et al. (2006, 2007); Raiteri et al. (2006, 2007).} (see Carosati et al., 2007). About one month later, the GASP observed a bright phase of the source, triggering new AGILE and Swift observations. In the period from September to October 2007, S5 0716+714 showed intense activity with strong optical flaring episodes and a rare contemporaneous optical-radio outburst (Villata et al. 2008).  The results of a multiwavelength campaign on S5 0716+714 with simultaneous Swift and AGILE observations in October 2007 are discussed in Giommi et al. (2008). Throughout this paper, the quoted uncertainties are given at the 1--$\\sigma$ significance level, unless otherwise stated. ",
        "conclusions": "\\label{0716:discussion}  To analyze the gamma-optical correlation, we have applied the Discrete Correlation Function (DCF; see Edelson $\\&$ Krolick (1988) and Hufnagel $\\&$ Bregman (1992)) to the $\\gamma$-ray and $R$-band light curves. The DCF is a statistical method developed to analyze unevenly sampled data sets. The $R$-band flux densities were averaged over 0.1 day bins to smooth the intranight variability. The result is shown in Fig. 3. The DCF displays a significant peak (DCF $\\sim$ 0.9) for a time-lag of -1 day. Notwithstanding the large uncertainty due to poor $\\gamma$-ray sampling, this result suggests a possible delay in the $\\gamma$-ray flux variations with respect to optical variations of the order of 1 day. The uncertainty in the delay can be estimated by Monte Carlo simulations based on the ``flux randomization / random subset selection'' method (see Peterson et al. (2001) and Raiteri et al. (2003)). By performing 2000 simulations we derived a 1-$\\sigma$ uncertainty level in the lag of 1.1 days.  \\begin{figure}[!t] % \\centering \\includegraphics[angle=0,scale=0.40]{fig4.eps} \\caption[0716 DCFsep]{Discrete correlation function (DCF) between the $\\gamma$-ray and $R$-band light curves for S5 0716+714 in September-October 2007.  \\label{0716:fig:dcs}} \\end{figure}  In Fig. 1, it is clear that most of the DCF signal originates from the quasi-simultaneity of the $\\gamma$-ray and optical peaks of late October (JD $\\sim$ 2454396-397). As for the September AGILE detection, the strong $\\gamma$-ray flare lacks strictly simultaneous optical observations since it occurred at both the start of the GASP operation and the optical observing season.  \\begin{table}[!b] \\begin{center} \\caption{Parameters for the two SSC components.} \\begin{tabular}{|l|lll|} \\hline  &1$^{st}$ SSC comp &2$^{nd}$ SSC comp &Units \\\\ \\hline $\\delta$             &14         &25     &       \\\\ $\\Gamma$             &7.5       &15      &       \\\\ $R$                  &40        &40     & [$10^{15}$ cm]   \\\\ $B$                  &1         &0.5      &  [G]    \\\\ $\\gamma_{\\rm min}$   &200       &3 x 10$^{3}$     &     \\\\ $\\gamma_{\\rm break}$ &4 x 10$^{3}$  &6 x 10$^{3}$     &     \\\\ $p_{\\rm  low}$        &2.0      &2.0    &     \\\\ $p_{\\rm  high}$       &4.8      &4.8    &     \\\\ $n_{\\rm e}$          &2.2         &0.8   & [cm$^{-3}$] \\\\ $\\theta$             & 2         & 2          &  [deg] \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  We note that when the $\\gamma$-ray fluxes are $\\la$ 120 $\\times$ 10$^{-8}$\\, photons cm$^{-2}$ s$^{-1}$, the corresponding optical flux densities are around 25--30 mJy. In contrast, the October $\\gamma$-ray peak reaching $\\sim$ 200 $\\times$ 10$^{-8}$\\, photons cm$^{-2}$ s$^{-1}$ has an optical counterpart of 40--45 mJy (see Fig. 1). This suggests that a significant optical event occurred at the same time as the $\\gamma$-ray flare and in September was missed. Moreover, while the ratio between the high and low $\\gamma$-ray flux levels is about 2.5, in the optical band the same ratio is of the order of 1.5. Hence, the gamma variability appears to depend on the square of changes in optical flux density. This would favour a SSC interpretation, in which the emission at the synchrotron and IC peaks is produced by the same electron population, which self-scatters the synchrotron photons. The 1-day time-lag in the high-frequency peak emission found from the DCF could then be due to the light travel time of the synchrotron seed photons that scatter the energetic electrons.  The Spectral Energy Distribution for the AGILE and GASP-WEBT data of September 2007 is shown in Fig.~\\ref{0716:fig:spectrum} as green dots. The blue dashed line shows a simple SSC model that fits simultaneous observations of a ground state (see Tagliaferri et al. 2003 and references therein) and non-simultaneous EGRET data (empty blue circles). Because the high state of mid-September 2007 cannot be fitted by a one-zone SSC component alone, we used a model with two SSC components.  To the first SSC component that reproduces the ground state, we add a second SSC component that dominates the optical and $\\gamma$-ray bands. Both the components are reproduced  with a double power-law electron distribution: the spectral index is $p_{low}$ from $\\gamma_{min}$ to $\\gamma_{break}$ and $p_{high}$ above $\\gamma_{break}$. The parameters of the two SSC components are reported in detail in Table 1.  We cannot exclude a second component due to an external seed photon field, for example mirrored by a putative broad line region, which could also account for the possible 1 day time lag. Nevertheless, the large amplitude of $\\gamma$-ray variability with respect to that of the optical favors a SSC explanation.    The luminosities observed in the optical and $\\gamma$-ray ranges are both $10^{48}$ erg s$^{-1}$, with a dissipated power in the jet rest-frame of $2 \\times 10^{43} (\\delta/15)^{-4}$ erg s$^{-1}$. This output is significantly large with respect to other BL Lacs, and we estimate a global power transported into the jet $L_{tot}>3 \\times 10^{45}$ erg s$^{-1}$. This may exceed the maximum power generated by a spinning black hole of mass $10^{9}$ M$_\\odot$ in most widely known models (see e.g. Cavaliere $\\&$ D'Elia 2002).  \\begin{figure}[!h] % \\centering \\includegraphics[angle=0,scale=0.08]{fig5.eps} \\caption[SED of S5~0716+714.]{ The SED of S5 0716+714, including GASP-WEBT optical data quasi-simultaneous with a AGILE-GRID gamma-ray observation in September (green dots). Historical data over the entire electromagnetic spectrum relative to a ground state of the source and EGRET non-simultaneous data are represented with blue dots. Red dots represent historical data simultaneous with a high X-ray state. \\label{0716:fig:spectrum}} \\end{figure}"
    },
    "0808/0808.3673_arXiv.txt": {
        "abstract": "{We report the $\\gamma$-ray activity from the Intermediate BL Lac S5 0716+714 during 2007 September--October observations by the AGILE satellite, coincident with a period of intense optical activity of the source monitored by GASP--WEBT.  AGILE observed the source with its two co-aligned imagers, the Gamma-Ray Imaging Detector (GRID) and the hard X-ray imager (Super-AGILE) sensitive in the energy range 30~MeV--50~GeV and 18--60 keV respectively, in two different periods: the first between 4 and 23 September 2007, the second between 24 October and 1 November 2007.  Over the period 7--12 September, AGILE detected $\\gamma$-ray emission from the source at a significance level of 9.6-$\\sigma$ with an average flux (E$>$100~MeV) of $(97 \\pm 15) \\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$, increasing by a factor of at least four within three days. No emission was detected by Super-AGILE in the energy range 18--60~keV, with a 3-$\\sigma$ upper limit of 10 mCrab in 335 ksec. The $\\gamma$-ray flux of S5 0716+714 detected by AGILE is the highest ever detected for this blazar and one of the most intense $\\gamma$-ray fluxes detected from a BL Lac object. The Spectral Energy Distribution (SED) of mid-September seems to be consistent with the synchrotron self-Compton (SSC) emission model, but only by including two SSC components with different variability.  In October 2007 AGILE repointed toward S5 0716+714 following an intense optical flare, measuring an average flux of $(47 \\pm 11) \\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$ at a significance level of 6.0-$\\sigma$.  The $\\gamma$-ray flux during both AGILE pointings appears to be highly variable on timescales of 1 day.}  \\FullConference{Workshop on Blazar Variability across the Electromagnetic Spectrum\\\\ April 22-25 2008\\\\ Palaiseau, France}  \\begin{document} ",
        "introduction": "The source S5 0716+714 was classified by Biermann (1981) as a BL Lac object, because of its featureless optical spectrum and high linear polarization. The optical continuum is so featureless that every attempt to determine the spectroscopic redshift of the source has failed; however, very recently through optical imaging of the underlying galaxy was estimated a redshift of z = 0.31 $\\pm$ 0.08 (Nilsson et al. 2008).  The source belongs to the Intermediate BL Lac class according to its Spectral Energy Distribution. In fact, observations by BeppoSAX (Tagliaferri et al. 2003) and XMM-$Newton$ (Foschini et al. 2006, Ferrero et al. 2007) provide evidence for a concave X-ray spectrum in the 0.1--10~keV band, a signature of the presence of both the steep tail of the synchrotron emission and the flat part of the Inverse Compton spectrum. The detection in the X-ray band of fast variability only in the soft X-ray component can be interpreted as the contemporary presence of a slowly variable Compton component and a fast and erratic variable tail of the synchrotron component.  In general, the variability of this blazar is strong in every band on both long and short intraday timescales. The optical and radio historical behaviour has been analyzed by Raiteri et al. (2003), while the EGRET telescope onboard $CGRO$ (Hartman et al. 1999) detected S5 0716+714 several times in the $\\gamma$-rays (Lin et al. 1995). The integrated flux above 100 MeV varied between (13 $\\pm$ 5) and (53 $\\pm$ 13) x 10$^{-8}$ photons cm$^{-2}$ s$^{-1}$.  We present the analysis of the AGILE data obtained during the S5 0716+714 observations in September--October 2007, in particular two flaring episodes: the first in mid-September, the other on 22--23 October 2007. Preliminary results were communicated in Giuliani et al. (2007) and a more detailed analysis is presented in Chen et al. (2008).  The strong $\\gamma$-ray flare detected by AGILE in mid-September triggered observations by the GLAST-AGILE Support Program (GASP) of the WEBT\\footnote[1]{\\texttt{http://www.oato.inaf.it/blazars/webt/; see e.g. Villata et al. (2007)}} (see Carosati et al., 2007). About one month later the GASP observed a new very bright phase of the source, triggering Swift as well as new AGILE observations. In the period from September to October 2007, S5 0716+714 showed intense activity with strong optical flaring episodes and a rare contemporaneous optical-radio outburst (Villata et al. 2008).  The results of a multiwavelength campaign on S5 0716+714 with simultaneous AGILE and Swift observations in October 2007 are discussed in Giommi et al. (2008). Throughout this paper the quoted uncertainties are given at the 1--$\\sigma$ level, unless otherwise stated. ",
        "conclusions": "To analyze the gamma-optical correlation we applied the Discrete Correlation Function (DCF; see Edelson $\\&$ Krolick 1988; Hufnagel and Bregman 1992; Peterson 2001) to the $\\gamma$-ray and $R$-band light curves. The $R$-band flux densities were averaged over 0.1 day bins to smooth the intranight variability. The DCF is a statistical method that was developed to analyze unevenly sampled data trains. The DCF displays a significant peak (DCF $\\sim$ 0.9) at a lag of -1 day (Fig. 2, right panel). Despite the large uncertainty due to poor $\\gamma$-ray sampling, this result suggests a possible delay of the $\\gamma$-ray flux variations with respect to the optical ones on the order of 1 day. Looking at Fig. 1, one can see that most of the DCF signal comes from the quasi-simultaneity of the $\\gamma$-ray and optical peaks of late October (JD $\\sim$ 2454396-397).   We notice that when the $\\gamma$-ray fluxes are $\\leq$ 120 $\\times$ 10$^{-8}$\\, photons cm$^{-2}$ s$^{-1}$, the corresponding optical flux densities are around 25--30 mJy. In contrast, the October $\\gamma$-ray peak reaching $\\sim$ 200 $\\times$ 10$^{-8}$\\, photons cm$^{-2}$ s$^{-1}$ has an optical counterpart of 40--45 mJy (see Fig. 1). This suggests that a strong optical event simultaneous to the $\\gamma$-ray flare was missed in September, since it occured at the beginning of the optical observing season as well as the start of the GASP activity.  The gamma variability seems to depend on the optical flux density changes roughly quadratically and this would favour a SSC interpretation, in which the emission at the synchrotron and IC peaks is produced by the same electron population, which self-scatters the synchrotron photons. In this case, the 1-day time lag in the high-frequency peak emission found from the DCF could be due to the light travel time of the synchrotron seed photons which scatter the energetic electrons.  \\begin{figure}[!t] % \\centering \\includegraphics[angle=0,scale=0.083]{SED_pos2.eps} \\caption{ The SED of S5 0716+714 including GASP-WEBT optical data quasi-simultaneous with AGILE-GRID $\\gamma$-ray observation in September 2007 (green dots). Historical data over the entire electromagnetic spectrum relative to a ground state of the source together with EGRET non-simultaneous data is represented with blue triangles. Red triangles represent historical data simultaneous with a high X-ray state.} \\label{figure3} \\end{figure}  The Spectral Energy Distribution with the AGILE and GASP-WEBT data of September 2007 is shown in Figure~\\ref{figure3} as green dots. The blue solid line shows a simple SSC model fitting simultaneous observations of a ground state (see Tagliaferri et al. 2003 and references therein) together with non-simultaneous EGRET data (empty blue triangles). Because the high state of mid-September 2007 cannot be fitted by a one-zone SSC component alone, we used a model with two SSC components. Without simultaneous X-ray data the spectrum is poorly constrained, then we show two models: one with a high hard X-ray state (red dashed line) and one with a low hard X-ray state (green solid line).  The first SSC component dominates in the optical and X-ray bands and it is reproduced  with a double power law electron distribution: the spectral index is $p_{low}=2$ from $\\gamma_{min}$ to $\\gamma_{break}$ and $p_{high}=4.5$ above $\\gamma_{break}$. For the high X-ray state model $\\gamma_{min}$ = 500, while the low X-ray state model has $\\gamma_{min}$ = 700; in both cases $\\gamma_{break}$ = 10$^3$. The density at the spectral break is $n_{e}=40\\,cm^{-3}$, the blob radius $R=2 \\times 10^{16}$ cm and the magnetic field $B=3$ Gauss.  The second SSC component contributes primarily to the gamma range of the SED and, to a lesser extent, to the optical emission. The electron distribution is a single power law with $p=4.5$, $\\gamma_{min}=4 \\times 10^{3}$ and $n_{e}=50\\,cm^{-3}$. The blob radius is $R=10^{16}$ cm and the magnetic field $B=1.3$ Gauss. Both blobs are moving with bulk Lorentz factor $\\Gamma=15$, at an angle of $3^{\\circ}$ with respect the line of sight.  We cannot exclude a second component due to an external seed photon field (e.g. mirrored by a putative broad line region) which could also account for the possible 1-day time lag, but the large amplitude of $\\gamma$-ray variability with respect to that of the optical one favours a SSC explanation."
    },
    "0808/0808.4115_arXiv.txt": {
        "abstract": "Although it is well recognized that gamma-ray burst (GRB) afterglows are obscured and reddened by dust in their host galaxies, the wavelength-dependence and quantity of dust extinction are still poorly known. Current studies on this mostly rely on fitting the afterglow spectral energy distributions (SEDs) with template extinction models. The inferred extinction (both quantity and wavelength-dependence) and dust-to-gas ratios are often in disagreement with that obtained from dust depletion and X-ray spectroscopy studies. We argue that this discrepancy could result from the prior assumption of a template extinction law. We propose an analytical formula to approximate the GRB host extinction law. With the template extinction laws self-contained, and the capability of revealing extinction laws differing from the conventional ones, it is shown that this is a powerful approach in modeling the afterglow SEDs to derive GRB host extinction. ",
        "introduction": "} In addition to the Galactic foreground extinction, GRBs and their afterglows are subject to extinction caused by the dust within their host galaxies. Evidence for this includes --- \\begin{itemize} \\item ``{\\it Dark bursts}'' -- an appreciable fraction of GRBs with X-ray and/or radio afterglows lack an optical afterglow (Jakobsson et al.\\ 2004).\\footnote{% Prior to the launch of {\\it Swift}, nearly $\\simali$60\\% of the X-ray afterglows reportedly had no optical counterparts. Despite rapid and deep searches in the {\\it Swift} era, it was found that $\\simali$1/3 GRBs with bright X-ray afterglows remain undetected at optical wavelengths (Fiore et al.\\ 2007, Schady et al.\\ 2007). } A natural explanation for dark bursts is that they lie behind significant obscuring dust columns in their host galaxies which effectively suppresses the optical light [although some dark bursts may be intrinsically faint or occur at high redshifts (say, $z\\gtsim 5$) where the Ly$\\alpha$ break has moved through the optical bands, leading to absorption of the optical light by the Ly$\\alpha$ forest]. Indeed, Schady et al.\\ (2007) found that the X-ray afterglows of GRBs not detected by UVOT were more affected by extinction than those of GRBs with detected UVOT counterparts. The recent detection of the near infrared (IR) afterglows of some GRBs (which would have been considered as ``dark bursts'' since their afterglows were not detected in any bluer bands) provides another piece of evidence for dust obscuration (e.g. see Jaunsen et al.\\ 2008, Tanvir et al.\\ 2008). \\item {\\it Reddening} -- some GRB afterglows with low redshifts appear very red, due to effects of extinction -- ultraviolet (UV)/visible light is extinguished more by dust than red light (e.g. see Klose et al.\\ 2000, Levan et al.\\ 2006). Dust reddening is also indicated by the significant deviation of the optical/near-IR spectral energy distributions (SEDs) of many afterglows from that expected from standard models. Also because of dust reddening, the Balmer line ratios in the spectra of some GRB host galaxies (e.g. see Djorgovski et al.\\ 1998), known as the {\\it Balmer decrement}, deviate from the expected ratios for the standard Case B recombination, which are fairly independent of physical conditions (Osterbrock \\& Ferland 2006). \\item {\\it Depletion} -- dust-forming heavy elements such as Si and Fe were found to be substantially depleted from the gas phase in some host galaxies (e.g. see Savaglio et al.\\ 2003). This indirectly shows the presence of dust in GRB host galaxies since the missing heavy elements must have been locked up in dust grains. \\item {\\it Connection between long GRBs and massive stars} -- there are multiple strong lines of evidence that long-duration ($\\simgt 2\\s$) GRBs are associated with the death of massive stars, occurring in regions of active star formation embedded in dense clouds of dust and gas (see Woosley \\& Bloom 2006). \\end{itemize}  A precise knowledge of the extinction (quantity, wavelength-dependence) and the nature (size, composition, and quantity) of the dust in GRB host galaxies is crucial for \\begin{itemize} \\item Correcting for the extinction of afterglows from X-ray to near-IR wavelengths to derive their intrinsic luminosities -- this is particularly important for studying the luminosity distribution of GRB afterglows and their intrinsic SEDs (e.g. see Kann et al.\\ 2008); \\item Constraining the nature of the GRB progenitors (i.e. collapsing massive stars or merging neutron stars) -- if long-duration GRBs are indeed linked to the collapse of massive stars, it is most likely that their optical and near-IR afterglows will suffer from significant attenuation in the star-forming molecular clouds heavily enshrouded by dust -- the birth place of these short-lived ($\\simali 10^{6}\\yrs$) massive stars; \\item Tracing the physical conditions of (and processes occurring in) the environments where GRBs occur which hold clues for understanding the mechanism for making a burst, e.g., a flat or gray extinction law for GRB host galaxies would imply a dense circumburst environment where dust undergoes coagulational-growth or a preferential destruction of small grains; and \\item Probing the interstellar medium (ISM) of high-redshift galaxies and the cosmic star formation history -- because of their intense luminosity which allows their detection at cosmological distances, GRBs are a powerful tool to study the star formation history up to very high redshifts; e.g., the dust and extinction properties of GRB hosts would help understand the nature of dark bursts and the dark burst fraction which would place important constraints on the fraction of obscured star formation in the universe (e.g. see Djorgovski et al. 2001, Ramirez-Ruiz et al.\\ 2002). \\end{itemize}  However, our current understanding of the dust extinction in GRB host galaxies is still very poor. Existing studies on this often draw conclusions in conflict with each other (see \\S\\ref{sec:status} for details). We argue that this could be caused by the prior adoption of a {\\it template} extinction law in fitting the observed GRB afterglow spectra to derive dust extinction (\\S\\ref{sec:status}). We propose in this work an alternative, robust method based on an analytical formula which can restore the widely adopted template extinction laws (\\S\\ref{sec:approach}). For illustration, we apply this approach to GRB\\,000301C and GRB\\,021004 (\\S\\ref{sec:test}). We demonstrate in \\S\\ref{sec:discussion} the uniqueness of the derived extinction laws. The robustness of this approach will be discussed in a separate paper (Liang \\& Li 2008a) in which the afterglow SEDs of $>$\\,50 GRBs of a wide range of properties are successfully modeled and for which the inferred extinction curves are diverse, with some differing substantially from any of the template extinction curves. ",
        "conclusions": "} We have also fitted the afterglow SEDs of GRB\\,000301C and GRB\\,021004 in terms of the MW, SMC, LMC, Calzetti, and ``linear'' % template extinction curves (see Table \\ref{tab:grbmod} and Figs.\\,\\ref{fig:GRB000301C},\\ref{fig:GRB021004}). Since for a given template extinction law the wavelength-dependence of the extinction $A_\\lambda/A_V$ is fixed, we are now left with only three parameters: $\\Fo$, $\\beta$, and $A_V$. The models based on the MW and LMC extinction laws could not fit the observed SEDs at all. This is because the 2175$\\Angstrom$ extinction feature which is prominent in the MW and LMC curves is absent in the SEDs of GRB\\,000301C and GRB\\,021004. In contrast, the SMC and ``linear'' models closely fit the afterglow SEDs of these two bursts, better than the Drude model proposed here as measured by $\\chi^2/N_{\\rm d.o.f.}$ (see Table \\ref{tab:grbmod}).  While the Drude model has three more parameters than the SMC and ``linear'' models, the quality of the fitting of the Drude model is even not as good as that of the SMC or ``linear'' model. Then, why do not we simply adopt the SMC or ``linear'' model?  First of all, we should note that there are no physical reasons for a prior assumption of a known extinction law, either that of the SMC, LMC, ``linear'' or MW: the composition and size distribution (and therefore the extinction law) of the dust in the dense circumburst clouds of GRB hosts with a wide range of metallicities and evolutionary stages are not expected to resemble that of the MW, LMC, or SMC (e.g. see Dwek 2005). In literature, a SMC-type extinction is often assumed for low-metallicity environments. However, there is no physical basis for this (except the lack of grain growth in these regions because of the lack of raw dust materials -- the SMC dust, on average, is substantially smaller than that of the Milky Way [see Weingartner \\& Draine 2001]). Moreover, it is known that the GRB hosts have a wide range of metallicities. Indeed, the reasons why the MW, LMC and SMC laws are often used for GRB afterglow SED modeling are mainly (1) little is known about the extinction laws of other galaxies, and (2) the Pei (1992) formula for the MW, LMC and SMC extinction laws is numerically convenient for computer implementation.  Second, although the SMC-type extinction is preferred in most of the present afterglow SED modeling studies, only the Drude approach is capable of reproducing the SEDs of those reddened by gray extinction or by non-conventional extinction. Indeed, it was shown that the afterglow SED of GRB\\,050904 at a redshift of $z$\\,$\\approx$\\,6.3 cannot be explained by dust reddening with any of the conventional (MW, SMC, Calzetti) extinction curves; instead, it can be well reproduced by invoking the extinction curve inferred for a distant quasar at $z$\\,=\\,6.2 (Maiolino et al.\\ 2004), suggesting that the properties of dust may evolve beyond $z$\\,=\\,6 (Stratta et al.\\ 2007).  Third, the Drude model would at least complement the models using template extinction curves, particularly for those bursts for which the Drude model gives a larger $\\chi^2/N_{\\rm d.o.f.}$ (but still fits the observed SEDs well). Given that the derived extinction $A_V$ and the intrinsic spectral slope $\\beta$ differ appreciably among different approaches (see Table \\ref{tab:grbmod}), the SMC model (and other models) should be used along side with the Drude model to gain insight into the ``true'' extinction and the ``true'' spectral slope.  We finally demonstrate the uniqueness of the extinction curve inferred from the Drude approach. To this end, we generate three sets of afterglow ``photometry data'' by reddening the intrinsic afterglow spectrum $F_\\nu\\,(\\mu {\\rm Jy})$\\,=\\,$5.2\\times 10^8 \\left(\\nu/{\\rm Hz}\\right)^{-0.5}$ of a burst at $z\\approx 2$ respectively with three template extinction laws: MW, SMC, and Calzetti, each with $A_V=0.5\\magni$. We then apply the Drude approach to these three sets of artificially-created GRB afterglow data. As shown in Figure\\,\\ref{fig:test}, we uniquely restore the MW, SMC, and Calzetti extinction laws: the inferred extinction curves are almost identical to that used to redden the intrinsic spectrum (the derived parameters [see Table\\,\\ref{tab:test}] are essentially the same as those tabulated in Table\\,\\ref{tab:extcurv})."
    },
    "0808/0808.3818_arXiv.txt": {
        "abstract": "We present new proper motions from the 10 m Keck telescopes for a puzzling population of massive, young stars located within 3\\farcs5 (0.14 pc) of the supermassive black hole at the Galactic Center. Our proper motion measurements have uncertainties of only 0.07 mas yr$^{-1}$ (3 km s$^{-1}$), which is $\\gtrsim$ 7 times better than previous proper motion measurements for these stars, and enables us to measure accelerations as low as 0.2 mas yr$^{-2}$ (7 km s$^{-1}$ yr$^{-1}$). Using these measurements, line-of-sight velocities from the literature, and 3D velocities for additional young stars in the central parsec, we constrain the true orbit of each individual star and directly test the hypothesis that the massive stars reside in two stellar disks as has been previously proposed. Analysis of the stellar orbits reveals only one of the previously proposed disks of young stars using a method that is capable of detecting disks containing at least 7 stars. The detected disk contains 50\\% of the young stars, is inclined by $\\sim115^\\circ$ from the plane of the sky, and is oriented at a position angle of $\\sim100^\\circ$ East of North. Additionally, the on-disk and off-disk populations have similar K-band luminosity functions and radial distributions that decrease at larger projected radii as $\\propto r^{-2}$. The disk has an out-of-the-disk velocity dispersion of 28 $\\pm$ 6 km s$^{-1}$, which corresponds to a half-opening angle of $7^\\circ \\pm 2^\\circ$, and several candidate disk members have eccentricities greater than 0.2. Our findings suggest that the young stars may have formed {\\it in situ} but in a more complex geometry than a simple, thin circular disk. ",
        "introduction": "\\label{sec:intro}  The center of our Galaxy harbors not only a supermassive black hole \\citep[Sgr A*, $M_\\bullet \\sim 4 \\times 10^6$ \\msun;][]{ eckart96,genzel96,ghez98pm,ghez00nat,ghez03spec,ghez05orbits, schodel02,schodel03,eisenhauer06}, but also a population of massive (10-120 \\msun), young ($\\lesssim$10-100 Myr) stars whose existence is a puzzle. The origin of such young stars has been difficult to explain since the gas densities observed today are orders of magnitude too low for a gas clump to overcome the extreme tidal forces and collapse to form stars \\citep[e.g.][for reviews]{sanders92,morris93,ghez05orbits,alexander05review}. And yet, within the central parsec of our Galaxy, nearly 100 stars have been classified as OB main-sequence stars, more luminous OB giants and supergiants, and post-main-sequence Wolf-Rayet stars \\citep{allen90,krabbe91,blum95heI,krabbe95,tamblyn96, najarro97,ghez03spec,paumard06}, with the more evolved massive stars having ages as young as 6$\\pm$2 Myr \\citep{paumard06}. Populations of young stars have also been observed in the nuclei of other galaxies, such as M31 \\citep{bender05}, suggesting that star formation near a supermassive black hole may be a common, but not understood, phenomenon in galaxy evolution. The close proximity of the black hole at the center of the Milky Way provides a unique laboratory for studying this ''paradox of youth'' \\citep[e.g.][]{ghez03spec,ghez05orbits,schodel03,eisenhauer06}.  Proposed resolutions to the paradox of youth can be grouped into several broad categories, including (1) rejuvenation of an older population such that older stars appear young, (2) dynamical migration from larger radii, and (3) {\\it in situ} formation. Rejuvenation scenarios include stripping \\citep{davies98,davies05} or tidal heating of the atmospheres of old stars \\citep{alexander03}, or combining multiple low mass stars via collisional mergers to form a higher-mass hot star akin to a ``blue straggler'' \\citep{lee96gcmergers,morris93,genzel03cusp}. Although these processes may be candidates for explaining the closest young stars within the central arcsecond, they cannot account for the OB giants, OB supergiants, and Wolf-Rayet stars that are located at larger radii (1\\arcsec-14\\arcsec), since the rate of collisions is too low to produce the observed total numbers. Thus, it appears that these massive young stars must have formed, or were deposited, in the central region within the last 4-8 Myr. Dynamical migration scenarios attempt to resolve the paradox of youth with the formation of a massive star cluster at larger distances from the black hole (3-30 pc). Such a cluster would spiral in due to dynamical friction and deposit stars at smaller radii where they are observed today \\citep{gerhard01}. However, for a cluster to reach the central parsec in only a few million years, it must be very massive and centrally concentrated \\citep{kim03,pzwart03irs16,mcmillan03,gurkan05}, and it may even require the existence of an intermediate-mass black hole (IMBH) as an anchor in the cluster core \\citep{hansen03,kim04}. {\\it In situ} star formation scenarios can resolve the paradox of youth if a massive, self-gravitating gas disk was once present around the black hole \\citep{levin03}. Such a disk would be sufficiently dense to overcome the strong tidal forces, and gravitational instabilities would then lead to fragmentation and the formation of stars, as has been suggested in the context of both the Galactic Center circumnuclear disk and AGN accretion disks in other galaxies \\citep[e.g.][]{kolykhalov80,shlosman89,morris96,sanders98,goodman03,nayakshinCuadra05}.  Insight into the origins of the massive, young stars may be obtained through observations of the spatial distribution and stellar dynamics of this population. Already, high-resolution infrared imaging and spectroscopy have shown that the young stars between 0\\farcs5 and 14\\arcsec (0.02-0.6 pc) exhibit coherent rotation \\citep{genzel00}. Analyses of the statistical properties of the three-dimensional velocity vectors for these stars suggest that they may reside in two disks. The first proposed disk has a clockwise sense of rotation, as projected onto the plane of the sky \\citep[][hereafter: clockwise-rotating or CW disk]{levin03}, while the second proposed disk is counter-clockwise-rotating \\citep[CCW][]{genzel03cusp} and is nearly perpendicular to the first. The proposed disks extend from $\\sim$0\\farcs8 to at least 7\\arcsec \\citep{paumard06}. Other velocity vector analyses show that there are possible co-moving groups or clusters of stars, including the IRS 13 cluster, which is proposed to lie within the putative CCW disk \\citep{maillard04irs13,schodel05}, and the IRS 16SW co-moving group, which are also consistent with the proposed CW disk \\citep{lu05irs16sw}. The two proposed disks are inferred to be oriented with an inclination and angle to the ascending node of [$i_{CW}$=127$^\\circ \\pm$ 2$^\\circ$, $\\Omega_{CW}$=99$^\\circ \\pm$ 2$^\\circ$] and [$i_{CCW}$=24$^\\circ \\pm$ 4$^\\circ$, $\\Omega_{CCW}$=167$^\\circ \\pm$ 7$^\\circ$] and to have a finite angular thickness of $\\Delta\\theta_{CW} \\sim 14^\\circ$ and $\\Delta\\theta_{CW} \\sim 19^\\circ$ where $\\Delta\\theta$ is the standard deviation of the orbital inclinations distributed normally about the disk plane \\citep{paumard06}. The thickness of the stellar disks has been attributed to thickening as a result of gravitational interactions between the two disks, which provides an estimate of the disk masses \\citep{nayakshin06thick}. The derived mass is smaller than the mass inferred from the number of observed young stars, assuming a Salpeter initial mass function (IMF); accordingly, \\citet{nayakshin06thick} suggest that the disks have a top-heavy mass function. Both {\\it in situ} gas disk and in-spiraling star cluster formation scenarios have been used to explain the kinematics of this young star population and to predict that the stars should lie in a common orbital plane. However, the presence of two stellar disks with similarly aged populations requires either two nearly concurrent gas disks or two infalling star clusters; and both of these scenarios are difficult to produce. Therefore, to understand the recent star formation history, it is critical to measure the orbital planes of individual stars in order to confirm the existence of the two stellar disks previously derived from a statistical analysis of velocity vectors alone.  The {\\it in situ} gas disk and inspiraling star cluster formation scenarios predict different structures and evolutions for the resulting stellar disk, particularly with respect to the eccentricities and radial distribution of stars within the disk. Early models of a self-gravitating gas disk around the supermassive black hole at the center of the Milky Way produce stars with a steep radial profile in the disk surface density, $\\Sigma \\propto r^{\\alpha}$, with $\\alpha \\sim -2$ \\citep{linPringle87,levin06}. These models typically result in stars on circular orbits as would be the case for the slow build up of a gas disk that is circularized before there is sufficient mass for gravitational instabilities to set in \\citep{milosav04,nayakshinCuadra05,levin06}. The stellar eccentricities of an initially circular disk can relax to higher eccentricities up to $e_{rms} = \\sqrt{<e^2>} \\sim $0.15 for a normal IMF or $e_{rms} \\sim $0.3 for a top-heavy IMF \\citep{alexander07imf,cuadra08}. More recent models have also shown that star formation can occur rapidly before circularization in an initially eccentric disk as might result from the infall of a single massive molecular cloud or a cloud-cloud collision \\citep{sanders98,nayakshin07sims,alexander08}. These eccentric self-gravitating accretion disk models typically produce a more top-heavy IMF than initially circular disks. On the other hand, an inspiraling star cluster would dissolve into a disk of stars with a flatter radial profile \\citep[$\\Sigma \\propto r^{-0.75}$;][]{berukoff06} whose orbital eccentricities would reflect the eccentricity of the cluster's orbit, which could be either circular or eccentric \\citep{pzwart03irs16,mcmillan03,kim03,kim04,gurkan05,berukoff06}. Previous measurements of the radial distribution of young stars yields a steep radial profile consistent with {\\it in situ} formation \\citep{paumard06}. Also, the eccentricities of the stars have previously been estimated from observations by assuming that the stars orbit in a disk; however, there are conflicting results claiming that the stars in the clockwise-rotating disk are on nearly circular orbits \\citep{paumard06} or on eccentric orbits \\citep{beloborodov06}. Determining the radial profile and stellar eccentricities of stars in a disk may provide observational constraints on the origin of the young stars.  We present an improved proper motion study that yields an order of magnitude more precise proper motions and the first measurement of accelerations in the plane of the sky for stars outside the central arcsecond. By combining the stellar positions, proper motions, radial velocities, and accelerations, we estimate stellar orbital parameters and test whether the young stars reside on one or two stellar disks in a more direct manner than previous methods using only velocity information. This provides a {\\it direct} test of the existence, membership, and properties of these disks. The observations are described in \\S\\ref{sec:obs} and the astrometric analysis procedure and results are detailed in \\S\\ref{sec:astrometry}. Orbit analysis and results are presented in \\S\\ref{sec:orbitAnalysis} and \\S\\ref{sec:orbitResults} and a discussion of the implications for the origin of the massive, young stars at the Galactic Center is presented in \\S\\ref{sec:discussion}. ",
        "conclusions": "In summary, the advent of laser guide star adaptive optics has allowed us to retroactively improve our 11 year astrometric data set used for monitoring stars orbiting our Galactic Center. This has increased our proper motion precision, with resulting uncertainties of $\\sim$3 km s$^{-1}$, and allowed us, for the first time, to make measurements of and place limits on accelerations for stars outside the central arcsecond out to a radius of 3\\farcs5, with typical 3$\\sigma$ acceleration limits of -0.19 mas yr$^{-2}$. By combining our improved stellar positions and proper motions with radial velocity information from the literature, we compute orbits for individual young stars proposed to lie in stellar disks orbiting the supermassive black hole. The orbits for the young stars confirm only a single disk of young stars at a high inclination rotating in a clockwise sense and there is no statistically significant evidence for a second disk. Stars within the well-defined, clockwise disk have an out-of-the-disk velocity dispersion of 28 $\\pm$ 6 km s$^{-1}$ and several stars have high eccentricities. These disk properties suggest that star formation may have occurred in a single event, rather than the two events previously needed to explain two stellar disks; however, there are open questions as to how $\\sim$50\\% of all young stars can be perturbed out of the disk plane and whether the apparent compact cluster, IRS 13, which is not part of the stellar disk, requires a separate star formation or dynamical event. Future directions include (1) obtaining new LGSAO data sets with improved astrometry to measure accelerations for the young stars at all radii and (2) identifying new young stars within the central parsec in order to better constrain the orbital properties of these stars and to study in detail the distribution of eccentricities and semi-major axes for stars both in and out of the disk."
    },
    "0808/0808.2733_arXiv.txt": {
        "abstract": "{A large number of AGN have been monitored for nearly 30 years at 22, 37 and 87 GHz in Mets\\\"ahovi Radio Observatory. These data were combined with lower frequency 4.8, 8.0 and 14.5 GHz data from the University of Michigan Radio Astronomy Observatory, higher frequency data at 90 and 230 GHz from SEST, and supplementary higher frequency data from the literature to study the long-term variability of a large sample of AGN. Both the characteristics of individual flares from visual inspection and statistically-determined variability timescales as a function of frequency and optical class type were determined. Based on past behaviour, predictions of sources expected to exhibit large flares in 2008--2009 appropriate for study by GLAST and other instruments are made. The need for long-term data for properly understanding source behaviour is emphasised.}   \\FullConference{Workshop on Blazar Variability across the Electromagnetic Spectrum\\\\ April 22-25, 2008\\\\ Palaiseau, France}  \\begin{document} ",
        "introduction": "Active galactic nuclei (AGN) are variable across the whole electromagnetic spectrum. The radio regime is of special interest because we can directly observe the synchrotron radiation from the jet. The total flux density variations and flares seen in the flux curves are usually explained with shock-in-jet models where a disturbance creates a shock moving down in the jet, and as the shock develops we see a flare evolving from higher submm- and mm wavelengths towards lower radio frequencies \\cite{marscher85, hughes85}.  We have used a sample of 90 AGN to study their long-term variability timescales \\cite{hovatta07, hovatta08b} and flare characteristics \\cite{hovatta08a, nieppola08}. Our extensive database enables us to study the correspondence between the shock model and the observations, and also the statistical differences between the different AGN types (27 high polarisation quasars (HPQs), 33 low polarisation quasars (LPQs), 25 BL Lacertae objects (BLOs), and 5 radio galaxies (GALs)). Our sample includes bright sources that have flux density at least 1 Jy in the active state. ",
        "conclusions": "We studied the long-term radio variability of a sample of 90 sources using statistical timescale analysis methods and visual inspection of the flare parameters. Our main results are the following:  \\begin{itemize}  \\item{Fourier-based methods the DCF and the periodogram give very similar results as wavelets, but wavelets should be used when quasi-periodities are studied because they give information on the locality of the timescale. With wavelets it is possible to see if a timescale is long-lasting or just a short transient phenomenon in the flux curve. DCF or periodograms can then be used to verify the timescale more accurately.}  \\item{Variability behaviour is complex and no clear periodicities could be found at radio frequencies. Episodes of quasi-periodic behaviour are common, and therefore false periodicities may be found if the temporal coverage is inadequate.}  \\item{Flares are seen, on average, every 4 years in all the source types at 37\\,GHz but when intrinsic redshift-corrected timescales are studied, the quasars have shorter timescales of 2 years compared to the 3-4 years of BLOs. This could indicate that shocks are produced less frequently in BLOs than in quasars.}  \\item{Median duration of a flare is 2.5 years at 22 and 37\\,GHz, but the range in durations is between 0.3 and 13.2 years. When comparing the duration with intrinsic redshift- and Doppler-corrected peak luminosities, we found that the energy release in a flare does not increase with the duration of the flare.}  \\item{Flares adhere quite well to the predictions of the shock model but the scatter in the data, due to poor sampling and complicated structure of the flares, is still large.}  \\item{By combining the median duration of flares, 2.5 years, with the average time between the flares, 4 years, we see that multifrequency campaigns should last for 5-7 years in order to catch the source in both its highest and lowest activity states.}  \\item{Long-term monitoring is essential in understanding the true behaviour of these sources at radio frequencies.} \\end{itemize}"
    },
    "0808/0808.2505_arXiv.txt": {
        "abstract": "We present a comprehensive abundance analysis of 27 heavy elements in bright giant stars of the globular clusters M4 and M5 based on high resolution, high signal-to-noise ratio spectra obtained with the Magellan Clay Telescope. We confirm and expand upon previous results for these clusters by showing that (1) all elements heavier than, and including, Si have constant abundances within each cluster, (2) the elements from Ca to Ni have indistinguishable compositions in M4 and M5, (3) Si, Cu, Zn, and  all $s$-process elements are approximately 0.3 dex overabundant in M4 relative to M5, and (4) the $r$-process elements Sm, Eu, Gd, and Th are slightly overabundant in M5 relative to M4. The cluster-to-cluster abundance differences for Cu and Zn are intriguing, especially in light of their uncertain nucleosynthetic origins. We confirm that stars other than Type Ia supernovae must produce significant amounts of Cu and Zn at or below the clusters' metallicities. If intermediate-mass AGB stars or massive stars are responsible for the Cu and Zn enhancements in M4, the similar [Rb/Zr] ratios and (preliminary) Mg isotope ratios in both clusters may be problematic for either scenario. For the elements from Ba to Hf, we assume that the $s$- and $r$-process contributions are scaled versions of the solar $s$- and $r$-process abundances. We quantify the relative fractions of $s$- and $r$-process material for each cluster and show that they provide an excellent fit to the observed abundances. ",
        "introduction": "\\label{sec:intro}  Globular clusters continue to play a vital role in testing many aspects of stellar evolution and stellar nucleosynthesis. Observational and theoretical studies of globular clusters have focused heavily upon (1) the star-to-star light element abundance variations \\citep{cottrell81,langer95,gratton04}, (2) the cluster-to-cluster variation in the color distribution of horizontal branch stars, the so-called ``2nd parameter effect'' \\citep{sandage67,lee94,carretta06}, and (3) the multiple populations as inferred from large spreads in metallicity and/or detailed structure in color-magnitude diagrams \\citep{butler78,norris95b,bekki06}.  Abundance measurements of the $s$-process and $r$-process elements offer great insight into stellar nucleosynthesis and globular cluster chemical evolution. Aside from M15, a metal-poor cluster that displays a scaled-solar $r$-process abundance distribution \\citep{sneden97,sneden00b,otsuki06}, in general only a handful of $s$-process elements (e.g., Y, Zr, Ba, La) and the $r$-process element Eu have been measured in globular clusters.  The globular clusters M4 and M5 are particularly well suited for refining our understanding of stellar evolution and stellar nucleosynthesis, especially for the neutron-capture elements. \\citet{M4,M5} showed that these clusters have essentially identical metallicities, [Fe/H] = $-$1.2, based on high resolution spectra of large samples, 36 stars in each cluster. They also showed that the abundance similarities for these clusters extend to numerous $\\alpha$- and Fe-peak elements as well as the $r$-process element Eu. However, the $s$-process elements revealed striking abundance differences between these two clusters. Specifically, the heavy $s$-process elements Ba and La are overabundant in M4 relative to M5 \\citep{M4,M5}. With the notable exception of $\\omega$ Cen, M4 may be uniquely enriched in $s$-process elements among the Galactic globular clusters \\citep{pritzl05}.  \\citet{rbpbm4m5} recently extended the analysis of neutron-capture elements in these two clusters to Rb and Pb, $s$-process elements which may be overproduced in metal-poor asymptotic giant branch (AGB) stars \\citep{busso99,travaglio01}. Whereas Pb production is dominated by 2 to 4$M_{\\odot}$ low-metallicity AGB stars \\citep{travaglio01}, the intermediate-mass 4 to 8$M_{\\odot}$ AGB stars are predicted to dominate Rb production \\citep{vanraai08}. Not surprisingly, M4 again had higher abundance ratios [Rb/Fe] and [Pb/Fe] than M5. However, the abundance ratios [Rb/X] for X = Y, Zr, and La were very similar in the two clusters indicating that the nature of the $s$-process products is very similar for both clusters but that M4 formed from gas with a higher concentration of these products. A comprehensive study of $s$-process elements in these two clusters promises to provide a novel observational study of the $s$-process at low metallicities as well as valuable clues to the chemical evolution diversity of globular clusters. In this paper we present such an analysis, focusing upon a suite of $\\alpha$-, Fe-peak, $s$-process, and $r$-process elements. ",
        "conclusions": "\\label{sec:summary}  In this paper we present abundance ratios [X/Fe] for a large number of $\\alpha$-, Fe-peak, $s$-process, and $r$-process elements for 12 bright giants in the globular cluster M4 and 2 bright giants of the globular cluster M5. This comprehensive abundance analysis is only possible due to the large wavelength coverage, high resolution, and very high S/N spectra. For all elements in this study, we find no evidence for star-to-star abundance variations in either cluster. We confirm and extend upon previous results for these clusters by showing that (1) for the elements from Ca to Ni, M4 and M5 have identical abundance ratios, (2) M4 shows overabundances by roughly 0.3 dex for Si, Cu, Zn, and all $s$-process elements relative to M5, and (3) for the $r$-process elements, M5 may have slightly higher abundances than M4 by 0.1 dex.  We also measure Mg isotope ratios and find that the ratios are solar in both clusters, with no sign of any star-to-star variation within each cluster. The ratios $^{25}$Mg/$^{24}$Mg and $^{26}$Mg/$^{24}$Mg exceed values found in field halo stars at the same metallicity, e.g., Gmb 1830, which implies differences in the clouds from which globular clusters and field halo stars formed. However, we regard these ratios as preliminary since the spectral resolution was insufficient to accurately distinguish $^{25}$Mg from $^{26}$Mg.  There is no clear explanation for the M4-M5 Si abundance differences since the abundances of Ca should, but do not, follow the behavior of Si. For the elements from Ba to Hf, we find that the mean abundances in M4 and M5 are well explained by scaled versions of the solar $s$- and $r$-process abundances, albeit with different mixes of $s$- and $r$-process material for each cluster. Therefore, no new $s$-process site is required to explain the M4-M5 abundance differences for the elements from Ba to Hf. However, although the Th abundances lie above these predictions, the ratio [Th/Eu] is identical in both clusters indicating that the universality of the $r$-process extends to Th in these clusters and that no differential decay of Th has occurred, i.e., the clusters have identical ages. The Pb abundances lie below the predictions, by different amounts for each cluster. Therefore, the sources of the $s$-process may differ between M4 and M5, at least regarding the production of Pb via the strong component.  The abundance differences between M4 and M5 for Cu and Zn are particularly intriguing given that their nucleosynthetic origins continue to be debated. The $s$-process elements, produced in AGB stars (and massive stars), share a similar abundance behavior to Cu and Zn in M4 and M5. Updated, but preliminary, yields from AGB models indicate that small amounts of Zn may be produced only in the most massive AGB stars. These models also predict that the most massive AGB stars may produce Cu in contrast to our current understanding of Cu production. Massive AGB stars are expected to produce large amounts of the neutron-rich Mg isotopes, and observations of high $^{26}$Mg/$^{24}$Mg ratios in field and cluster stars confirm the AGB yields. However, preliminary measurements show no difference in the Mg isotope ratios between these two clusters which constrains the contribution of intermediate-mass AGB stars to the Cu and Zn enhancements in M4. Si, which is produced in massive stars, shows a similar abundance behavior to Cu and Zn in M4 and M5. Cu and Zn may be produced in massive stars via the weak $s$-process. While the abundance ratios [Rb/Sr], [Rb/Y], and [Rb/Zr] are predicted to increase via the weak $s$-process, our measurements do not reveal any cluster-to-cluster variations in these abundance ratios which suggests that either massive stars are not responsible for the Cu and Zn differences or that metal-poor massive stars do not alter the [Rb/Zr] ratio. Of great interest would be detailed chemical evolution modeling of these two clusters to gain insight into the origin of the Cu and Zn abundance differences and therefore their nucleosynthesis production sites."
    },
    "0808/0808.3579_arXiv.txt": {
        "abstract": "{ Hydromagnetic stresses in accretion discs have been the subject of intense theoretical research over the past one and a half decades. Most of the disc simulations have assumed a small initial magnetic field and studied the turbulence that arises from the magnetorotational instability. However, gaseous discs in galactic nuclei and in some binary systems are likely to have significant initial magnetisation. Motivated by this, we performed ideal magnetohydrodynamic simulations of strongly magnetised, vertically stratified discs in a Keplerian potential. Our initial equilibrium configuration, which has an azimuthal magnetic field in equipartion with thermal pressure, is unstable to the Parker instability. This leads to the expelling of magnetic field arcs, anchored in the midplane of the disc, to around five scale heights from the midplane. Transition to turbulence happens primarily through magnetorotational instability in the resulting vertical fields, although magnetorotational shear instability in the unperturbed azimuthal field plays a significant role as well, especially in the midplane where buoyancy is weak. High magnetic and hydrodynamical stresses arise, yielding an effective $\\alpha$-value of around 0.1 in our highest resolution run. Azimuthal magnetic field expelled by magnetic buoyancy from the disc is continuously replenished by the stretching of a radial field created as gas parcels slide in the linear gravity field along inclined magnetic field lines. This dynamo process, where the bending of field lines by the Parker instability leads to re-creation of the azimuthal field, implies that highly magnetised discs are astrophysically viable and that they have high accretion rates. ",
        "introduction": "Since the seminal work of \\cite{BalbusHawley1991} it has been widely recognised that hydromagnetic stresses play a central role in the dynamics of Keplerian gaseous discs. Much of the subsequent theoretical work has been devoted to the study of the magnetic dynamo driven by a combination of magnetorotational instability (MRI), rotation, and Keplerian shear \\citep[][see \\citealp{BrandenburgSubramanian2005}, for an excellent review of astrophysical dynamo theory]{Brandenburg+etal1995,Hawley+etal1996}. Typically, numerical simulations of the MRI-driven dynamo begin with an initial zero net flux magnetic field with an associated pressure which is a small fraction of the thermal pressure. The idea is that the MRI-driven turbulence should increase the characteristic coherence length of the magnetic field, and grow its mean strength to a significant fraction of the equipartition value. The simulations, however, have had at best a mixed success in explaining high and persistent inflow rates observed in astrophysical accretion discs. Firstly, a number of recent numerical experiments have shown that the effectiveness of the dynamo depends in a critical way on the value magnetic Prandtl number, defined as the ratio of the collisional kinematic viscosity to the magnetic diffusivity \\citep[][following suggestive non-disc simulations of \\citealp{Schekochihin+etal2005} and a careful study of previously published MRI-turbulence results by \\citealp{Pessah+etal2007}]{LesurLongaretti2007,Fromang+etal2007}. The upshot of this work is that when the Prandtl number is significantly less than one, as is expected in most accretion discs \\citep{BalbusHenri2008}, the dynamo seems to fail\\footnote{Also the magnitude of the turbulent magnetic stresses decrease as the grid resolution increases (decreasing the numerical diffusivity accordingly), with no convergence in sight \\citep{FromangPapaloizou2007}. In may well be, however, that this problem is related to an insufficient size of the simulation domain, and that bigger, stratified shearing boxes will show better convergence \\citep{RegevUmurhan2008}.}. Secondly, in the zero net flux simulations where all parameters are chosen so that the dynamo works, the measured value of the Shakura-Sunyaev parameter $\\alpha$ is typically of order $10^{-3}$ and at most $10^{-2}$ \\citep[e.g.][]{Brandenburg+etal1995,Sano+etal2004,JohansenKlahr2005,FromangNelson2006}. This is $1$--$2$ orders of magnitude smaller than what is required to explain the high accretion rates in dwarf novae systems \\citep{King+etal2007}.  The weak initial magnetic fields assumed in almost all disc-MRI simulations may be unrealistically small for many astrophysical discs. We give here two examples:\\newline \\hskip .1in 1. Extended accretion discs around central supermassive black holes \\citep[such as the ones traced by megamasers in nearby galaxies; see e.g.][]{Greenhill2007,Vlemmings+etal2007} are fed by molecular material in the interstellar medium. Molecular clouds are known to have large scale superthermal magnetic fields, and thus the initial magnetic fields in AGN discs are likely to be comparable to or larger than the thermal equipartition values. The $\\sim$2 pc molecular circumnuclear disc in our galactic centre is permeated by large scale equipartition magnetic fields \\citep{WardleKonigl1990,Hildebrand+etal1990} which are certain to play an important dynamical role in its subsequent evolution.\\newline \\hskip .1in 2. Accretion discs in close binary systems, which are fed by a Roche-lobe overflow from a tidally distorted low mass star, may be initially magnetised. The magnetisation $1/\\beta$ of the gas as it departs from the donor star can be estimated as \\begin{equation}\\label{binary1} {1\\over \\beta}\\sim {B_\\star^2 \\over 8\\pi \\rho c_s^2}, \\end{equation} where $B_\\star$ (measured in Gauss) is the stellar magnetic field at the surface, $c_s$ is the speed of sound at the stellar surface, and $\\rho$ is the density of gas as it becomes detached from the star. The density $\\rho$ can be estimated as \\begin{equation}\\label{binary2} \\rho\\sim {\\dot{M}\\over 4\\pi R_\\star^2 c_s}, \\end{equation} where $\\dot{M}$ is the average accretion rate of the companion, and $R_\\star$ is the radius of the star. We get \\begin{equation} {1\\over \\beta}\\sim 0.3 \\left({\\dot{M}\\over 10^{-9}M_{\\odot}/\\hbox{yr}}\\right)^{-1} \\left({c_s\\over 3\\hbox{km/s}}\\right)^{-1}\\left({B_\\star \\over 1\\hbox{G}}\\right)^2 \\left({R_\\star\\over 10^{11}\\hbox{cm}}\\right)^2. \\label{binary3} \\end{equation} The streaming and shearing motion after the gas detaches from the star can further amplify the magnetic field. Clearly, there is a realistic parameter range where the initial magnetisation is high.  Motivated by these astrophysical considerations, in this paper we perform numerical experiments on gas discs which contain initially strong magnetic fields, with magnetic pressure comparable to that of the  gas. Specifically, we initialise the azimuthal field so that it is subject to the Parker instability \\citep[PI,][]{Parker1966} in the vertical stratification.  In terms of initial conditions our physical set up is close to \\cite*{Machida+etal2000}. However, \\cite{Machida+etal2000} focused on the formation and evolution of a disc corona, and the spirit or their numerical experiment was different from ours. They have simulated the whole circular disc, and have introduced reflecting boundary conditions at the midplane of the disc, which suppresses field anchoring in the midplane. By contrast, our purpose is to understand the long-term dynamics of the disc, possible field confinement, and the dynamo processes related to strong magnetic fields. Therefore, our shearing-box simulations focus on a small part of the disc and study it with high numerical resolution. We simulate the fluid both below and above the midplane, and thus have no artificial boundary conditions at the midplane of the disc.  We find that both the short term and, more importantly, the long term behaviour of initially strongly magnetised discs is radically different from that of their weakly magnetised counterparts. We observe the following three-step dynamics: (a) the Parker instability expels azimuthal field in huge arcs, creating vertical field which becomes the seed for a strong magnetorotational instability, (b) matter sliding down inclined field lines stretches the azimuthal magnetic field and creates a vertically dependent large scale mean radial field, and (c) the Keplerian shear recreates azimuthal field from the stretching of the radial field. The latter step closes the dynamo cycle, in much the same way as was sketched by \\cite{ToutPringle1992} a decade and a half ago, although non-axisymmetric magnetorotational instability in the azimuthal field also plays an important role in creating accretion stresses in our simulations \\citep{BalbusHawley1992,FoglizzoTagger1995,TerquemPapaloizou1996}. The azimuthal field remains strongly concentrated towards the disc midplane; this is in contrast with the simulations of Parker instability which do not include strong Keplerian shear in the radial direction \\citep{Kim+etal1998}. We show that the dynamo is robust and stable over at least tens of orbital periods, and that the accretion torque {\\it increases} if we use a finer grid. We observe $\\alpha$-values as high as $0.1$ in our highest resolution run.   The plan of the paper is as follows. In \\S\\ref{s:numerics} we detail the mechanics of our numerical experiments, and in the following section (\\S\\ref{s:noshear}) we test our computer code by comparing the results of Parker instability simulations without Keplerian shear to the extensive literature that exists on the PI under rigid rotation. In \\S\\ref{s:shear} we perform simulations with the Keplerian shear included, and present our main results. We devote the next section (\\S\\ref{s:confinement}) to analysing the confinement of azimuthal flux to the disc by a dynamo process. In \\S\\ref{s:conclusions} we conclude with the discussion of the astrophysical implications and possible future improvements of our work. ",
        "conclusions": "\\label{s:conclusions}  In this paper we consider the evolution of strongly magnetised Keplerian accretion discs. Our numerical experiments show that the hydromagnetic state of the gas flow is very different from what is seen in zero net flux simulations. The Parker instability leads to huge magnetic arcs rising several scale heights from the disc midplane, and the magnetorotational instability in turn feeds off the vertical fields and creates a highly turbulent flow, an interaction that was predicted analytically by \\cite{ToutPringle1992}. Although the flow is stochastic and time fluctuating, we have identified an underlying dynamo process that couples the vertically dependent mean radial and azimuthal magnetic field components. As gas slides down inclined field lines, it obtains a helical motion due to Coriolis forces, and thus the azimuthal field lines are twisted in such a way as to create a mean electromotoric force in the direction of the unperturbed field line -- a configuration prone to create radial field. In turn the large scale radial field regenerates the azimuthal field through Keplerian shear. Although Parker instability dominates the linear growth phase, we have found evidence for magnerotational instability in the azimuthal field as well. In the midplane of the disc, where the buoyancy is weak, azimuthal MRI drives the initial evolution towards turbulence \\citep{FoglizzoTagger1994,FoglizzoTagger1995,TerquemPapaloizou1996}. These two related instabilities, magnetorotational instability in the vertical fields created by the Parker instability and magnetorotational swing instability in the azimuthal fields, both rely on azimuthal flux confinement and can coexist in the linear as well as in the non-linear state of transmagnetic accretion discs.  Such a path to accretion, based on the interaction of Parker and magnetorotational instabilities, has at least two appealing traits. First of all that the vertical fields that feed the magnetorotational instability are created in a transparent way by the Parker instability. Zero magnetic flux models must most likely rely on a small scale dynamo in order to create vertical fields, and there is mounting evidence that such a dynamo would not operate in the bulk part of accretion discs where the magnetic Prandtl number is much lower than unity \\citep{Schekochihin+etal2005,Fromang+etal2007}. The second appealing result of our model is that the Maxwell and Reynolds stresses are significant ($\\alpha\\approx0.1$). Such high accretion stresses could solve the problem that observed accretion rates are often one or two orders of magnitude higher than the accretion rates obtained in zero net flux MRI simulations \\citep{King+etal2007}.  The regeneration of azimuthal field by the shearing of an appropriate radial field was seen in all our simulations that included Keplerian shear. As magnetohydrostatic equilibrium is compromised by the Parker instability, gas streams down along inclined field lines. Coriolis force diverts the gas to the right, and a radial magnetic field is created as the azimuthal field is subjected to shear-regions typically the size of the Parker instability. Eventually magnetic reconnection leads to a coherent large scale radial magnetic field. This dynamo was predicted by \\cite{Parker1992} and subsequently observed in the rigid rotation simulations of \\cite{Hanasz+etal2002}. To our knowledge we are the first to point out the relevance of Parker's fast galactic dynamo to accretion discs and how it closes the accretion loop by replenishing the azimuthal field that is lost by magnetic buoyancy.  The fine-tuned initial conditions with a purely azimuthal magnetic field and a constant ratio of magnetic to thermal pressure may be questioned. However our experiments with a combined azimuthal and vertical field shows that the Parker instability is robust even if the azimuthal field coexists with a moderately strong net vertical field, and that the additional vertical field component may indeed increase the accretion rate further.  Our results may also be relevant for star formation in the galactic centre. Although there is currently no coherent accretion disc structure, the population of young, massive stars in a disc-like structure close to the galactic centre points towards the brief existence of an accretion disc some million years ago \\citep{LevinBeloborodov2003,MilosavljevicLoeb2004,Nayakshin+etal2007,Alexander+etal2008}. The disc was likely to be initially strongly magnetised, as indicated by the current high magnetisation of the circumnuclear molecular ring. Hence the type of MHD processes studied in this paper may be of central significance for the disc dynamics. The presented simulations of transmagnetic ($\\beta\\sim1$) discs argues that discs dominated by magnetic pressure $\\beta \\ll 1$ are astrophysically viable. The existence of such discs was conjectured by \\cite{Pariev+etal2003} and their limitations and  observational consequences were explored by \\cite{BegelmanPringle2007}. A number of problems in accretion disc theory are alleviated by the presence of super-equipartition magnetic fields, among them is the long-standing issue of self-gravity and fragmentation of AGN discs \\citep{Goodman2003}.  Future research into strongly magnetised accretion discs should also focus on the self-consistent modelling of the magnetisation of the material that feeds accretion discs in various environments and on the evolution of magnetised disc coronae, in light of effects such as reconnection, shearing of foot points, Ohmic heating and radiative cooling that take place there."
    },
    "0808/0808.1056_arXiv.txt": {
        "abstract": "A rudimentary calculation is employed to evaluate the possible effects of $\\beta$-decays of excited state nuclei on the astrophysical r-process. Single particle levels calculated with the FRDM are adapted to the calculation of $\\beta$-decay rates of these excited state nuclei.  Quantum numbers are determined based on proximity to Nilson model levels. The resulting rates are used in an r-process network calculation in which a supernova hot-bubble model is coupled to an extensive network calculation including all nuclei between the valley of stability and the neutron drip line and with masses 1$\\le$A$\\le$283. $\\beta$-decay rates are included as functional forms of the environmental temperature. While the decay rate model used is simple and phenomenological, it is consistent across all 3700 nuclei involved in the r-process network calculation.  This represents an approximate first estimate to gauge the possible effects of excited-state $\\beta$-decays on r-process freezeout abundances. ",
        "introduction": "\\label{Introduction} The r-process is responsible for the synthesis of roughly half of all nuclei heavier than A$\\sim$70 and all of the actinides \\cite{cowan90,wallerstein97}. The solar system r-process abundances act as the canonical constraint to r-process theories as well as the prime indicator of the success of r-process models. Several r-process sites have been proposed; the hot-bubble region of a type II supernova (SNII) has been modeled fairly successfully. The composition of the environment in which the r-process occurs might be expected to have a profound effect on the final abundance distribution. Observations indicate that the r-process site is primary \\cite{sneden98} and further evidence may suggest that the r-process is also unique \\cite{cowan99}; it may occur in a single site or event.  The uniqueness of the r-process site, however, remains a subject of study \\cite{qian98}.  Nuclear properties also constrain the r-process, and the purpose of this work is the examination of one particular characteristic - $\\beta$-decay - as it relates to the r-process.  The $\\beta$-decay inputs, and other nuclear physics inputs, have been shown \\cite{woosley94,meyer95} to have important effects on the success or failure of r-process models.  This is somewhat unfortunate, as properties of only a few nuclei on the neutron closed shells closer to stability have been experimentally determined, while data for the rest are relegated to calculation \\cite{cowan90}.  Of paramount importance is the determination of nuclear masses and $\\beta$-decay rates. Nuclear mass formulae based on the microscopic properties of nuclei are slowly replacing the empirical droplet models, and these can change resulting reaction rates by factors as large as 10$^8$ \\cite{goriely96}.  As well, the r-process path is affected by the choice of mass formula, since the path roughly follows a line of constant S$_n$ \\cite{cowan90}.  A successful r-process calculation must predict an abundance peak for nuclei in the A$\\sim$195 region - a difficulty over much of the r-process parameter space. However, as discussed below, $\\beta$-decays from excited state nuclei may help to mitigate this difficulty.  They do so by allowing the r-process to proceed at a faster rate, thereby enhancing the abundances at higher masses.  For the purposes of this study, the most recent semi-gross theory of $\\beta$-decay \\cite{nakata97} has been adapted to neutron-rich nuclei relevant to the r-process.  The ability of this model to determine decay properties of an extremely wide range of nuclei with reasonable accuracy and speed makes it ideal for this preliminary calculation.  In particular, the semi-gross theory has good agreement for very neutron-rich nuclei \\cite{ichikawa05,asai99}.  It has also been used to improve the accuracy of decay rates for astrophysical calculations by incorporating first-forbidden transition strengths \\cite{moller03}. In its original form, the gross theory of $\\beta$-decay assumed that the energy states of a nucleus consist of a smoothed distribution with transition strengths that peak at or near the energy of the isobaric analog state \\cite{takahashi72,takahashi69}.  Subsequent evolutions of the gross theory incorporated strength functions allowing for transitions of higher forbiddenness \\cite {takahashi71}, as well as improvements over the original theory to include odd-odd effects \\cite{nakata95}, sum rules \\cite{tachibana90}, even-odd mass differences \\cite{koyama70}, and improvements on the strength functions \\cite {kondoh85}.  While the precision of the Fermi gas model used is consistent with that of the gross theory itself, more accurate models might be used.  Recently, shell effects in the parent nucleons have been taken into account by using an energy distribution for single-particle states \\cite{nakata97,kondoh76}; this is denoted as the ``semi-gross theory.'' In these models, the energy distribution of the daughter states is still assumed to be smooth. Thus, the transition strength functions depend on the quantum numbers only of the initial parent states and are independent of the energy of the parent state; transition types are based on a statistical weight for a particular parent state to make a particular type of transition.  The advantage of using single-particle strength functions is that quantum numbers can be assigned to the states easily, lending a better notion of the actual strength of each type of transition involved.  Since decay rates of nearly all of the nuclei along the r-process path have yet to be studied in a laboratory, r-process calculations rely heavily on calculated decay rates. Further, the temperature of the r-process environment (~10$^9$K) necessitates accounting for nuclei in excited states, especially given the expected high level density of these far-from-stability nuclei. Some of the effects that might be expected if one considers excited state nuclei in r-process simulations include increased (n,$\\gamma$) and ($\\gamma$,n) rates, which might shift the r-process path, but would tend to counteract each other as the r-process is generally presumed to proceed at (n,$\\gamma$)$\\leftrightarrow$($\\gamma$,n) equilibrium, increased neutrino spallation rates, tending to enhance smoothing in post-processing, and increased beta-decay rates.  However, at signicant excitation, neutron separation energies are low enough that neutron emission may be a dominant decay mode.  This is discussed briefly along with the discussion of excited-state decays in the network calculation.  Section \\ref{grosstheory} of this paper is an overview of $\\beta$-decay calculations and average properties in calculating decay rates calculated using the more recent semi-gross theory. This includes a brief review of the semi-gross theory of $\\beta$-decay and a description of the calculations of single-particle states and their relationship to $\\beta$-decay in \\S\\ref{spstates}.  Energy levels are calculated in \\S\\ref{spcalc} using the Finite Range Droplet Model (FRDM).  Results of $\\beta$-decay calculations of excitated state nuclei with the adapted FRDM are discussed in \\S\\ref{results}, along with their potential astrophysical importance.  Application of these results to an r-process network calculation is made in \\S\\ref{r-process} with results from a model with several environmental parameter sets described in \\S\\ref{network_results}. Future work and possibilities are discussed in \\S\\ref{conclusion}, along with experimental possibilities. ",
        "conclusions": "\\label{conclusion} This work provides a study of the effects of excited state $\\beta$-decays on the r-process.  A preliminary method was used to evaluate the possible effects of $\\beta$-decay rates of excited-state nuclei. Though the accuracy of the model is limited by the knowledge of single particle levels, an approximate treatment allows one to gauge the magnitude of effects on the r-process and provide impetus for further study.  An empirical calculation was employed to find single-particle levels, and quantum numbers were deduced based on the level proximity to those of the spherical shell model. Although minor effects were found to result from inclusion of the excited state decays, there are also other possible effects that the inclusion of excited state nuclei may have on the r-process. One can imagine that if the excitation is due to the promotion of neutrons to higher-lying single-particle orbitals, then the photoneutron Q-value will decrease, and the ($\\gamma$,n) reaction rate may increase.  Thus the effect of an increased rate might be to shift the r-process path closer to $\\beta$-stability.  All models used in this calculation do a reasonable qualitative job of reproducing the solar abundance distribution for the mass region $80\\le{A}\\le{130}$.  The fact that the abundance at low mass is roughly independent of the type of model used (hot or cold) is an indicator that the nuclei in this region have decay rates not as heavily dependent on temperature as some of the higher mass nuclei. In the more massive nuclei, the sensitivity of the decay rates on temperature resulted in a more pronounced shift in the path as the r-process evolves through freezeout.  As mentioned previously, the dynamical treatment of the r-process is an important factor here in that the path continues to evolve even during freezeout, a result of $\\beta$-delayed neutron emission.  With the mass formula used in this evaluation, it was found that the low mass nuclei have a higher probability of emitting two neutrons during $\\beta$-decay than the higher mass nuclei, which have a higher probability of emitting a single neutron during $\\beta$-decay.  These available neutrons are recaptured, with the cross section roughly increasing with mass.  The net result is  a slight shift in the A$\\sim$195 abundance peak.  Despite the ability to predict a more reasonable abundance of the A$\\sim$190 nuclei, there is still a discrepancy between the predicted abundance distribution in the rare earth region and that of the solar system. The abundances of the rare earth elements (A$\\sim$165) are predicted to be lower in abundance by about an order of magnitude than the A$\\sim$130 peak, in fair agreement with that of the solar system.  This corresponds to the argument of the authors of reference \\cite{surman97}, who state that the rare-earth region is a robust feature of any dynamical calculation including post-production of the r-process progenitors.  However, nuclei with 130$<$A$<$160 are overproduced slightly, removing the effect of the rare earth region being manifest as a peak, hence the appearance of the abundance distributions in the figure.  It is obvious that no shell quenching has been employed in this preliminary model, as can be seen by the dip in the A$\\sim$180 abundance. While the abundance of the A$\\sim$180 nuclei relative to the A$\\sim$130 peak is similar to that of the solar system, the width of this dip is greater than that of the solar system abundance distribution. It has been mentioned \\cite{kratz93} that the underproduction of nuclei in this region might vanish if the quenching of shell closures for the very neutron-rich nuclei in this mass region is properly included.  Because of the large number of nuclei involved, the improved semi-gross theory has been used to globally calculate decay rates.  It is understood that the initial model represents a first attempt to gauge the effects of nuclear $\\beta$-decays in the hot environment of the r-process, and further study is warranted. In particular, a more accurate global calculation of decay rates is desired.  Currently, the measurements of $\\beta$-decay rates are limited to either ground-state nuclei or long-lived isomers.  However, with the advent of large neutron flux devices\\cite{petrasso96}, the measurement of the GT strength functions of nuclei in excited states from may become feasible within the next decade, allowing for experimental confirmation of the transition strengths of excited-state nuclei. \\ack The authors wish to acknowledge the helpful insight provided by K. Takahashi and comments by B.A. Brown. This work was funded by NSF grants PHY 9901241 and PHY9905241 and by the WMU Faculty Research and Creative Activities Support Fund (FRACASF) grant \\#06-005."
    },
    "0808/0808.3786_arXiv.txt": {
        "abstract": "In continuation of previous work, numerical results are presented, concerning relativistically counter-streaming plasmas. Here, the relativistic mixed mode instability evolves through, and beyond, the linear saturation -- well into the nonlinear regime. Besides confirming earlier findings, that wave power initially peaks on the mixed mode branch, it is observed that, during late time evolution wave power is transferred to other wave numbers. It is argued that the isotropization of power in wavenumber space may be a consequence of weak turbulence. Further, some modifications to the ideal weak turbulence limit is observed. Development of almost isotropic predominantly electrostatic -- partially electromagnetic -- turbulent spectra holds relevance when considering the spectral emission signatures of the plasma, namely bremsstrahlung, respectively magneto-bremsstrahlung (synchrotron radiation and jitter radiation) from relativistic shocks in astrophysical jets and shocks from gamma-ray bursts and active galactic nuclei. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.3924_arXiv.txt": {
        "abstract": "In the study of stars, the high energy domain occupies a place of choice, since it is the only one able to directly probe the most violent phenomena: indeed, young pre-main sequence objects, hot massive stars, or X-ray binaries are best revealed in X-rays. However, previously available X-ray observatories often provided only crude information on individual objects in the Magellanic Clouds. The advent of the highly efficient X-ray facilities XMM-Newton and Chandra has now dramatically increased the sensitivity and the spatial resolution available to X-ray astronomers, thus enabling a fairly easy determination of the properties of individual sources in the LMC. ",
        "introduction": "In the range of astronomical tools, X-rays are of particular interest. Indeed, the high-energy domain unveils the most energetic phenomena taking place in our Universe. Such processes are usually difficult to perceive at other wavelengths, though they provide important constraints for astrophysics.\\\\  In this context, the MCs are targets of choice. Their advantages are multiple: the known and small distance, together with the small angular size, small inclination, different metallicities, and recent star formation episodes are here as crucial as at other wavelengths. In addition, the low obscuration towards the MCs renders soft X-ray observations much easier while the numerous available data taken at other wavelengths ensure a correct, global analysis of the LMC X-ray sources.\\\\  X-rays associated with the LMC were first detected 40 years ago by a rocket experiment (\\cite{mar69}): the source appeared extended and soft, with a total luminosity estimated to 4$\\times 10^{38}$\\,erg\\,s$^{-1}$. It did not take long to distinguish a few individual sources in this X-ray emission, nicknamed LMC\\,X-1 to 6, thanks to the joint effort of the satellites Uhuru, Copenicus, OSO-7, and Ariel V (\\cite{leo71}, \\cite{rap74}, \\cite{mar75}, \\cite{gri77}). In the following decade, the Einstein observatory increased the number of known X-ray sources in the direction of the LMC to about a hundred (\\cite{lon81}, \\cite{wan91}). Finally, a sensitive survey undertaken by ROSAT provided another ten-fold increase in the total source number (\\cite[Haberl \\& Pietsch 1999a, Sasaki \\etal\\ 2000]{hab99a,sas00}).\\\\  However, only half of the detected X-ray sources truly belonged to the LMC and an even smaller fraction appeared to be associated with LMC stars. For example, of the 758 ROSAT-PSPC sources, only 144 were identified at first (\\cite[Haberl \\& Pietsch 1999a]{hab99a}): 15 as background AGNs or galaxies behind the LMC, 57 as foreground Galactic stars, 46 as SNRs and SNRs candidates, 17 as X-ray binaries (XRBs) and candidates, 9 as Supersoft sources (SSSs) and candidates. Using the observed X-ray properties (especially the hardness ratios), \\cite{hab99a} further proposed additional identifications (3 AGNs, 27 foreground stars, 9 SNRs and 3 SSSs) which yields a fraction of 20\\% of X-ray sources associated with LMC stars. Similar results were obtained with the ROSAT-HRI (\\cite[Sasaki \\etal\\ 2000]{sas00}): 397 detections among which 138 in common with the PSPC and 115 identified (10 AGNs, 52 foreground stars, 33 SNRs, 12 XRBs, 5 SSSs, and 3 hard sources which could either be AGNs or XRBs). If one considers the variable X-ray sources, the contamination by non-LMC objects is smaller. For the PSPC survey, the proposed identification of the 27 variable X-ray sources is 12 XRBs, 5 SSSs, 9 foreground stars and 1 Seyfert galaxy (\\cite[Haberl \\& Pietsch 1999b]{hab99b}); for the HRI survey, 26 variable sources were detected among which 8 XRBs, 4 SSSs, 6 foreground stars, 2 AGNs and 1 nova (\\cite[Sasaki \\etal\\ 2000]{sas00}): the fraction of LMC stellar objects among variable X-ray sources is thus 60\\%.\\\\  The current facilities possess much higher sensitivities and spatial/spectral resolution but unfortunately they also have a smaller field-of-view. Its non-zero extension penalizes the LMC, especially in comparison with the SMC ($\\sim$59 sq. deg. vs. $\\sim$18). This explains why, up to now, no full survey of the LMC has been performed with \\xmm\\ or \\ch. Nevertheless, smaller fields have been observed and it should be underlined that, though its coverage is patchy, the 2XMM catalog currently lists 5421 entries in the area of the ROSAT surveys! In addition, \\cite{hab03} reported the analysis of one deep \\xmm\\ observation of a northern region of the LMC (on the rim of the supergiant shell LMC4). While ROSAT had detected 34 sources in this field, \\xmm\\ data reveal 150 objects (detection limit 6$\\times 10^{32}$\\,erg\\,s$^{-1}$). In a selection of 20 bright or peculiar sources, the majority (10) are AGNs, but there are also 3 foreground stars, 2 SNRs, 4 HMXBs, and 1 SSSs \\footnote{This source was not confirmed by \\cite{kah08}.}. In addition, \\cite{sht05} analyzed 23 \\xmm\\ archival observations covering 3.8 sq. degrees of the LMC. With a detection limit of 3$\\times 10^{33}$\\,erg\\,s$^{-1}$, they detected 460 sources in the 2--8\\,keV band, in vast majority AGNs, and focused on 9 good XRBs candidates and 19 possible XRBs (see below for more details). Finally, a sensitive survey of the LMC in the hard X-ray domain (15\\,keV--10\\,MeV) has been performed with {\\sc Integral} (\\cite{got06}). Only a few sources have been detected: the X-ray binaries LMC\\,X-1, LMC\\,X-4, and PSR B0540--69, as well as two hard sources which might correspond to LMC binaries. These encouraging first results in both the soft and hard X-ray domains enlight the detection potential of the sensitive observatories of the current generation. \\\\ ",
        "conclusions": ""
    },
    "0808/0808.1921_arXiv.txt": {
        "abstract": "We present cosmological constraints arising from the first measurement of the  radial (line-of-sight)  baryon acoustic oscillations (BAO) scale in the large scale structure traced by the galaxy distribution. Here we use these radial BAO measurements at $z=0.24$ and $z=0.43$ to derive new constraints on dark energy and its equation of state for a flat universe, without any other assumptions on the cosmological model: $w = -1.14 \\pm 0.39$ (assumed constant), $\\Omega_m = 0.24^{+0.06}_{-0.05}$. If we drop the assumption of flatness and include previous cosmic microwave background and supernova data, we find $w = -0.974 \\pm 0.058$, $\\Omega_m = 0.271 \\pm 0.015$, and $\\Omega_k = -0.002 \\pm 0.006$, in good agreement with a flat cold dark matter cosmology with a cosmological consant. To our knowledge, these are the most stringent constraints on these parameters to date under our stated assumptions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.0320_arXiv.txt": {
        "abstract": "We analyse intermediate-resolution VLT FLAMES/Giraffe spectra of six ultra-compact dwarf (UCD) galaxies in the Fornax cluster. We obtained velocity dispersions and stellar population properties by full spectral fitting against {\\sc pegase.hr} models. Objects span a large range of metallicities (--0.95 to --0.23~dex), 4 of them are older than 8~Gyr. Comparison of the stellar and dynamical masses suggests that UCDs have little dark matter at best. For one object, UCD3, the Salpeter initial mass function (IMF) results in the stellar mass significantly exceeding the dynamical one, whereas for the Kroupa IMF the values coincide. Although, this object may have peculiar dynamics or/and stellar populations, the Kroupa IMF seems more realistic. We find that UCDs lie well above the metallicity--luminosity relation of early-type galaxies. The same behaviour is demonstrated by some of the massive Milky Way globular clusters, known to contain composite stellar populations. Our results support two following UCD formation scenarii: (1) tidal stripping of nucleated dwarf elliptical galaxies; (2) formation of tidal superclusters in galaxy mergers. We also discuss some of the alternative channels of the UCD formation binding them to globular clusters. ",
        "introduction": "A new class of compact dwarf galaxies, called UCDs, was discovered a decade ago in the Fornax cluster \\citep{Hilker+99,Drinkwater+00,PDGJ01}.  These objects, being brighter and much larger than globular clusters (GCs), but still far below luminosities and sizes of both dwarf elliptical (dE) and compact elliptical (cE) galaxies, fill an empty region on the Fundamental Plane \\citep{DD87}.  Their origin still remains a matter of debate; several alternatives are considered: (1) UCDs are the result of the evolution of primordial density fluctuations \\citep{PDGJ01}; (2) they have been formed through mergers of GCs or simply represent the extreme high-luminosity end of the GC luminosity function \\citep{MHI02}; (3) UCDs are nuclei of tidally stripped (``threshed'') nucleated dE (dE,N) galaxies \\citep{BCDS03} or dE,Ns with very low surface brightness; (4) UCDs are created as tidal superclusters during major mergers of galaxies \\citep{FK05,KJB06}.  Mass-to-light ratios of UCDs vary quite significantly \\citep{Drinkwater+03,Hasegan+05,Hilker+07}, suggesting the presence of dark matter in some of them. \\cite{Hasegan+05} propose to use $M/L$ ratio (i.e. presence of dark matter) as a criterion to distinguish between ``UCD galaxies'' and massive GCs.  Developing this idea, we conclude that the presence of dark matter in a compact stellar system rejects two formation scenarii -- in this case UCDs cannot be GCs neither they can be created as tidal superclusters during galaxy mergers. On the other hand, if the stellar population is not old and metal-poor, the primordial density fluctuation scenario will become implausible, leaving the only channel of UCD formation to be the tidal stripping of dE,Ns.  Stellar population analysis may also help to choose the formation scenario. Presently published data \\citep{MHIJ06,EGDH07} based on the analysis of absorption line strengths (Lick indices, \\citealp{WFGB94}) suggest that UCDs are old and rather metal-poor.  In this paper, we present stellar population parameters for 6 Fornax cluster UCDs, compare them with dE,N nuclei, and derive stellar masses to check for the presence of dark matter assuming different stellar IMFs. ",
        "conclusions": "\\subsection{Comparison of Stellar and Dynamical Masses}  Given the stellar population parameters and luminosities, we derive stellar masses of UCDs in our sample. The stellar mass estimates, computed from the mass-to-light ratios provided by {\\sc pegase.2} \\citep{FR97} for Salpeter and Kroupa et al. IMFs are given in Table~\\ref{tabmlcomp}. In the fourth column we provide the corrected dynamical masses, derived by re-normalising the values of \\cite{Hilker+07} by our velocity dispersion measurements ($M_{\\rm{d, corr}} = M_{\\rm{d}} (\\sigma / \\sigma_{\\rm{lit}})^2$. The fifth and sixth column contain the dark matter fractions estimated from Salpeter and Kroupa SSPs and corrected dynamical masses.  For UCD1, 2, 4, and 5 the Salpeter SSP stellar masses are consistent with the dynamical ones within uncertainties, although in average the stellar masses tend to be lower. The Kroupa et al. IMF decreases them more resulting in a 40--50 per cent (65 for UCD5) upper limits of the dark matter content. However, for UCD3 the stellar mass derived using the Salpeter IMF becomes significantly larger than the dynamical estimate (note negative dark matter fraction), whereas Kroupa IMF provides almost a perfect match between the two, suggesting zero dark matter content. Therefore, if we assume no object-to-object IMF variation, the observations are more in favour of the Kroupa et al. IMF.  There is a possibility that the dynamical model of UCD3 used by \\cite{Hilker+07} was not correct (for example, (1) the outer component of the UCD3 is not spherically symmetric or (2) there is significant rotation not taken into account, or (3) velocity dispersions are anisotropic), which may lead to an underestimated dynamical mass. At the same time, we cannot exclude that SSP models do not represent well the spectrum (e.g. the object contains a metal-rich sub-population, which is not properly modeled). In this case the stellar mass may be overestimated.  Given the large uncertainties of stellar population parameters and, consequently, stellar mass-to-light ratios, we cannot give a decisive answer on a question \\emph{``Is there dark matter in UCDs?''} However, the main conclusion we draw is that \\emph{UCDs are not dark matter dominated objects} and at present level of detection, the dark matter content can be explained by uncertainties of the measurements and looseness of the models used to derive dynamical and stellar masses.  \\begin{table} \\caption{Comparison of stellar masses of 6 UCDs for Salpeter (2) and Kroupa (3) IMFs; and the dynamical masses for 5 objects (5) from Hilker et al. (2007) corrected using our velocity dispersion estimations, (6)--(7) the dark matter content (per cent) for the Salpeter and Kroupa et al. IMF.\\label{tabmlcomp}} \\begin{tabular}{cccccc} \\hline $n$ & $M_{*{\\rm{Salp.}}}$ & $M_{*{\\rm{Kroupa}}}$ & $M_{\\rm{d, corr}}$ & DM$_{\\rm{S.}}$  & DM$_{\\rm{K.}}$ \\\\ & 10$^7 M_{\\odot}$ & 10$^7 M_{\\odot}$ & 10$^7 M_{\\odot}$ & \\% & \\% \\\\ \\hline 1 &   2.9$\\pm$0.7 &  2.1$\\pm$0.5 &  3.7$\\pm$0.5  &  20 & 45 \\\\ 2 &   2.3$\\pm$0.8 &  1.4$\\pm$0.5 &  2.4$\\pm$0.3  &   0 & 40 \\\\ 3 &  17.5$\\pm$3.0 & 12.8$\\pm$2.2 & 12.0$\\pm$2.4  &--45 &  0 \\\\ 4 &   2.9$\\pm$0.8 &  1.9$\\pm$0.6 &  4.0$\\pm$1.0  &  30 & 50 \\\\ 5 &   0.9$\\pm$0.2 &  0.5$\\pm$0.1 &  1.4$\\pm$0.5  &  35 & 65 \\\\ 6 &   2.3$\\pm$0.4 &  1.6$\\pm$0.4 & & & \\\\ \\hline \\end{tabular} \\end{table}  \\subsection{Metallicity-$M_B$ and metallicity-$\\sigma$ relations}  \\begin{figure} \\includegraphics[width=0.95\\hsize]{UCD_A496_sigma_met.ps}\\\\ \\includegraphics[width=0.95\\hsize]{UCD_A496_MB_met.ps} \\caption{Metallicity-velocity dispersion (top) and metallicity-luminosity (bottom) relations for early-type galaxies, UCDs, and GCs. The data sources are described in the text. Outlined crosses represent Galactic GCs with multi-component stellar populations (see Section 4.2 for details). \\label{figSLZ}} \\end{figure}  In Fig.~\\ref{figSLZ} we plot metallicities versus velocity dispersions and luminosities of the 6 Fornax cluster UCDs and 2 dE,N nuclei from our sample. They are compared to: (a) Milky Way GCs from \\cite{MvdM05}, outlined crosses show GCs with direct evidences of multiple stellar populations revealed by the analysis of colour-magnitude diagrams \\citep{Piotto08}; (b) Local Group dEs and dwarf spheroidals (dSph) from \\cite{Mateo98}; (c) Abell~496 low-luminosity early-type galaxies from \\cite{Chil08A496}; (d) a sample of intermediate-luminosity and bright early-type galaxies from \\cite{SGCG06a,SGCG06c}; two compact stellar systems in the Virgo cluster with the spectra available in SDSS DR6 \\citep{SDSS_DR6}: transitional UCD/cE ``M59cO'' \\citep{CM08} and VUCD~7, where the data have been processed exactly in the same way as for M59cO. The aperture size for the Abell~496 dE/dS0 is around 0.8~kpc, so the dE nuclei do not dominate the light, therefore metallicities and velocity dispersions should be closer to the global than to the central values, given flat velocity dispersion profiles usually observed in dEs \\citep{SP02, GGvdM03, vZSH04}.  In the metallicity-luminosity relation (Fig.~\\ref{figSLZ}, bottom panel) there is a continuous sequence $Z \\varpropto L_B^{0.45}$, spanning over 6 orders of magnitude in luminosity, formed by early-type galaxies: from the faintest dSph on the left to the brightest cluster ellipticals on the right. UCD galaxies lie significantly above (0.7--1.0~dex) this sequence, compared to the brightest Local Volume dSph's, having similar luminosities (see \\citealp{Mateo98} and references therein). In this sense, UCDs are similar to cE galaxies (e.g. \\citealp{Chilingarian+07}) having high metallicities for their luminosities, which probably represent end-products of the galaxy tidal threshing \\citep{BCD01}.  At the same time, on the metallicity-velocity dispersion plot (Fig.~\\ref{figSLZ}, top panel) the loci of UCDs practically coincide with those of dEs. We consider this as an argument for the scenario of tidal threshing of dE,Ns as a way to create UCDs. In this case, a velocity dispersion of a compact nucleus will not change very strongly, while the total luminosity of a progenitor will drop by several magnitudes.  For comparison with massive GCs \\citep{MvdM05} we chose a subsample of Galactic GCs with available measurements of velocity dispersions. Many GCs follow the behaviour of early-type galaxies. The three strongest outliers, namely NGC~104, NGC~6388, and NGC~6441, similarly to UCDs, reside significantly above the sequence of early-type galaxies on the metallicity -- luminosity plot (Fig~\\ref{figSLZ}). It is remarkable that the latter two exhibit direct evidences of multiple stellar populations \\citep{Piotto08}. Among the three other GCs (NGC~1851, NGC~2808, $\\omega$~Cen) demonstrating composite stellar populations only $\\omega$~Cen follows exactly the trend defined by early-type galaxies. We notice, that the third strongest outlier, 47~Tuc (NGC~104), being at least as massive as NGC~6388 does not have an evident double main sequence \\citep{Piotto08}.  In the frame of the tidal stripping scenario, we can also propose an explanation for the large spread of UCD metallicities ($-$0.25 to $-$0.93) on the [Fe/H]~vs.~$\\sigma$ plot. It may be a superposition of the two factors: (1) the relatively high spread of metallicities of dE progenitors of UCDs due to their own environmentally-driven evolution (see discussion in \\citealp{Chil08A496}); (2) different conditions during the tidal stripping, which may lead to some changes of the velocity dispersion values compared to the progenitors.  \\subsection{Comparison of dE,N nuclei and UCDs}  In Fig.~\\ref{figtZ} we compare ages and metallicities of the 6 Fornax cluster UCDs with nuclei of dE,Ns in the Fornax (2 objects, this study) and Virgo (26 objects) clusters, transitional cE/UCD object M59cO \\citep{CM08} and VUCD7, another UCD in the Virgo cluster. Stellar population parameters for 22 Virgo cluster galaxies (shown in red), as well as for VUCD7 and M59cO are obtained by analysing SDSS DR6 spectra. For the four remaining Virgo dE,Ns shown in light blue we used the results based on the 3D spectroscopic observations presented in \\cite{CPSA07,CSAP07}: diamonds with error-bars correspond to the ages and metallicities of the nuclei and the blue vectors point to the parameters of the ``main bodies'' of the galaxies.  Apart from the 2 intermediate age objects (UCD4 and UCD5), all UCDs are old, at the same time spanning a large range of metallicities. Most of the dE,Ns exhibit considerably younger stellar populations than UCDs. However, there is a number of old dE,N nuclei (including FCC~182) with ages comparable to those of UCDs. A scenario, assuming that dE,N nuclei are results of repeated or extended star formation episodes in the dE centres, leading to metal enrichment, can explain the observed quite high metallicities of dE,N nuclei.  \\begin{figure} \\includegraphics[width=0.95\\hsize]{agemet_ucd_new.ps} \\caption{Comparison of ages and metallicities of UCDs with a sample of Virgo cluster dE,N galaxies from SDSS.\\label{figtZ}} \\end{figure}  \\subsection{Origin of UCD galaxies}  The stellar population parameters obtained by our SSP fitting, namely metallicities higher than $-$1.0 dex, allow us to exclude the scenario of evolving primordial density fluctuations \\citep{PDGJ01}, because one would expect much lower metallicities for objects at this mass range.  Low dark matter content leaves space for all remaining channels of the UCD formation: \\cite{BCDS03} showed that in case of dE stripping (``threshing'') the progenitor's nucleus must not be dark matter dominated; the two other alternatives, GC merging and formation of UCDs as tidal superclusters, assume zero dark matter content.  However, the scenario of merging GCs \\citep{MHI02} fails to explain why we do not observe metal-poor UCDs. It is known that GCs exhibit a dichotomy in the metallicity distribution (e.g. review by \\citealp{BS06}), but observed UCDs correspond only to metal-rich GCs. Why don't metal-poor GCs merge? Moreover, in case of a merger of metal-poor and metal-rich GCs of the same mass, the resulting luminosity-weighted metallicity in our wavelength range will be lower than the mean value, because metal-poor stellar populations have lower $M/L$ ratios than metal-rich ones. Composite stellar populations observed in massive Galactic GCs \\citep{Piotto08} comprise two to four SSPs, sometimes significantly different in metallicities, which is evident from deep colour-magnitude diagrams. These objects look good candidates for the GC merger scenario. In addition, we do see metal-poor ``composite'' GCs (NGC~1851, NGC~2808, $\\omega$~Cen). Indeed, they are an order of magnitude (except $\\omega$~Cen, where at least four SSPs are evident) fainter than the UCDs we are discussing here.  On the statistical basis, UCDs are too frequent to be representatives of the high-luminosity end of the GC luminosity function (GCLF). After the initial discussion in \\cite{MHI02} a significant number of UCDs was discovered. The extrapolation of the GCLF (see e.g. \\citealp{ML07}) towards bright objects results in the statistical over-population of $M_V < -11$ objects. Although, the exact value depends on the adopted GCLF parameters and representation (i.e. Gaussian or $t_5$), we consider this fact important, making this channel of UCD formation scarcely probable.  The old ages of most UCDs, compared to dE,N nuclei, suggest that if we consider the dE,N tidal stripping scenario, it must have happened a long time ago. However, there is a difficulty in explaining the formation of the most metal-rich UCDs, because 8--10~Gyr ago dE,N nuclei must have been less metal rich than presently observed. A possible explanation is a tidal stripping of the most massive dE representatives (dE/E transitional objects) such as IC~3653 or FCC~182.  Another possibility is to create them as stellar superclusters (Fellhauer \\& Kroupa 2005) during interactions of massive galaxies, in this case metal-rich population formed from the metal-rich gas of the progenitors will be observed in the UCDs. Both scenarii are compatible with low dark matter content. A possible diagnosis is to measure $\\alpha$/Fe abundance ratios: populations formed in a short and intense star formation episode will be $\\alpha$-overabundant (e.g. \\citealp{Matteucci94}). With the present low S/N UCD data we are not able to carry out this test.  Finally, we are left with the two alternatives of UCD formation: UCDs with low metallicities ($[\\mbox{Fe/H}] < -0.5$~dex) are in favour of dE,N tidal stripping, while tidally created superclusters better explain metal-rich UCDs. At present we cannot exclude the diversity of the UCD origin suggested by \\cite{MHIJ06}."
    },
    "0808/0808.0044_arXiv.txt": {
        "abstract": "We consider the problem of optimal weighting of tracers of structure for the purpose of constraining the non-Gaussianity parameter $\\fnl$. We work within the Fisher matrix formalism expanded around fiducial model with $\\fnl=0$ and make several simplifying assumptions. By slicing a general sample into infinitely many samples with different biases, we derive the analytic expression for the relevant Fisher matrix element. We next consider weighting schemes that construct two effective samples from a single sample of tracers with a continuously varying bias. We show that a particularly simple ansatz for weighting functions can recover all information about $\\fnl$ in the initial sample that is recoverable using a given bias observable and that simple division into two equal samples is considerably suboptimal when sampling of modes is good, but only marginally suboptimal in the limit where Poisson errors dominate. ",
        "introduction": "\\label{sec:introduction}  The currently most attractive theory for the emergence of structure in the Universe is inflation \\cite{Starobinsky:1979ty,1981PhRvD..23..347G,1982PhLB..108..389L,1982PhRvL..48.1220A}. It is generically successful at diluting the primordial defects to undetectable densities and predicts a nearly-flat universe with nearly scale invariant spectrum of primordial fluctuations that are normally distributed and extend to scales larger than horizon \\cite{1981JETPL..33..532M,1982PhLB..115..295H,1982PhRvL..49.1110G,1982PhLB..117..175S,1983PhRvD..28..679B}. To understand details of the inflation, on must look at detailed predictions of different models. Non-Gaussianity of the primordial curvature perturbations, i.e small departures from the normal distribution of fluctuations is one aspect in which models of inflation differ.  Recently, non-Gaussianity of the local $\\fnl$ type has received a renewed attention.  This type of non-Gaussianity is characterised by a quadratic correction to the potential \\citep{1990PhRvD..42.3936S,1994ApJ...430..447G,2000MNRAS.313..141V,2001PhRvD..63f3002K}: \\begin{equation} \\Phi = \\phi + \\fnl \\left( \\phi^2 - \\left<\\phi^2 \\right> \\right), \\label{fnl} \\end{equation} where $\\phi$ is the primordial potential assumed to be a Gaussian random field and $\\fnl$ describes the amplitude of the correction during the matter domination era.  There are two main reasons for this renewed interest. First, there is a hint of a detection in the cosmic microwave data \\cite{2008PhRvL.100r1301Y} and several non-detections \\cite{2007JCAP...03..005C,2008arXiv0803.0547K,2008arXiv0802.3677H}. Second, a new method for its detection has been recently proposed in \\cite{Dalal:2007cu}. This method uses biased tracers of structure for which it can be shown that local-type of non-Gaussianity leads to a very particular scale-dependence of the bias \\begin{equation} \\Delta b = \\fnl (b-1) u (k), \\label{eq:4} \\end{equation} where $\\Delta b$ is the bias induced by non-Gaussianity, $b$ is the tracer's intrinsic bias and $u$ is given by \\begin{equation} u(k) = \\frac{3 \\delta_c \\Omega_m H_0^2}{c^2 k^2 T(k) D(z)}, \\label{eq:3} \\end{equation} where $T(k)$ is the matter transfer function normalised to unity at $k=0$, $D(z)$ is the growth function normalised to $(1+z)^{-1}$ in the matter era, $\\delta_c=1.68$ is the linear over-density at collapse for the spherical collapse model and other symbols have their usual meaning. Note that $\\Delta b$ becomes significant only at large scales, where non-linearities and scale-dependent bias are expected to be small and therefore offers a surprisingly clean probe of non-Gaussianity.  This equation has been re-derived, scrutinised and better understood in the subsequent work \\cite{2008ApJ...677L..77M,2008arXiv0805.3580S,2008arXiv0806.1046A,2008arXiv0806.1061M}.   A first application of this method to the real data using a wide variety of tracers of large scale has recently shown the promise of this method \\cite{2008arXiv0805.3580S,2008arXiv0806.1046A}. The derived constraints are already competitive with those coming from the cosmic microwave background.  In that work, the constraints were derived by comparing the power spectrum of the distribution of tracers with those predicted by the theory. At largest scales, where the effect coming from the non-Gaussianity is the largest, the method suffers from the sample variance. In other words, the finite number of large-scale modes in any survey severely limits our ability to measure the power spectrum.  Recently, Seljak has suggested a method of circumventing this limitation \\citep{2008arXiv0807.1770S}. This method essentially considers two differently biased tracers that sample the same volume. The ratio of amplitudes of a single mode for the two tracers will give the ratio of the two biases $(b_1+\\Delta b_1(\\fnl))/(b_2+\\Delta b_2(\\fnl))$, but the amplitude of the primordial mode cancels out. One thus measures the auto correlation power spectra of the two tracers in the same volume. By taking the ratio of these two spectra, one can put a constraint on the value of $\\fnl$, which is independent on the primordial field and thus unaffected by the sample variance. The biases $b_1$ and $b_2$ can be derived from the amplitude of small scale fluctuations, where sampling variance is not a problem and hence, one extremely well measured large-scale mode is in principle enough to constrain $\\fnl$.  A more robust technique would be to assume nothing about the matter power spectrum and derive limits on $\\fnl$ from limits on the scale dependence of ratio of $b+\\Delta b$. This would protect measurements of $\\fnl$ from systematics arising from, for example,  massive neutrinos.   In practice, one rarely has two distinct samples with a well-defined bias.  In this work, we extended the analysis by considering a single tracer of the underlying field that spans a range of biases and attempt to answer the question of how to optimally analyse such tracer. The approach we take is to create two effective samples and to optimally weight the tracer's constituents. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have analysed the problem of optimal weighting of biased tracers of structure with the goal of extracting maximum information about the non-Gaussianity parameter $\\fnl$. We have derived the minimum error on $\\fnl$ by considering slices that are infinitely thin in bias. We have shown that a simple weighting scheme of Equations \\eqref{eq:1} and \\eqref{eq:2} obtains the same constraining power. General division of the full sample into two subsamples can be considerably sub-optimal even when mass at which the samples are divided is carefully chosen.  The optimal weighting scheme of Equations \\eqref{eq:1} and \\eqref{eq:2} is surprisingly simple. In fact, the product $P\\bar{n}$ does not come into weighting at all - this is a lucky coincidence, which allows us to use the same optimal weighting for every mode, rather to optimize weighting around some fiducial wave-vector.  The result in this paper is subject to the assumptions outlined in the Section \\ref{sec:appr-limit-this} of this paper. If these assumptions are violated, the weighing is sub-optimal, but probably nevertheless beneficiary. Since any division into two samples by e.g. an absolute magnitude cut requires some knowledge of bias, the implementation of the scheme proposed in this paper is likely to be very simple.  How can this be put in practice? In this work we have used halo mass $M$ as a proxy for the bias. However, our analysis is completely general and one can replace the host halo mass with any variable that is monotonically linked to the bias. For example, one could take luminous red galaxies (LRGs) and determine their bias by splitting the entire sample into several subsamples in different luminosity bins and the constrain a smooth function $b(L)$, which describes the variation of galaxy bias with its luminosity, using modes which are not affected by the $\\fnl$. One would next construct two effective samples by optimally weighting the original sample using Equations \\eqref{eq:1} and \\eqref{eq:2} and replacing $b(M)$ with $b(L)$. In the next step, auto and cross-correlation power spectra of these two samples should be calculated, taking into account the Poisson error correlation between the two. At this step, one can use the cross-correlation spectra to check for the amount of stochasticity, which has been assumed to be negligible in this work. Finally, $\\fnl$ should be constrained using these power spectra as input."
    },
    "0808/0808.2870_arXiv.txt": {
        "abstract": "We report new precision measurements of the properties of our Galaxy's supermassive black hole. Based on astrometric (1995-2007) and radial velocity (2000-2007) measurements from the W. M. Keck 10-meter telescopes, a fully unconstrained Keplerian orbit for the short period star S0-2 provides values for the distance (R$_0$) of 8.0 $\\pm$  0.6 kpc, the enclosed mass (M$_{bh}$) of 4.1 $\\pm$ 0.6 $\\times$ $10^6 M_{\\odot}$, and the black hole's radial velocity, which is consistent with zero with 30 km/s uncertainty. If the black hole is assumed to be at rest with respect to the Galaxy (e.g., has no massive companion to induce motion), we can further constrain the fit and obtain R$_0$ = 8.4 $\\pm$ 0.4 kpc and M$_{bh}$ = 4.5 $\\pm$ 0.4 $\\times$ $10^6 M_{\\odot}$. More complex models constrain the extended dark mass distribution to be less than 3-4 $\\times$ $10^5 M_{\\odot}$ within 0.01 pc, $\\sim$100x higher than predictions from stellar and stellar remnant models. For all models, we identify transient astrometric shifts from source confusion (up to 5x the astrometric error) and the assumptions regarding the black hole's radial motion as previously unrecognized limitations on orbital accuracy and the usefulness of fainter stars.  Future astrometric and RV observations will remedy these effects. Our estimates of R$_0$ and the Galaxy's local rotation speed, which it is derived from combining R$_0$ with the apparent proper motion of Sgr A*, ($\\theta_0$ = 229 $\\pm$ 18 km s$^{-1}$), are compatible with measurements made using other methods. The increased black hole mass found in this study, compared to that determined using projected mass estimators, implies a longer period for the innermost stable orbit, longer resonant relaxation timescales for stars in the vicinity of the black hole and a better agreement with the M$_{bh}$-$\\sigma$ relation. ",
        "introduction": "\\label{sec:intro}  Ever since the discovery of fast moving (v $>$ 1000 km s$^{-1}$) stars within 0.$\\tt''$3 (0.01 pc) of our Galaxy's central supermassive black hole (Eckart \\& Genzel 1997; Ghez et al. 1998), the prospect of using stellar orbits to make precision measurements of the black hole's mass (M$_{bh}$) and kinematics, the distance to the Galactic center (R$_0$) and, more ambitiously, to measure post-Newtonian effects has been anticipated (Jaroszynski 1998, 1999; Salim \\& Gould 1999; Fragile \\& Mathews 2000; Rubilar \\& Eckart 2001; Weinberg, Milosavlejic \\& Ghez 2005; Zucker \\& Alexander 2007; Kraniotis 2007; Will 2008). An accurate measurement of the Galaxy's central black hole mass is useful for putting the Milky Way in context with other galaxies through the apparent relationship between the mass of the central black hole and the velocity dispersion, $\\sigma$, of the host galaxy (e.g., Ferrarese  \\& Merrit 2000; Gebhardt et al. 2000; Tremaine et al. 2002). It can also be used as a test of this scaling, as the Milky Way has the most convincing case for a supermassive black hole of any galaxy used to define this relationship. Accurate estimates of R$_0$ impact a wide range of issues associated with the mass and structure of the Milky Way, including possible constraints on the shape of the dark matter halo and the possibility that the Milky Way is a lopsided spiral (e.g., Reid 1993; Olling \\& Merrifield 2000; Majewski et al. 2006).  Furthermore, if measured with sufficient accuracy ($\\sim$1\\%), the distance to the Galactic center could influence the calibration of standard candles, such as RR Lyrae stars, Cepheid variables and giants, used in establishing the extragalactic distance scale. In addition to estimates of $M_{bh}$ and R$_0$, precision measurements of stellar kinematics offer the exciting possibility of detecting deviations from a Keplerian orbit. This would allow an exploration of a possible cluster of stellar remnants surrounding the central black hole, suggested by Morris (1993), Miralda-Escud{\\'e} \\& Gould(2000), and Freitag et al. (2006). Estimates for the mass of the remnant cluster range from $10^4 - 10^5 M_{\\odot}$ within a few tenths of a parsec of the central black hole.  Absence of such a remnant cluster would be interesting in view of the hypothesis that the inspiral of intermediate-mass black holes by dynamical friction could deplete any centrally concentrated cluster of remnants.  Likewise, measurements of post-newtonian effects would provide a test of general relativity, and, ultimately, could probe the spin of the central black hole.  Tremendous observational progress has been made over the last decade towards obtaining accurate estimates of the orbital parameters of the fast moving stars at the Galactic center.  Patience alone permitted new astrometric measurements that yielded the first accelerations (Ghez et al. 2000; Eckart et al. 2002), which suggested that the orbital period of the best characterized star, S0-2, could be as short as 15 years.  The passage of more time then led to full astrometric orbital solutions (Sch\\\"odel et al. 2002, 2003; Ghez et al. 2003, 2005a), which increased the implied dark mass densities by a factor of $10^4$ compared to earlier velocity dispersion work and thereby solidified the case for a supermassive black hole.  The advent of adaptive optics enabled radial velocity measurements of these stars (Ghez et al. 2003), which permitted the first estimates of the distance to the Galactic center from stellar orbits (Eisenhauer et al. 2003, 2005).  In this paper, we present new orbital models for S0-2.  These provide the first estimates of the distance to the Galactic center and limits on the extended mass distribution based on data collected with the W. M. Keck telescopes. The ability to probe the properties of the Galaxy's central supermassive black hole has benefitted from several advancments since our previous report (Ghez et al. 2005).  First, new astrometric and radial velocity measurements have been collected between 2004 and 2007, increasing the quantity of kinematic data available.  Second,  the majority of the new data was obtained with the laser guide star adaptive optics system at Keck, improving the quality of the measurements (Ghez et al. 2005b; Hornstein et al. 2007). These new data sets are presented in \\S\\ref{sec:obs}. Lastly, new data analysis has improved our ability to extract radial velocity estimates from past spectroscopic measurements, allowing us to extend the radial velocity curve back in time by two years, as described in \\S\\ref{sec:data_analysis}.  The orbital analysis, described in \\S\\ref{sec:orbit}, identifies several sources of previously unrecognized biases and the implications of our results are discussed in \\S\\ref{sec:disc}. ",
        "conclusions": ""
    },
    "0808/0808.2277_arXiv.txt": {
        "abstract": "{The quasi-Hilda comets (QHCs), being in unstable 3:2 Jovian mean motion resonance, are considered a major cause of temporary satellite capture (TSC) by Jupiter. Though the QHCs may be escaped Hilda asteroids, their origin and nature have not yet been studied in sufficient detail. Of particular interest are long TSCs/orbiters. Orbiters -- in which at least one full revolution about the planet is completed -- are rare astronomical events; only four have been known to occur in the last several decades. Every case has been associated with a QHC: 82P/Gehrels 3; 111P/Helin-Roman-Crockett; P/1996 R2 (Lagerkvist); and the possibly QHC-derived D/1993 F2 (Shoemaker-Levy 9, SL9). } {We focus on long TSC/orbiter events involving QHCs and Jupiter. Thus we survey the known QHCs, searching for further long TSCs/orbiters over the past century. } {First, we confirmed the long TSC/orbiter events of 82P, 111P, and 1996 R2 in order to test our method against previous work, applying a general $N$-body Newtonian code. We then used the same procedure to survey the remaining known QHCs and search for long TSC/orbiter events. } {We newly identified another long TSC/orbiter: 147P/Kushida-Muramatsu from 1949 May 14$^{+97 {\\rm days}}_{-106 {\\rm days}}$--1961 July 15. Our result is verified by integrations of 243 cloned orbits which take account of the present orbital uncertainty of this comet. The event involves an $L_{\\rm 2} \\to L_{\\rm 1}$ transition as with 82P and 1996 R2; this may represent a distinct subtype of TSCs from QHC derived ($L_{\\rm 1} \\to$) longer captures exemplified by 111P and (probably) SL9, though this classification is still only based on a small database of TSCs. } {This is the third long TSC and the fifth orbiter to be found, thus long TSC/orbiter events involving Jupiter have occurred once per decade. Two full revolutions about Jupiter were completed and the capture duration was $12.17^{+0.29}_{-0.27}$ years; both these numbers rank 147P as third among long TSC/orbiter events, behind SL9 and 111P. This study also confirms the importance of the QHC region as a dynamical route into and out of Jovian TSC, via the Hill's sphere. } ",
        "introduction": "Among all the asteroids that have been recorded up to the present in the asteroid database, e.g., the ``JPL Small-Body Database\" ({\\tt http://ssd.jpl.nasa.gov/}), a large number (more than 1000 including unnumbered objects) are known to populate the region of the 3:2 mean motion resonance (MMR) with Jupiter, in the outer main belt. These are the ``Hilda asteroids\" (Schubart \\cite{schubart68}, \\cite{schubart82}, \\cite{schubart}; Ip \\cite{ip}; Yoshikawa \\cite{yoshikawa}; Franklin et al. \\cite{franklin}; Nesvorn\\'y \\& Ferraz-Mello \\cite{nesvorny}). Their semimajor axes, $a$, concentrate in the range 3.7 AU $\\le a \\le$ 4.2 AU, at eccentricities $e \\le 0.3$, and inclinations $i \\le 20^\\circ$ (Zellner et al. \\cite{zellner}). This results in a range of the Tisserand parameter with respect to Jupiter, $T_{\\rm J}$, of $\\sim 2.90$--3.05, where $T_{\\rm J} = a_{\\rm J}/a + 2 \\sqrt {a/a_{\\rm J}(1 - e^2)} \\cos I$, with $a_{\\rm J}$ being the semimajor axis of Jupiter and $I$ the mutual inclination between the orbits. The critical arguments for the Hildas, $\\phi = 3\\lambda_{\\rm J}-2\\lambda-\\varpi$, librate about $0^\\circ$, being stable in the long-term ($\\lambda$ is mean longitude, $\\varpi$ is longitude of perihelion, and J indicates Jupiter). As regards physical properties, the low-albedo D- and P-types are more abundant in Hildas' surface colours than the small fraction of C-types (Dahlgren \\& Lagerkvist \\cite{dahlgren}; Dahlgren et al. \\cite{dahlgren97}; Gil-Hutton \\& Brunini \\cite{gil-hutton}; Licandro et al. \\cite{licandro}). The surface colour of D- and P-type asteroids, such as Hildas and Trojans in the outer main belt, corresponds well with that of cometary nuclei (Fitzsimmons et al. \\cite{fitzsimmons}; Jewitt \\cite{jewitt02}), which means that they are covered with a similar mineralogical surface, suggestive of a common origin.  More than 50 Jovian irregular satellites are known at present (Jewitt \\& Haghighipour \\cite{jewitt}). \\'Cuk \\& Burns (\\cite{cuk}) pointed out that the progenitor of the main prograde cluster, the Himalia family, was plausibly derived from the Hildas long ago, if it was captured by a gas-drag assisted mechanism. Thus tracing the origin and nature of irregular satellites is very significant for studying the accretion processes in the early solar system. Satellite capture mechanics in the circular restricted three-body problem (CR3BP) or the $N$-body problem has often been investigated (e.g., H\\'enon \\cite{henon}; Huang \\& Innanen \\cite{huang}; Tanikawa \\cite{tanikawa}; Murison \\cite{murison}; Brunini et al. \\cite{brunini}; Nesvorn\\'y et al. \\cite{nesvorny03}, \\cite{nesvorny07}). Reflectance spectra of Jovian irregulars, being dominated by D- and C-types (Luu \\cite{luu}; Grav et al. \\cite{grav}), are comparable to those of Hildas.  During the past half century, several Jupiter family comets (JFCs; cf.\\ Levison \\cite{levison}) have stayed in or near the Hilda zone, although being in unstable 3:2 MMR with Jupiter. Some of them have been transferred from outside to inside, \\textit{vice versa}, or from inside to inside of Jupiter's orbit by undergoing a temporary satellite capture (TSC) by Jupiter (e.g., Carusi \\& Valsecchi \\cite{carusi79}; Tancredi et al. \\cite{tancredi}). Such a JFC is called a ``quasi-Hilda comet\" (QHC) by Kres\\'ak (\\cite{kresak}), who identified three such objects: 39P/Oterma, 74P/Smirnova-Chernykh, and 82P/Gehrels 3. Di Sisto et al. (\\cite{di sisto}) integrated the motions of 500 fictitious Hildas for $\\sim 10^9$ years, and found that most of them escaped from the Hilda zone into the JFC population, i.e., left the 3:2 MMR and evolved quickly on to unstable orbits: such a chaotic diffusion from the Hilda zone has also been demonstrated by Nesvorn\\'y \\& Ferraz-Mello (\\cite{nesvorny}). In addition, large-scale collisional processes, such as the late heavy bombardment, might also release small bodies from the Hilda zone into JFCs, e.g., see Gil-Hutton \\& Brunini (\\cite{gil-hutton00}). Hence, some QHCs may indeed be such escaped Hildas themselves. Recently, Toth (\\cite{toth}) updated the QHC list, finding a total of 17 members (see Section 3). This includes bodies such as the renowned Comet D/1993 F2 (Shoemaker-Levy 9, SL9) that have undergone TSC by Jupiter and then disappeared after colliding with the planet. Surface spectroscopic (or colorimetric) measurements for QHCs have only been carried out for 82P (De Sanctis et al. \\cite{de sanctis}), the results of which also indicate a taxonomic D-type.  Among the QHCs, 39P/Oterma was the first known to be temporarily captured by Jupiter, in 1936--1938 (Marsden \\cite{marsden}). However, this comet flew through the region near Jupiter over a rather short time, during which the comet did not complete a full revolution orbiting about the planet. In contrast, unlike 39P's ``fly-through\" capture, there is a different kind of TSC, in which at least one full revolution about the planet is completed; we deal with these in the present paper. Following Kary \\& Dones (\\cite{kary}) we call such objects ``orbiters\". These are often characterized by a long capture with very small perijove distance, usually lasting for $\\sim 10$ years or more. Not all orbiters become such long TSCs ($> 10$ yr), although of course they last longer than the fly-through type. SL9 is a representative case for both long TSCs and orbiters. This comet was pointed out to have possibly been QHC-derived before its tidal disruption on passing through perijove at less than 1.5 Jovian equatorial radii ($R_{\\rm J}$, where $R_{\\rm J}=71492.4$ km), i.e., within the Roche limit, in 1992 July and its subsequent collision with Jupiter in 1994 July (Nakano \\& Marsden \\cite{nakano93}; Sitarski \\cite{sitarski}; Benner \\& McKinnon \\cite{benner}). If it is QHC-derived, then SL9 is the only QHC so far that has been orbiting the planet as a TSC at the time of discovery (Benner \\& McKinnon \\cite{benner}). By numerically integrating SL9's pre-collision orbital motion, several studies showed that the TSC duration of SL9 lasted for 50 years or more, during which the comet completed more than 30 revolutions orbiting about Jupiter, making it nominally the longest known TSC (Carusi et al. \\cite{carusi}; Benner \\& McKinnon \\cite{benner}; Chodas \\& Yeomans \\cite{chodas}). However, according to Benner \\& McKinnon (\\cite{benner}), SL9 was the most chaotic known object in the solar system with an effective Lyapunov time of only $\\sim10$ years on its jovicentric orbit. Thus it is difficult to determine with certainty SL9's pre-capture orbit and its true TSC duration. Nevertheless, Benner \\& McKinnon's additional statistical analysis of distributions in $a$-$e$ space and $T_{\\rm J}$ values, based on back integrations of the various SL9 fragments, revealed a possible QHC origin of SL9. The work of Kary \\& Dones (\\cite{kary}) supports this possibility: they traced the motions of numerous fictitious JFCs for $\\sim 10^5$ years and found that half the SL9-like very long captures $>50$ years were due to QHCs. They estimated that the frequency of such a very long TSC (= ``long capture\" as designated by them) is extremely rare, only 0.02\\% of all the TSC events in their simulations. Interestingly, impacts on Jupiter are more frequent than the very long TSCs by a factor of 8--9. They also evaluated that long TSCs (= ``orbiters bound $> 10$ yr\" as designated by them) and orbiters are still rare events, at the respective levels of 0.8\\% ($\\supset$ very long TSCs) and 2\\% ($\\supset$ long TSCs) relative to all events, with about 98\\% being the short TSC type which contains the fly-through events. Carusi \\& Valsecchi (\\cite{carusi79}) had earlier confirmed the rarity of orbiter events, simulating motions of fictitious small bodies as well.  In the jovicentric Keplerian system, a TSC (especially a long TSC/orbiter) occurs whenever a small body passes near one of the collinear libration points $L_{\\rm 1}$ or $L_{\\rm 2}$ in the CR3BP of the Sun-Jupiter-(third) body system with very low velocity, i.e., effectively becoming bound by Jupiter when it enters the Hill's region with near-zero velocity. After that, the bound small body revolves about Jupiter on an elliptical jovicentric orbit until it again passes through the region near either $L_{\\rm 1}$ or $L_{\\rm 2}$ and escapes from the Jovian system. Considering such a transition using newly developed dynamical systems techniques based on a Hamiltonian formulation in the CR3BP, Koon et al. (\\cite{koon}) and Howell et al. (\\cite{howell}) demonstrated that a TSC by Jupiter occurs when the small body passes through a region inside the invariant manifold structure related to periodic halo orbits around $L_{\\rm 1}$ or $L_{\\rm 2}$ in the Hill's region. The TSC (or its duration) is usually defined by the jovicentric Kepler energy, $E_{\\rm J}$, being negative, $E_{\\rm J} < 0$, with the additional condition that the bound small body must be within the jovicentric sphere of gravitational influence: Kary \\& Dones (\\cite{kary}) set its boundary at 3 Hill's sphere radii (=1.065 AU) of Jupiter. However, Howell et al. (\\cite{howell}) defined the TSC duration as the residence time in the Hill's region. The former generally lasts longer than the latter, and here we regard the former as the TSC duration. The dynamics involved in TSC is quite different from that of quasi-satellites in 1:1 libration with Jupiter far outside the Hill's region (Wiegert et al. \\cite{wiegert}; Kinoshita \\& Nakai \\cite{kinoshita}).  Apart from SL9, only three orbiters have been known to occur. Every case has been associated with a QHC: 82P; 111P/Helin-Roman-Crockett; and, though it is not in Toth's (\\cite{toth}) QHC list, P/1996 R2 (Lagerkvist). These TSCs were found by Rickman (\\cite{rickman}), Tancredi et al. (\\cite{tancredi}), and Hahn \\& Lagerkvist (\\cite{hahn}), respectively, and are discussed further in Section 2. The occurrence of these few events during the last several decades is consistent with the rarity of long TSC/orbiter events suggested by Kary \\& Dones (\\cite{kary}). Although 74P encountered Jupiter at distances of 0.24 AU and 0.47 AU in 1955 October and 1963 September respectively, it was not bound to Jupiter (Rickman \\cite{rickman}; Carusi et al. \\cite{carusi85b}); this comet is, however, expected to experience a TSC by Jupiter in this century (Carusi et al. \\cite{carusi85b} and see also Section 4). There further exist some JFCs that encountered Jupiter more closely than several QHCs involved in TSC, e.g., 16P/Brooks 2; D/1770 L1 (Lexell); 81P/Wild 2; but they were not captured by the planet owing to their high-velocity encounters (Carusi et al. \\cite{carusi85a}; Emel'yanenko \\cite{emel'yanenko}).  Therefore studying the origin and nature of QHCs is of great importance from various astronomical points of view mentioned above, especially as regards their origin and being possibly related to the Hildas and the Jovian irregular satellites. Such studies may provide unique knowledge and clues about formation processes in the early solar system. Here we focus on long TSC/orbiter events involving Jupiter and QHCs. First (Section 2), applying a general $N$-body Newtonian code, we reconfirmed the long TSC/orbiter events of 82P, 111P, and 1996 R2. Then (Section 3) we used the same procedure to search Toth's (\\cite{toth}) QHCs list for other objects that have become long TSCs/orbiters in the past century. Eventually, we successfully found another long TSC/orbiter, occurring in the mid-20th century, 147P/Kushida-Muramatsu. ",
        "conclusions": "On the basis of our investigations above, we have presented a newly identified TSC of a comet by Jupiter. This is in the rare, orbiter class of TSCs and involves 147P from 1949 May 14$^{+97 {\\rm days}}_{-106 {\\rm days}}$--1961 July 15. This is the third long TSC of $>10$ years and the fifth orbiter found, so that Jupiter's long TSCs/orbiters have occurred once per decade. The completion of two full revolutions about Jupiter and the capture duration of $12.17^{+0.29}_{-0.27}$ years rank 147P as third in both these numbers among known orbiters, behind SL9 and 111P.  \\subsection{TSC Classification}  Following Kary \\& Dones (\\cite{kary}), we classify the known TSCs as follows: 39P as fly-through, 82P and 1996 R2 as orbiters, 111P and 147P as long TSCs ($>10$ years), and SL9 as a very long TSC ($>50$ years).  On the other hand, depending on the TSC characteristics, Howell et al. (\\cite{howell}) defined two TSC types: Type 1 as a 39P-like fly-through; Type 2 as a long lasting capture like 111P, where the comet experiences more than one close encounter with Jupiter while in the TSC region. Here we tentatively divide Type 2 into two subtypes: Type 2A as the $L_{\\rm 2} \\to L_{\\rm 1}$ transition such as 82P, 1996 R2 and 147P; Type 2B as the QHC-derived ($L_{\\rm 1} \\to$) longer capture than 2A, as with 111P and SL9. Further TSC classifications may be possible in the future, if different kinds of TSC are found.  \\subsection{Future TSCs}  We also surveyed all the known QHCs, integrating their orbital motions forward for 100 years to check for future long TSCs/orbiters. We found that 111P will undergo a long TSC/orbiter phase with 6 perijove passages, and minimum $E_{\\rm J} \\sim -3.19$, from 2068 April 20--2086 June 09 (duration $\\sim 18.14$ years). We identified a rather long capture of 82P with minimum $E_{\\rm J} \\sim -2.26$ though it is not an orbiter but instead follows a symmetric trajectory about the $x$-axis, from 2056 February 18--2064 July 26 (duration $\\sim 8.43$ years). There appear to be two TSCs for 74P if we take its initial parameters from the JPL Small-Body Database, from 2025 August 10--2031 May 29 (duration $\\sim 5.80$ years, minimum $E_{\\rm J} \\sim -0.74$) and 2081 January 29--2085 January 29 (duration $\\sim 4.00$ years, minimum $E_{\\rm J} \\sim -0.77$).  Such future TSCs have already been predicted by Carusi et al. (\\cite{carusi85a}, \\cite{carusi85b}) for 74P and 82P and Tancredi et al. (\\cite{tancredi}) for 111P, and our simulations are almost identical in profile with those. Meanwhile, it is unlikely that 1996 R2 and 147P will encounter Jupiter as a long TSC/orbiter for the next 100 years from each initial epoch. No further long TSCs/orbiters were found among the remaining QHCs.  \\subsection{Tidal Effects}  Here we discuss the Jovian tidal force acting on 147P and other captured QHCs, around their perijove passages. Well-known cometary tidal splitting events have occurred twice: 16P in 1886 (Sekanina \\& Yeomans \\cite{sekanina}) and SL9, both passing perijove within the Roche limit for comets, $\\sim 2.7 R_{\\rm J}$. In contrast, the close encounter of 147P with Jupiter around 1952 August 26 (Table~\\ref{tbl:147PTSC}) was at a perijove distance of $14.61^{+2.94}_{-2.61} R_{\\rm J}$, comfortably outside the Roche limit. The radius, $R_{\\rm c}$, of 147P's nucleus was estimated at 0.21 km by Lamy et al. (\\cite{lamy}), the smallest among all their measured cometary nuclei.  The Jovian tidal stress, $\\sigma_{\\rm T}$, acting on small bodies (QHCs here), is given by: \\begin{equation} \\sigma_{\\rm T} \\approx Gm_{\\rm J}\\rho R_{\\rm c}^2/r_{\\rm J}^3, \\end{equation} where $m_{\\rm J}$ is the mass of Jupiter. Compared to the $\\sigma_{\\rm T}$ acting on SL9 at perijove $< 1.5 R_{\\rm J}$ in 1992 (Scotti \\& Melosh \\cite{scotti}; Asphaug \\& Benz \\cite{asphaug94}, \\cite{asphaug}), the $\\sigma_{\\rm T}$ for 147P was extremely small, amounting to $< 0.005 \\%$ if we take the original, pre-encounter $R_{\\rm c}$ of SL9 as $\\sim 1$ km (Scotti \\& Melosh \\cite{scotti}; Asphaug \\& Benz \\cite{asphaug94}) and $\\rho$ of both the comets as being equal. Hence, although our treatment here is approximate, there seems no reason to believe that 147P was affected by Jovian tides. On the other hand, 82P, which passed its perijove at $3.01 R_{\\rm J}$ in 1970, just outside the cometary Roche limit, may have suffered some effect due to the Jovian tides, even though it was not broken up. Taking $R_{\\rm c}$ for 82P as 0.73 km (Lamy et al. \\cite{lamy}), and equal $\\rho$ as above, implies that the $\\sigma_{\\rm T}$ acting on 82P could be up to $7 \\%$ of that for SL9. If 82P is a friable and furthermore a loose rubble-pile object, the heating energy from continual Jovian tides might sufficiently affect the comet's nucleus structure so as to allow H$_2$O ice, if present, to sublimate even though 82P is beyond the usual heliocentric distance for H$_2$O sublimation. If our hypothetical scenario is true, such tidal heating effects could trigger cometary activity on 82P and produce an outburst. In this event, the coma size and consequent brightness would peak several days (or more) after the maximum tidal effects.  Another intriguing point is whether or not other tidal effects, e.g., tidal distortion and tidal torques leading to rotation state changes (Scheeres et al.\\ \\cite{scheeres}), are detectable in the light-curve observations of the captured QHCs. In the physical database of cometary nuclei by Lamy et al. (\\cite{lamy}), we can notice a rather high axis ratio ($a/b > 1.6$) and long rotation period $\\sim 50$ hr in the dataset of 82P, though an interpretation in terms of such tidal effects is still speculative.  \\subsection{Comets or asteroids?}  At any rate, we have considered here only a small subset of the TSC events, focusing on the long TSCs/orbiters, for which there still exist only limited available data. Therefore we know little about such unique and extraordinary astronomical events, which are still full of ambiguities. In addition, we do not know whether QHCs are comets or asteroids. Indeed, several QHCs have sometimes been discovered or recovered as asteroids because of their cometary activity being weak: e.g., 36P (= 1925 QD = 1940 RP), 39P (= 1950 CR), 74P (= 1967 EU = 1978 NA$_6$ = 1981 UH$_{18}$ = 1982 YG$_3$), D/1977 C1 (= 1977 DV$_3$), P/1999 XN$_{120}$, P/2001 YX$_{127}$, and P/2003 CP$_7$. Further observations and research for the QHCs will be necessary to unlock their origin and nature."
    },
    "0808/0808.2661_arXiv.txt": {
        "abstract": "{} {Using new Chandra X-ray observations and existing XMM-Newton X-ray and Hubble far ultraviolet observations, we aim to detect and identify the faint X-ray sources belonging to the Galactic globular cluster \\mbox{\\object{NGC 2808}} in order to understand their role in the evolution of globular clusters.} {We present a Chandra X-ray observation of the Galactic globular cluster \\mbox{\\object{NGC 2808}}. We classify the X-ray sources associated with the cluster by analysing their colours and variability. Previous observations with XMM-Newton and far ultraviolet observations with the Hubble Space Telescope are re-investigated to help identify the Chandra sources associated with the cluster. We compare our results to population synthesis models and observations of other Galactic globular clusters.} {We detect 113 sources, of which 16 fall inside the half-mass radius of \\mbox{\\object{NGC 2808}} and are concentrated towards the cluster core. From statistical analysis, these 16 sources are very likely to be linked to the cluster. We detect short-term (1~day) variability in X-rays for 7 sources, of which 2 fall inside the half-mass radius, and long-term (28~months) variability for 10 further sources, of which 2 fall inside the half-mass radius. Ultraviolet counterparts are found for 8 Chandra sources in the core, of which 2 have good matching probabilities and have ultraviolet properties expected for cataclysmic variables. We find one likely neutron star-quiescent low-mass X-ray binary and 7 cataclysmic variable candidates in the core of \\mbox{\\object{NGC 2808}}. The other 8 sources are cataclysmic variable candidates, but some could possibly be active binaries or millisecond pulsars. We find a possible deficit of X-ray sources compared to \\mbox{\\object{47 Tuc}} which could be related to the metallicity content and the complexity of the evolution of \\mbox{\\object{NGC 2808}}.} {} ",
        "introduction": "Globular clusters (GCs) are old, gravitationally bound stellar systems which can have extremely high stellar densities, especially in their core regions. In such an environment, dynamical interactions between the cluster members are inevitable, leading to a variety of close binary (CB) systems and other exotic stellar objects. The observed overabundance of neutron star (NS) low-mass X-ray binaries (LMXBs) in GCs relative to the Galactic field was explained by the dynamical processes occurring in the dense cores of GCs \\citep{fabian}. In contrast, evolution of a primordial binary into an LMXB in a GC is considered to be much less likely \\citep{VH87}. Observations also support the fact that quiescent LMXBs (qLMXBs) in GCs scale with the cluster encounter rate \\citep{GBW03,Pooley+03}, implying that qLMXBs are formed through dynamical processes in the dense cores. As white dwarfs (WDs) are far more common than NSs, we would then also expect many more close binaries containing an accreting WD primary, i.e. cataclysmic variables (CVs).  The dynamically-formed CBs are expected to be found in the cores of GCs, where the stellar densities are at a maximum. The less dense regions outside the cores might be populated by CBs that evolved from primordial binaries \\citep[e.g.][]{Davies97} which are unlikely to survive in the dense core region. \\citet{HAS07} found that the combined effects of new binary creation and mass segregation exceed the destruction of primordial binaries in the central region of GCs, leading to a marked increase of the binary fraction in the central regions. Thus, we expect the majority of CBs, which are more massive than the mean stellar mass, to be located inside the half-mass radius. Outside the half-mass radius, the primordial binary fraction is well preserved \\citep{HAS07}.  CBs are important for our understanding of GC evolution, since the binding energy of a few, very close binaries can rival that of a modest-sized GC \\citep[e.g.][and references therein]{EHI87,Hut+92,Hut+03}. In the core, binaries are subject to encounters and hard binaries become harder while transferring their energy to passing stars. Thus, CBs can significantly affect the dynamical evolution of the cluster. If there are only a few CBs, thermal processes dominate the cluster evolution, leading to core collapse followed by GC disruption on a timescale shorter than the mean age of GCs, estimated to be $12.9\\pm2.9$~Gyr \\citep{CGCFP00}. In contrast, the presence of many CBs leads to violent interactions, which heat the cluster, delay the core collapse, and promote its expansion. This depends critically on the number of CBs, which is still poorly known.  Finding and studying these systems has proven to be extremely difficult, since the spatial resolution and detection limits of most available telescopes are insufficient for their detection. Only with the improved sensitivity and imaging quality of XMM-Newton and Chandra in the X-ray \\citep[e.g.][]{Webb+04,WWB06,Heinke+03,Heinke+06} and HST in the ultraviolet (UV) to infrared (IR) wavebands \\citep[ and references therein]{GHEM01,Albrow+01,EGHG03,EGHG03b,Knigge+02,Knigge+03} has it become possible to finally detect significant numbers of CB systems in GCs.  The 13 bright X-ray sources found in the $\\sim$150 known Galactic GCs are LMXBs showing type I X-ray bursts \\mbox{\\citep[e.g.][]{1983adsx.conf...41L}}, whereas the faint sources belonging to the clusters are qLMXBs, CVs, active binaries (ABs, generally RS~CVn systems), or millisecond pulsars (MSPs). Multiwavelength studies can be used to identify the faint X-ray sources. For example, qLMXBs are usually identified by their soft blackbody-like or hydrogen atmosphere X-ray spectra \\mbox{\\citep[e.g.][]{GBW03,GBW03b}}, CVs can be confirmed by their blue, variable optical counterpart with hydrogen emission lines in their spectra \\mbox{\\citep[e.g.][]{Webb+04}}, ABs by their main-sequence, variable optical counterparts \\mbox{\\citep[e.g.][]{EGHG03}}, and MSPs by their radio counterpart \\mbox{\\citep[e.g.][]{GHEM01}}.  Here, we present an X-ray study of the massive (\\mbox{$\\sim10^{6} M_{\\sun}$}) GC \\mbox{\\object{NGC 2808}} (\\mbox{$\\alpha = 09^{h} 12^{m} 02^{s}$, $\\delta = $-$64^{\\circ} 51{\\arcmin} 47{\\arcsec}$}). This intermediate metallicity GC \\citep[\\mbox{$\\rm{[Fe/H]} = -1.36$}, ][]{Walker99} lies at a distance of 9.6~kpc and is reddened by $E_{B-V}=0.22\\pm0.01$ \\citep{Harris96}. An absorption column of N${_H = 1.2\\times10^{21}\\mathrm{~cm^{-2}}}$ is derived from the reddening with the relation computed by \\citet{PS95}. The cluster has a very dense and compact core ($0.26\\arcmin$), a half-mass radius of $0.76\\arcmin$, a tidal radius of $15.55\\arcmin$, and a half-mass relaxation time of $1.35\\times10^{9}$~yrs \\citep{Harris96}.  \\mbox{\\object{NGC 2808}} has received considerable attention in the literature and has been observed in the optical in detail as this GC is one of the most extreme examples with an unusual horizontal branch (HB) morphology, as first noted by \\citet{Harris74}. It shows a bimodal HB and one of the longest blue HB tails, the so-called extreme HB (EHB), with prominent gaps between the red HB (RHB), blue HB (BHB) and EHB \\citep[see also ][]{Bedin+00,Carretta+06}. Recently, \\citet{Piotto+07} found that \\mbox{\\object{NGC 2808}}'s main sequence (MS) is separated into three branches, which might be associated with the complex HB morphology and abundance distribution, and might be due to successive rounds of star formation with different helium abundances. \\mbox{\\object{NGC 2808}} is proposed as a good candidate to harbour an intermediate mass black hole (IMBH) in its core, due to its optical luminosity profile and EHB morphology \\citep{Miocchi07}.  The core of \\mbox{\\object{NGC 2808}} has been imaged with the Space Telescope Imaging Spectrograph (STIS) on board the Hubble Space Telescope (HST) in the far-UV (FUV) and the near-UV (NUV). \\citet{Dieball+05} used the data set with an emphasis on the dynamically-formed stellar populations like CVs and blue stragglers (BSs) and young WDs. They found $\\sim$40~WD, $\\sim$60~BS and $\\sim$60~CV candidates in the field of view that covers the core of the cluster. Two of the CV candidates are variable (FUV sources~222 and 397).  \\mbox{\\object{NGC 2808}} has also been observed with XMM-Newton in Feb. 2005. \\citet{Servillat+08-a} found 96~sources in the field of view (equivalent to the tidal radius), of which five fall inside the half-mass radius and are likely to be linked to the cluster. One qLMXB candidate and four CV candidates were discovered in the core of \\mbox{\\object{NGC 2808}}. However, several sources remained unresolved.  In Sect.~\\ref{obs}, we present the new Chandra X-ray data, and then compare them to the XMM-Newton observations (Sect.~\\ref{xmm}). We present HST and XMM-Newton Optical Monitor UV counterparts in Sect.~\\ref{uv}. We finally discuss our results in Sect.~\\ref{disc}. ",
        "conclusions": "We presented Chandra observations of \\mbox{\\object{NGC 2808}} coupled with previous XMM-Newton observations \\citep{Servillat+08-a}, HST FUV observations \\citep{Dieball+05}, VLT and ATCA observations. We have shown that 16 Chandra sources are likely to be linked to the cluster, with possibly a 17th close to the half mass radius. One of these is consistent with the X-ray emission of a qLMXB, confirming the previous detection with XMM-Newton. Two Chandra sources (7 and 14) have FUV counterparts that show emission compatible with a CV. Chandra source 10 is likely to be a CV from its X-ray emission, but no UV counterparts was found to confirm its nature. Another highly variable source (16) in the core is likely to be a CV, as well as two other variable sources (Chandra 3 and XMM-Newton C5). Two other Chandra sources (8 and 11) have optical counterparts compatible with the expected emission of CVs. We have thus identified 7 CV candidates (plus XMM-Newton source C5) and the observations indicate that there may be as many as 15 in the Chandra observations (although some of the faintest may be ABs or MSPs), along with $\\sim30$ CV candidates in the HST UV observations. This significant population of close binaries is likely to play an important role in slowing down the core collapse of this cluster. Compared to the number of X-ray sources detected in \\mbox{\\object{47 Tuc}} and expected from dynamical formation, we found a possible deficit of X-ray sources in \\mbox{\\object{NGC 2808}}. This might indicate a true deficit of CVs, which is possibly linked to the metallicity content and the complexity of the evolution of \\mbox{\\object{NGC 2808}}."
    },
    "0808/0808.1512_arXiv.txt": {
        "abstract": "We examine the temperature-dependent electroweak phase transition in extensions of the Standard Model in which the electroweak symmetry is spontaneously broken via strongly coupled, nearly-conformal dynamics. In particular, we focus on the low energy effective theory used to describe Minimal Walking Technicolor at the phase transition. Using the one-loop effective potential with ring improvement, we identify significant regions of parameter space which yield a sufficiently strong first order transition for electroweak baryogenesis. The composite particle spectrum corresponding to these regions can be produced and studied at the Large Hadron Collider experiment. We note the possible emergence of a second phase transition at lower temperatures. This occurs when the underlying technicolor theory possesses a nontrivial center symmetry. ",
        "introduction": "The experimentally observed baryon asymmetry of the universe may be generated at the electroweak phase transition (EWPT) \\cite{Shaposhnikov:1986jp,Shaposhnikov:1987tw,Shaposhnikov:1987pf, Farrar:1993sp,Farrar:1993hn,Gavela:1993ts,Gavela:1994ds}. {}For the mechanism to be applicable it requires the presence of new physics beyond the Standard Model (SM) \\cite{Nelson:1991ab,Joyce:1994bi,Joyce:1994fu,Joyce:1994zn, Joyce:1994zt,Cline:1995dg}. {An essential condition for electroweak baryogenesis is that the  baryon-violating interactions induced by electroweak sphalerons are sufficiently slow immediately after the phase transition to avoid the destruction of the baryons that have just been created.  This is achieved when the thermal average of the Higgs field evaluated on the ground state,  in the broken phase of the electroweak symmetry, is large enough compared to the critical temperature at the time of the transition (see for example ref.~\\cite{Cline:2006ts} and references therein), \\beq \\phi_c/ T_c   > 1. \\label{cond} \\eeq In the SM, the bound (\\ref{cond}) was believed to be satisfied only for very light Higgs bosons \\cite{Carrington:1991hz,Arnold:1992fb,Arnold:1992rz, Anderson:1991zb,Dine:1992wr}.  However, this was before the mass of the top quark was known.  With $m_t=175$ GeV, nonperturbative studies of the phase transition \\cite{Kajantie:1995kf} show that the bound (\\ref{cond}) cannot be satisfied for {\\it any} value of the Higgs mass. In addition to the difficulties with producing a large enough initial baryon asymmetry, the impossibility of satisfying the sphaleron constraint (\\ref{cond}) in the SM (Standard Model) provides an incentive for seeing whether the situation improves in various extensions of the SM. We refer to \\cite{Cline:2006ts} for a summary of the different attempts in this direction.  In this paper we explore the electroweak phase transition in a model in which the electroweak symmetry is broken dynamically \\cite{Weinberg:1979bn,Susskind:1978ms}. A dynamical origin behind the spontaneous breaking of the electroweak symmetry is a natural extension of the SM. However, electroweak precision data and constraints from flavor changing neutral currents both disfavor an underlying gauge dynamics resembling too closely a scaled-up version of Quantum Chromodynamics (QCD) (see \\cite{Sannino:2008ha,Hill:2002ap,Lane:2002wv} for recent reviews).  Since technicolor models have been less fashionable than supersymmetric models in the last decade, it is worthwhile to review the recent progress that has enhanced their attractiveness from the particle physics perspective.  One area of progress  is in the understanding of the phase diagram \\cite{Sannino:2004qp,Dietrich:2006cm,Ryttov:2007sr,Ryttov:2007cx}, as function of the number of flavors and colors, of any SU(N) non-supersymmetric gauge theory with fermionic matter transforming according to various representations of the underlying gauge group. This has made it possible to provide the first classification of the possible theories one can use to break the electroweak symmetry \\cite{Dietrich:2005jn,Dietrich:2006cm}.  New analytic tools such as the all-order beta function \\cite{Ryttov:2007cx} allow the determination, for the first time, of the anomalous dimension of the mass of the fermions at the nonperturtative infrared fixed point. This information is crucial for walking technicolor models \\cite{Holdom:1984sk,Eichten:1979ah,Holdom:1981rm,Yamawaki:1985zg,Appelquist:1986an,Lane:1989ej}, {\\it  i.e.}, the ones for which the underlying gauge dynamics is nearly conformal.  A key realization that enabled further progress was that  gauge theories with fermions in two-index (symmetric or adjoint) representations of the underlying gauge group have interesting features \\cite{Sannino:2004qp,Dietrich:2005jn,Dietrich:2006cm,Ryttov:2007sr,Ryttov:2007cx}, such as the possibility of the existence of a nonperturbative infrared fixed point for a very low number of flavors \\cite{Sannino:2004qp}, naturally reducing the tension with precision data \\cite{Sannino:2004qp,Dietrich:2005jn,Foadi:2007ue,Foadi:2007se}. These properties make them intriguing candidates for walking technicolor type models \\cite{Sannino:2004qp,Dietrich:2005jn} (related studies can be found in \\cite{Christensen:2005cb}).   In contrast, the naive scaling up of QCD, which is far from conformal, is strongly contradicted by phenomenological constraints \\footnote{The reader will find in Appendix F of \\cite{Sannino:2008ha} a complete account of alternative large N limits one can use to gain information on the spectrum of theories with matter in higher dimensional representation}.  Another important development occurred in  first principle lattice simulations of the minimal walking technicolor theories, carried out in refs.\\ \\cite{Catterall:2008qk, Catterall:2007yx,Shamir:2008pb,DelDebbio:2008zf,DelDebbio:2008wb}. These studies give preliminary support to the analytical arguments that these theories are near or actually already conformal. The case of fermions in the fundamental representation has been investigated in  \\cite{Catterall:2008qk,Appelquist:2007hu,Deuzeman:2008sc}.  On the astrophysical side, technicolor models are capable of providing interesting dark matter candidates, since the  new strong interactions confine techniquarks in technimeson and technibaryon bound states. The spin of the technibaryons depends on the representation according to which the technifermions transform, and the numbers of flavors and colors. The lightest technimeson is short-lived, thus evading BBN constraints, but the lightest technibaryon has typically \\footnote{there may be situations in which the technibaryon is a goldstone boson of an enhanced flavor symmetry} a mass of the order \\begin{equation} m_{TB}  \\sim  1-2\\ {\\rm TeV} \\ . \\end{equation}  Technibaryons are therefore natural dark matter candidates \\cite{Nussinov:1985xr,Barr:1990ca,Gudnason:2006yj}. In fact it is possible to {naturally} understand the observed ratio of the dark to luminous matter mass fraction of the universe if the technibaryon possesses an asymmetry \\cite{Nussinov:1985xr,Barr:1990ca,Gudnason:2006yj}. If the latter is due to a net $B-L$ generated at some high energy scale, then this would be subsequently distributed among {\\em all} electroweak doublets by fermion-number violating processes in the SM at temperatures above the electroweak scale \\cite{Shaposhnikov:1991cu,Kuzmin:1991ft,Shaposhnikov:1991wi}, thus naturally generating a technibaryon asymmetry as well. To avoid experimental constraints the technibaryon should be constructed in such a way as to be a complete singlet under the electroweak interactions \\cite{Barr:1990ca,Dietrich:2006cm} while still having a nearly conformal underlying gauge theory \\cite{Dietrich:2006cm}. In this case it would be hard to detect it in current earth-based experiments  such as CDMS \\cite{Bagnasco:1993st,Gudnason:2006yj,Kouvaris:2008hc,Akerib:2004fq,Akerib:2005kh}. Other possibilities have been envisioned in \\cite{Kouvaris:2007iq,Khlopov:2007ic} and possible astrophysical effects studied in \\cite{Kouvaris:2007ay}.  One can  alternatively obtain dark matter from possible associated new sectors instead of the technicolor sector \\cite{Kainulainen:2006wq}, including those which are not gauged under the electroweak interactions \\cite{Dietrich:2006cm}.   In \\cite{Sannino:2008ha} the reader will find an up-to-date summary of the recent efforts in this direction.   Coming to the main topic of this paper,  the order of the electroweak phase transition (EWPT) depends on the underlying type of strong dynamics and  plays an important role for baryogenesis \\cite{Cline:2002aa,Cline:2006ts}. The technicolor chiral phase transition at finite temperature is mapped onto the electroweak one. Attention must be paid to the way in which the electroweak symmetry is embedded into the global symmetries of the underlying technicolor theory.  An interesting preliminary analysis dedicated to earlier models of technicolor has been performed in \\cite{Kikukawa:2007zk}.  In this work, we wish to investigate the EWPT in a class of realistic and viable technicolor models. An explicit phenomenological realization of walking models consistent with the electroweak precision data   is termed Minimal Walking Technicolor (MWT) \\cite{Foadi:2007ue}. It is based on an SU(2) gauge theory coupled to two flavors of adjoint techniquarks. This model is thought to lie close, in theory space, to theories with nontrivial infrared fixed points \\cite{Sannino:2004qp,Ryttov:2007cx}. Indeed it is possible that it already has such a fixed point itself. In the vicinity of such a zero of the beta function, the coupling constant flows slowly (``walks''). This theory possesses an SU(4) global symmetry. At the LHC one will observe the composite states which are classified according to irreducible representations of the stability group left invariant by the technifermion condensate. The stability group, here, corresponds to the SO(4) symmetry which contains the SU(2)  custodial symmetry of the SM. We choose the natural SM embedding, as detailed in the following section.   In ref.\\ \\cite{Foadi:2007ue} a comprehensive Lagrangian was introduced for this model, taking into account the global symmetries of the underlying gauge theory, the walking dynamics via the modified Weinberg sum rules \\cite{Appelquist:1998xf}, and the constraints coming from precision data \\cite{Foadi:2007se}. The effective theory contains composite scalars and spin-one vectors. Compatibility between the electroweak precision constraints and tree-level unitarity of $WW$-scattering was demonstrated  in \\cite{Foadi:2008ci}.    The study of longitudinal WW scattering unitarity versus precision measurements within the effective Lagrangian approach demonstrated that it is possible to pass the precision tests while simultaneously delaying the onset of unitarity \\cite{Foadi:2008ci}.  In the present work we will use as a template the low energy effective theory developed in  \\cite{Foadi:2007ue}. We start in section \\ref{sect2} by summarizing the basic theory, highlighting the degrees of freedom relevant near the phase transition.  In section \\ref{sect3} the finite-temperature effective potential is then computed at the one-loop order, including the resummation of ring diagrams.  Our analysis is presented in section \\ref{sect4}.   As a preliminary investigation we adopt the high-temperature expansion results for the effective potential.  We then explore the region of the effective theory parameters yielding a first-order phase transition and study its strength. The ratio of the composite Higgs thermal expectation value at the critical temperature divided by the corresponding temperature is determined as function of the parameters of the low energy effective theory. We identify a significant region of parameter space where this ratio is sufficiently large to induce electroweak baryogenesis. {The spectrum of the composite spin-zero states directly associated to these regions can be investigated and the related particles produced at the Large Hadron Collider experiment. In the subsection \\ref{sect4.d} we note  the possible emergence of a second phase transition at lower temperatures, {\\it i.e.}, the confinement/deconfinement one. This transition occurs when the underlying technicolor theory possesses a nontrivial center symmetry.} Several appendices are provided, which give details concerning our analytical results. ",
        "conclusions": ""
    },
    "0808/0808.3384_arXiv.txt": {
        "abstract": "Despite compelling arguments that significant discoveries of physics beyond the standard model are likely to be made at the Large Hadron Collider, it remains possible that this machine will make no such discoveries, or will make no discoveries directly relevant to the dark matter problem. In this article, we study the ability of astrophysical experiments to deduce the nature of dark matter in such a scenario. In most dark matter studies, the relic abundance and detection prospects are evaluated within the context of some specific particle physics model or models ({\\textit e.g.,} supersymmetry).  Here, assuming a single weakly interacting massive particle constitutes the universe's dark matter, we attempt to develop a model-independent approach toward the phenomenology of such particles in the absence of any discoveries at the Large Hadron Collider. In particular, we consider generic fermionic or scalar dark matter particles with a variety of interaction forms, and calculate the corresponding constraints from and sensitivity of direct and indirect detection experiments. The results may provide some guidance in disentangling information from future direct and indirect detection experiments. ",
        "introduction": "The consensus of the astrophysics community is that a large fraction of the universe's mass consists of non-luminous, non-baryonic material, known as dark matter \\cite{Bertone:2004pz}. Although the nature of this substance or substances remains unknown, weakly interacting massive particles (WIMPs) represent a particularly attractive and well motivated class of possibilities. Although the most studied WIMP candidate is the lightest neutralino \\cite{neutralino} in supersymmetric models, many other possibilities have also been proposed, including Kaluza-Klein states in models with universal \\cite{universal,universal2} or warped \\cite{warped} extra dimensions, stable states in Little Higgs theories \\cite{tparity}, and many others.  In each of the above mentioned cases, many new particle species, in addition to the WIMP itself, are expected to lie within the discovery reach of the Large Hadron Collider (LHC), making the task of deducing the nature of the WIMP immeasurably simpler. In supersymmetry, for example, gluinos and squarks are expected to be produced prolifically. By studying the cascades produced in the decays of such particles, the masses of several superparticle masses, including the lightest neutralino, are likely to be determined. If squarks, gluinos, and other additional superpartners are too heavy to be produced, however, the lightest neutralino will also be very difficult to study at the LHC, even if rather light itself. More generally speaking, in the absence of heavier particles with shared quantum numbers, WIMPs will not be easily detected or studied at the LHC. Although an electroweak scale, cold thermal relic particle, if it exists, would almost certainly be produced at the LHC, identifying and characterizing the nature of the WIMP simply from missing energy studies is a daunting, perhaps impossible, task \\cite{Birkedal:2004xn,Feng:2005gj}.  Although the usual list of prospective WIMPs mentioned above contains some very attractive and well-motivated candidates for dark matter, there are certainly many possible forms of dark matter that have not yet been considered. As the first observations of particle dark matter might well come from direct and/or indirect detection experiments, it is possible that these results may be misinterpreted as a result of theoretical bias, anticipating dark matter to have the properties of a neutralino or other often studied candidates. To avoid such confusion, model-independent studies of dark matter phenomenology can play an important role (for previous work in this direction, see Refs.\\ \\cite{Birkedal:2004xn,Kurylov:2003ra,Giuliani:2004uk}).  In this article, rather than consider a WIMP candidate from a specific theoretical model, we study model-independent WIMPs with different combinations of spins and interaction forms with standard model particles.  These interactions are limited only by the requirements of Lorentz invariance and a consequent WIMP abundance consistent with cosmological observations.  For each spin and interaction form, we evaluate the constraints from and prospects for direct and indirect detection of WIMPs in current and future experiments. Although we will be forced to adopt some assumptions in order to make the problems at hand tractable, we attempt to be as general as possible throughout our study. Beyond the starting point that the dark matter is a WIMP in the form of a single species of a cold thermal relic, we adopt only two assumptions:  \\begin{enumerate}  \\item Any new particle species in addition to the WIMP has a mass much larger than the WIMP.  \\item The WIMP interactions with standard model particles are dominated by those of one form (scalar, vector, \\textit{etc.}).  \\end{enumerate}  An implication of the first assumption is that the WIMP's thermal abundance is not affected by resonances or coannihilations. At a later stage of this paper, we will discuss the impact of relaxing these assumptions.  The remainder of this paper is structured as follows. In Sec.\\ \\ref{fermion} we explore the phenomenology of a generic fermionic WIMP, including its annihilation cross section and relic abundance, elastic scattering cross section and direct detection prospects, and indirect detection prospects in the form of a neutrino flux from the Sun and gamma rays and charged particles produced in galactic annihilations. In Sec.\\ \\ref{scalar}, we repeat this exercise for the case of a scalar WIMP. In each of these two sections, we also consider dark matter candidates from specific particle physics frameworks and discuss how they fit into our model-independent analysis. In Sec.\\ \\ref{conclusion} we summarize our results and present our conclusions. ",
        "conclusions": "Even if the Large Hadron Collider does not reveal physics beyond the standard model, a dark matter candidate in the form of a weakly interacting massive particle may still exist. In this article, we have studied how the nature of such a WIMP could be deduced by its signatures in astrophysical experiments. In our analysis, we have taken a general and model-independent approach, considering fermionic or scalar WIMPs with a variety of interaction forms.  In Table \\ref{t1}, we summarize our findings. For each combination of spin and interaction form, we indicate the constraints placed by and the prospects for direct detection experiments, high-energy neutrino telescopes, and indirect detection experiments using gamma-rays or charged particles. Under the column of direct detection, we use the phrases ``strongly excluded,'' ``weakly excluded,'' or ``within near future reach,'' to denote the sensitivity or prospects for each case. By ``strongly excluded,'' we indicate instances in which the effective couplings to quarks, as relevant to elastic scattering with nuclei, must be suppressed by more than a factor of ten relative to the value required to generate thermally the observed dark matter abundance. As we have discussed, such a suppression could result from resonant annihilations, coannihilations, or annihilations to gauge/Higgs bosons. The label ``weakly excluded,'' in contrast, indicates only that the case is excluded if the effective couplings to quarks are not suppressed by such effects. Lastly, the label ``within near-future reach'' indicates an elastic scattering cross section (without suppression) that is within approximately two order of magnitudes of current direct detection limits.  Under the column of neutrino telescopes, we classify each case as either not sensitive or sensitive over a range of WIMP masses (for next generation experiments, such as IceCube).  This evaluation depends on the annihilation products of the WIMP, however, and thus are highly approximate. Under the column of gamma-rays and charged cosmic rays, we simply indicate whether the WIMP's annihilations are or are not suppressed by the square of the WIMP's velocity. If such velocity suppression is present, it is highly unlikely that GLAST, PAMELA or other planned indirect detection experiments will be capable of detecting dark matter.  This leads us to the most obvious and important question: Will the information provided by direct and indirect detection experiments be able to be used to infer the particle nature of the dark matter? Although there are certainly cases in which measurements by these experiments will not lead to an unambiguous identification, there are many scenarios in which a great deal could be learned. For example, if IceCube or another high-energy neutrino telescope were to observe neutrinos from WIMP annihilations in the Sun, we would be able to conclude that the WIMP is likely fermionic,\\footnote{More precisely, we could conclude in this case that the dark matter particle is not a scalar. Vector WIMPs, which we have not studied in this paper, could also potentially generate an observable flux of high-energy neutrinos \\cite{Hooper:2002gs}.} and that it possesses an axial interaction with light quarks. By studying the precise rate observed, one could also potentially determine whether the WIMP's axial interaction played a dominant or only subdominant role in the process of thermal freeze-out. This could be combined with observations from direct detection experiments to further constrain the possible interactions possessed by the WIMP.  As a second possible scenario, imagine that near future direct detection experiments observe a WIMP with a mass of a few hundred GeV and that, shortly afterward, GLAST observes a corresponding gamma-ray signal from WIMPs annihilating in the halo.  From Table \\ref{t1}, we can see that this leads us to only three likely possibilities: the WIMP is either a fermion with vector interactions, a fermion with pseudoscalar-scalar interactions, or a scalar with Yukawa-like scalar interactions.   Although previous studies have shown that dark-matter experiments have the potential to constrain the parameters of supersymmetry \\cite{susydetermine} or even to help identify the theoretical framework from which the dark matter arises \\cite{otherdetermine}, here we have demonstrated that a far more model-independent approach can also be fruitful. In particular, without assuming any particular theoretical framework or model, we have shown that direct and indirect dark matter experiments can be used to considerably constrain the spin and interactions of the dark matter, even in the absence of any discoveries at the LHC.  \\vspace{0.5cm} \\begin{table*}[!ht] \\hspace{0.0cm} \\begin{center} \\begin{large} Fermionic Dark Matter \\end{large} \\end{center} \\begin{ruledtabular} \\begin{tabular} {c c c c c c c c} Interaction &\\vline& Direct Detection & \\,\\vline \\, & Neutrino Telescopes &\\,\\vline \\, &  $\\gamma$-rays, $e^{\\pm}$, $\\bar{p}$ &  \\\\ \\hline \\hline Scalar  &\\vline& Strongly Excluded $M_{\\chi} \\approx 10-100$ GeV & \\vline & Not Sensitive  & \\vline &  Suppressed by $v^2$ &  \\\\ ($G_f \\propto m_f$) &\\vline& Weakly Excluded $M_{\\chi} \\approx 100-200$ GeV &\\vline  &   & \\vline &  & \\\\ &\\vline& Within Near Future Reach $M_{\\chi} \\approx 200-300$ GeV  &\\vline &   & \\vline &  & \\\\ \\hline Scalar  &\\vline& Strongly Excluded $M_{\\chi} \\approx 10$ GeV$-10$ TeV & \\vline & NA & \\vline&  Suppressed by $v^2$ &  \\\\ ($G_f$ {\\rm Universal}) &\\vline&  & \\vline &  & \\vline& &  \\\\ \\hline Pseudoscalar  &\\vline& Not Sensitive & \\vline & Not Sensitive & \\vline&  Unsuppressed &  \\\\ \\hline Vector/Tensor  &\\vline& Strongly Excluded $M_{\\chi}\\approx 10-350$ GeV & \\vline & Not Sensitive & \\vline&  Unsuppressed &  \\\\ &\\vline& Weakly Excluded $M_{\\chi} \\approx 350$ GeV$-2$ TeV & \\vline  &   & \\vline &  & \\\\ \\hline Axial   &\\vline& Not Sensitive & \\vline & Sensitive $M_{\\chi} \\sim 100-500$ GeV & \\vline&  Suppressed by $v^2$ &  \\\\ \\hline Scalar-Pseudoscalar   &\\vline& Not Sensitive & \\vline & Not Sensitive & \\vline&  Suppressed by $v^2$ &  \\\\ \\hline Pseudoscalar-Scalar   &\\vline& Weakly Excluded $M_{\\chi} \\approx 10-180$ GeV & \\vline & Not Sensitive & \\vline&  Unsuppressed &  \\\\ ($G_f \\propto m_f$) &\\vline& Within Near Future Reach $M_{\\chi} \\approx 180-800$ GeV &\\vline  &   & \\vline &  & \\\\ \\hline Vector-Axial   &\\vline& Not Sensitive & \\vline & Not Sensitive & \\vline&  Unsuppressed &  \\\\ \\hline Axial-Vector   &\\vline& Strongly Excluded $M_{\\chi} \\approx 10$ GeV$-2$ TeV & \\vline & Not Sensitive & \\vline&  Unsuppressed &  \\\\ &\\vline& Weakly Excluded $M_{\\chi} \\approx 2-10$ TeV & \\vline  &   & \\vline &  & \\\\  \\end{tabular} \\end{ruledtabular} \\vspace{0.5cm} \\begin{center} \\begin{large} Scalar Dark Matter \\end{large} \\end{center} \\begin{ruledtabular} \\begin{tabular} {c c c c c c c c} Interaction &\\vline& Direct Detection & \\,\\vline \\, & Neutrino Telescopes &\\,\\vline \\, &  $\\gamma$-rays, $e^{\\pm}$, $\\bar{p}$ &  \\\\ \\hline \\hline Scalar  &\\vline& Weakly Excluded $M_{\\phi} \\approx 10-70$ GeV & \\vline & Not Sensitive  & \\vline &  Unsuppressed &  \\\\ ($F_f \\propto m_f$) &\\vline& Within Near Future Reach $M_{\\phi} \\approx 70-200$ GeV  &\\vline  &   & \\vline &  & \\\\ \\hline Scalar  &\\vline& Strongly Excluded $M_{\\phi} \\approx 10$ GeV$-10$ TeV & \\vline & NA & \\vline&  Unsuppressed &  \\\\ ($F_f$ {\\rm Universal}) &\\vline&  & \\vline &  & \\vline& &  \\\\ \\hline Vector  &\\vline& Strongly Excluded $M_{\\phi}\\approx 10$ GeV$-1$ TeV & \\vline & Not Sensitive & \\vline&   Suppressed by $v^2$ &  \\\\ &\\vline& Weakly Excluded $M_{\\phi} \\approx 1-5$ TeV  &\\vline  &   & \\vline & & \\\\ \\hline Scalar-Pseudoscalar  &\\vline& Not Sensitive & \\vline & Not Sensitive  & \\vline &  Unsuppressed &  \\\\ \\hline Vector-Axial  &\\vline& Not Sensitive & \\vline & Not Sensitive  & \\vline &  Suppressed by $v^2$ &  \\\\ \\end{tabular} \\end{ruledtabular} \\caption{A summary of our results, describing the sensitivity and prospects for the direct and indirect detection of dark matter particles in the various cases we have considered. See the text for explanations for the labels used.} \\label{t1} \\end{table*}   The results presented in Table \\ref{t1} rely upon the set of assumptions we have adopted. It must be noted that if dark matter consists of non-thermally produced WIMPs, or of multiple species of particles, our conclusions could be altered considerably. Furthermore, one might worry that the effects of resonances, coannihilations, or annihilations to gauge/Higgs bosons, which we have largely neglected in our analysis, might dramatically change our conclusions. To some degree, however, the impact of such processes are encapsulated in our definitions of ``strongly excluded'' and ``weakly excluded'', as used in Table \\ref{t1}. For example, if a WIMP annihilates largely through a narrow resonance such that twice the mass of the WIMP lies within approximately 5\\% of the exchanged particle, then the effective couplings relevant for elastic scattering can be reduced by a factor of ten without the WIMP being overproduced in the early universe (see Sec.\\ \\ref{resonance}). This mildly (5\\% or less) fine-tuned resonance corresponds to the ``weakly excluded'' label used in the table. Anything labeled ``strongly excluded'' would require the masses to be tuned even more precisely to the resonance value to remain viable. Similarly, if a significant fraction of WIMP annihilations in the early universe proceeded to a combination of gauge or Higgs bosons, or occurred through coannihilations with another species of particle, the elastic scattering cross section could be suppressed. For the scenarios we have labeled as ``strongly excluded'' to have remained hidden from direct detection experiments, however, about $99$\\% or more of the annihilations/coannihilations of WIMPs in the early universe must have occurred through such processes. So although the conclusions we have reached here are not entirely immune to the inclusion of such effects, they are quite robust in all but the most extreme cases.  In conclusion, we find that in the case that the Large Hadron Collider does not discover physics beyond the standard model, astrophysical experiments may still be able to constrain the nature of the dark matter, even without assuming supersymmetry or any other specific particle physics framework. In particular, the spin and interaction forms of dark matter can potentially be identified by combining results from direct detection experiments, neutrino telescopes, and indirect detection experiments using gamma-rays or charged cosmic rays.    \\vspace{0.5cm}"
    },
    "0808/0808.3451_arXiv.txt": {
        "abstract": "% {}{While observational evidence shows that most of the decline in a star's X-ray activity occurs between the age of the Hyades ($\\sim 8 \\times 10^8$ yrs) and that of the Sun, very little is known about the evolution of stellar activity between these ages. To gain information on the typical level of coronal activity at a star's intermediate age, we studied the X-ray emission from stars in the 1.9 Gyr old open cluster NGC~752.} {We analysed a $\\sim$ 140 ks \\chandra\\ observation of NGC~752 and a $\\sim$ 50 ks \\xmm\\ observation of the same cluster. We detected 262 X-ray sources in the \\chandra\\ data and  145 sources in the \\xmm\\ observation.  Around 90\\% of the catalogued cluster members within \\chandra's field-of-view are detected in the X-ray. The X-ray luminosity of all observed cluster members (28 stars) and of  11 cluster member candidates was derived.} {Our data indicate that, at an age of 1.9 Gyr, the typical X-ray luminosity of the cluster members with $M=0.8-1.2~M_{\\sun}$  is $L_{\\rm X} = 1.3 \\times 10^{28}$ \\es, so approximately a factor of 6 less intense than that observed in the younger Hyades. Given that $L_{\\rm X}$ is proportional to the square of a star's rotational rate, the median $L_{\\rm X}$ of NGC~752 is consistent, for $t \\ga 1$ Gyr, to a decaying rate in rotational velocities $v_{\\rm rot} \\propto t^{-\\alpha}$ with $\\alpha \\sim 0.75$, steeper than the Skumanich relation ($\\alpha \\simeq 0.5$) and significantly steeper than observed between the Pleiades and the Hyades (where $\\alpha <0.3$), suggesting that a change in the rotational regimes of the stellar interiors is taking place at $t\\sim 1$ Gyr.}{} ",
        "introduction": "\\label{sec:intro}  Investigations of solar-type stars in the better-studied nearby open clusters have provided a basis for understanding the evolution of stellar activity and  rotation velocity. The early {\\it Einstein} and {\\sc Rosat} studies of open clusters revealed that the mean X-ray luminosity of G, K and early M stars decreases steadily with age from pre-main sequence stars (PMS) through the Pleiades ($\\sim 8\\times 10^7$ yrs), the Hyades ($\\sim 8\\times 10^8$ yrs) and the old disc population (\\citealp{randich2000}; \\citealp{micela2002}).  In main-sequence stars, late-type stellar magnetic activity is regulated principally by rotation (\\citealp{pgr+81}; \\citealp{bds+95}) and its decay with stellar age is attributed to rotational spin-down. Magnetic braking models indicate that surface rotational velocities should decline as $v_{\\rm rot} \\propto t^{-\\alpha}$ with $0.38 < \\alpha < 0.75$, depending on whether the magnetic field geometry is radial or closer to a dipole configuration (\\citealp{kawaler88}). Models considering possible differences in the pre-main sequence disc-locking time, solid-body versus differential rotation in the interior, and the onset of magnetic saturation give decay laws over the range $0.2 < \\alpha< 0.8$ for 1$-$5 Gyr old solar-mass stars (\\citealp{kpb+97}).  Although \\cite{skuma72} found the decay in rotational velocities and the strength of the magnetic activity indicators Ca\\,{\\sc ii} to be consistent with pure magnetic breaking ($\\alpha \\simeq 0.5$) from the age of the Pleiades to that of the Sun, more recent evidence points to a decay rate of the stellar activity which is less steep between the Pleiades and the Hyades and drops significantly faster from the Hyades ages to that of the Sun or field stars (e.g. \\citealp{shb85} for the evolution of transition-region lines; \\citealp{micela02} and \\citealp{pf05} for the evolution of coronal emission). This may indicate a change, between 1 and 5 Gyr, in the rotational regime of stellar interiors.  Helioseismology has shown that in the Sun there is no difference between the angular rotation of the radiative core and that of the convective envelope, while models suggest that a significant radial gradient is present at earlier stages (\\citealp{mb91}, \\citealp{bcm99}, \\citealp{silpin00}).  \\cite{qam+98} reported the surface rotational evolution between the Pleiades and the Hyades to be less steep than the \\cite{skuma72} relation, and interpreted this as evidence of angular momentum transport from the core to the envelope. Determination of the evolution of coronal emission beyond the Hyades age can therefore be an effective probe of the evolution of stellar interiors in older stars.   Although  most of the decline in X-ray activity occurs between the age of the Hyades and that of the Sun, the evolution of stellar activity in solar-type stars older than the Hyades is still poorly known.  For stars older than the Hyades, the data on coronal emission level are based on field stars and the Sun.  While the Sun's age is well constrained at 4.5 Gyr, it is only one star and therefore cannot be assumed to be representative of the behavior of solar-type stars at its current age. The age of field stars are poorly constrained and they represent in general a mixture of stellar populations and chemical composition.  The best approach to explore the stellar activity of solar-type stars at an age intermediate between those of the Hyades and the Sun is the X-ray observation of an old open cluster, which naturally provides a sample of similar aged stars, at the same distance and with similar chemical abundances. This obvious approach is difficult for observational reasons: there are few old clusters in the solar neighbourhood, since many of them are dissipated in few billion years through dynamical evolution, and their intrinsically low X-ray luminosity requires extremely long exposures with current X-ray instrumentation.  The only known, sufficiently nearby, intermediate age, open cluster is NGC~752. At a distance of $\\sim$ 400 pc and with age between 1.7 and 2 Gyr, this is the only open cluster in which one can measure the X-ray luminosity of individual dG stars with exposure times below 200 ks -- unfortunately, though, the cluster is rather dispersed.  We present here the results of the analysis of two X-ray observations of NGC~752: a $\\sim$ 140 ks \\chandra\\ observation and a $\\sim$ 50 ks \\xmm\\ observation.  The analysis focuses on solar-type stars, with mass $0.8 < M < 1.2~M_{\\sun}$. The paper is organised as follows. After a summary, below, of the properties of NGC~752, the observations and data analysis are described in Sect.\\,\\ref{sec:obs}. Results are presented in Sect.\\,\\ref{sec:res} and the implications are discussed in Sect.\\,\\ref{sec:disc}.   \\subsection{NGC~752} \\label{sec:ngc752}  The most comprehensive study of the open cluster NGC~752 to date is the one by \\cite{dlm+94} (hereafter DLM94). These authors collected and systematised a number of existing proper-motion and radial-velocity studies of this cluster, supplemented them with their own measures of radial velocities, and combined them with existing photometric and spectroscopic data. Thanks to this effort, they claim to have obtained an effectively complete census of all probable and possible cluster members down to the observable limit of the unevolved main sequence which, in their case, corresponds to $V \\sim 13.5-14$. The electronic version of their catalogue available from {\\sc Vizier} lists 255 stars of which 157 are classified as probable or possible members of NGC~752.  Combining the spectroscopic and photometric approaches, DLM94 establish the cluster metallicity to be [Fe/H] $= -0.15 \\pm 0.05$ dex. They also derive a reddening value of $E(B-V) = 0.035 \\pm 0.005$, which corresponds to an absorbing column density of $N({\\rm H}) = 2 \\times 10^{20}$ \\cmtwo\\ or $A_V = 0.1$ (\\citealp{cox00}). From intermediate-band photometry and main-sequence fitting, DLM94 derive a distance modulus of $(m-M)=8.25 \\pm 0.10$ mag which corresponds to a distance of $\\sim 430$ pc (assuming $A_V = 0.1$). Comparing the data with theoretical models, they estimate for this cluster an age of $1.9\\pm0.2$ Gyr from classical isochrones and $1.7\\pm0.1$ from isochrones with over-shooting, in good agreement with the results from isochrone fitting of \\cite{mmm93} (1.8 Gyr) and of \\cite{ddg+95} (2.0 Gyr).  Since the first proper-motion study by \\cite{ebbi39}, it has been known that in NGC~752 the unevolved main sequence is deficient in stars relative to the turn-off. This is likely due to mass-segregation effects and the cluster relaxation, which leads to the evaporation of the less massive members.  NGC~752 was observed in the X-ray with {\\sc Rosat} by \\cite{bv96}. They catalogued 49 sources, seven of which were identified with cluster members. ",
        "conclusions": "\\label{sec:disc}  The aim of this paper is to determine the level of stellar activity of solar type stars at an intermediate age between that of the Hyades and those of field stars.  In Fig.\\,\\ref{fig:lxvage}, we compare the luminosities of members of NGC~752 with masses in the range $0.8-1.2~M_{\\sun}$ with the luminosity distributions of the Pleiades, the Hyades and the field stars. These three points are from \\cite{micela02} and refer to stars within the same mass range ($0.8-1.2~M_{\\sun}$). As summarised in Table\\,\\ref{tab:lx}, in NGC~752 we detected in the X-ray 6 cluster members within this mass range (the points in the figure), 5 stars within the \\chandra\\ FOV and one outside (significantly brighter than the others) detected by \\xmm. For stars detected both by \\chandra\\ and \\xmm\\ we used in the figure the \\chandra\\ derived flux estimate. Within the \\chandra\\ FOV there is one additional known member with solar mass: Id 790 in the work by DLM94 (easily identifiable in the CMD in Fig.\\,\\ref{fig:colmag}). The upper limit on the X-ray luminosity of this star is also given in the figure.  In the figure, the crosses give the median value of the distributions (including upper limits) and the size of the vertical bars are determined by the 25\\% and 75\\% quantiles.  In Sect.\\,\\ref{sec:res} we showed that we expect the optical sample by DLM94 to be complete within the \\chandra\\ FOV for $M > 0.8~M_{\\sun}$, thus the sample of the 5 \\chandra\\ detected solar-mass cluster members plus the upper limit for the one undetected source is very likely to include all the cluster members with masses in the $0.8-1.2~M_{\\sun}$ range, in this field.  The median luminosity value for this sample is $1.3 \\times 10^{28}$\\es; the quantiles bars also refer to this sample.  Including the two candidate members with $M \\sim 0.8 M_{\\sun}$ (Ids 136 and 246 -- see Table\\,\\ref{tab:lx}) does not change the median value; including the one \\xmm\\ detected (solar-mass) member outside the \\chandra\\ FOV does not change this value either, although it increases the size of the 75\\% quantile bar.  The error on the distance of the cluster ($430\\pm 20$ pc) is relatively small and does not impact significantly on the position of the median luminosity of NGC~752 in the plot; for the distance range of $410-450$ pc, the value of the median X-ray luminosity stays within the two quantile vertical bars.  The X-ray luminosities of the seven points in NGC~752 have a spread of more than one order of magnitude, confirming the large spread in luminosity of stars with the same age and similar mass already observed in the Pleiades, the Hyades (e.g. \\citealp{ssk95}; \\citealp{msk+96}) and in the sample by \\cite{nb98} studied by \\cite{micela02}. The observed spread in $L_{\\rm X}$ at a fixed age is associated to the spread in rotational periods (\\citealp{pmm+03}), which appears to depend on the coupling between the circumstellar disc and the star in the pre-main sequence phase; in some objects this coupling prevents a young star from spinning up during its PMS contraction, yielding a large spread in the initial angular momentum distribution (e.g. \\citealp{bouvier94}; \\citealp{ch96}). The analysis of $\\sim$ 500 pre-main sequence and recently arrived main-sequence stars by \\cite{hm05} supports this view.  The median X-ray luminosity of 1.3 $ \\times 10^{28}$\\es for solar-type stars in NGC~752 is in good agreement with the decaying trend of X-ray luminosity from the Hyades to the field stars. The value is about 6 times lower than the median value for the Hyades and 6.5 times higher than the median value from field stars, consistent with a steepening of the X-ray luminosity scaling law after the age of the Hyades.  In a study of nine solar-like G-stars with ages ranging from 70 Myr to 9 Gyr, \\cite{ggs97} found $L_{\\rm X} \\propto t^{-\\beta}$ with $\\beta \\sim 1.5$, for stars with ages beyond a few 100 Myr, in agreement with the earlier results by \\cite{msv+87}. The study of a sample of 11 late-type stars in the Chandra Deep Field-North by \\cite{fhm+04} is also consistent with this scaling law, for $ 1 < t < 11$ Gyr, although an excellent fit to their data is found for $\\beta=2.0$. \\cite{pp04} find evidence for a very steep decay of the chromospheric activity between 0.5 and 2 Gyr, their sample at $\\sim 2$ Gyr consisting of seven stars.  Comparison of the median X-ray luminosity of NGC~752 with that of the Hyades is fully consistent with a scaling law with $\\beta \\sim 1.5$. This scaling law is also consistent with the decay in X-ray luminosity from NGC~752 to the field stars, given the age uncertainties of the field stars.  Since $L_{\\rm X} \\approx v_{\\rm rot}^{2}$ (e.g. \\citealp{pgr+81}; \\citealp{pmm+03}), this result implies $\\alpha \\sim 0.75$ for the scaling law of rotational velocities, for $t \\ga 1$ Gyr. This is significantly steeper than found by \\cite{skuma72} and would require a nearly-dipolar magnetic configuration to be explained in terms of magnetic breaking (\\citealp{kawaler88}). The effect of differential rotation in the stellar interior and the onset of magnetic saturation, however, may also play a role (\\citealp{kpb+97}) and it is possible that the coupling efficiency between outer and inner layers of stars weakens with age or that it is mass dependent (\\citealp{barnes03}).  Comparison of rotational velocities in the Pleiades with those in older clusters such as M34 and the Hyades shows that, within the age interval of the Pleiades and Hyades, a star's rotational rate typically decreases less steeply than predicted by a pure magnetic braking law, that is $\\alpha < 0.3$ (\\citealp{qam+98}). The decay of the median $L_{\\rm X}$ of stars with mass $M=0.8-1.2~M_{\\sun}$ from the Pleiades to the Hyades reported by \\cite{micela02} (the points shown in Fig.\\,\\ref{fig:lxvage}) and that reported by \\cite{pf05} for stars with $M=0.9-1.2~M_{\\sun}$ for the same clusters, confirm this result ($\\alpha < 0.25$).  As discussed by \\cite{qam+98}, a value of $\\alpha < 0.3$ could be an indication that angular momentum tapped in the radiative core of slow rotators on the zero age main sequence resurfaces into the convective envelope between the Pleiades and Hyades ages.  We know from helioseismology that there is no gradient between the angular velocities of the core and the envelope in the Sun, thus our data suggest a change in rotation regimes of the stellar interior at $t\\sim $ 1 Gyr.  The shape of the temporal evolution of the X-ray luminosity of solar-mass stars also has an important consequence in the evolution of close-in exoplanets $-$ within 0.5 AU. As shown by \\cite{pml08}, the later the timing of the transition between the two scaling laws of $L_{\\rm X}$, the smaller the fraction of gaseous planets which at 4.5 Gyr retain most of their initial mass. Our result indicates this transition to be at around 1 Gyr.   \\begin{figure}[] \\begin{center} \\leavevmode  \\epsfig{file=10042f10.ps, width=8.0cm}  \\caption{X-ray luminosity of members of NGC~752 with mass between 0.8 and 1.2 $M_{\\sun}$ (filled points for \\chandra\\ detected stars and empty point for the one detected by \\xmm\\ outside \\chandra\\ FOV), compared with median luminosities of stars within the same mass range in the Pleiades, Hyades and a sample of field stars. The triangle indicates the upper limit for star Id 790 from DLM94. Error bars indicate the median value with the size of the vertical bar indicating the 25\\% and 75\\% quantiles of the distribution (see text). The line at age 4.5 Gyr connects the minimum and maximum of the Sun.}  \\label{fig:lxvage} \\end{center} \\end{figure}"
    },
    "0808/0808.4047_arXiv.txt": {
        "abstract": "This Letter is to investigate the physics of a newly discovered phenomenon --- contracting flare loops in the early phase of solar flares. In classical flare models, which were constructed based on the phenomenon of expansion of flare loops, an energy releasing site is put above flare loops. These models can predict that there is a vertical temperature gradient in the top of flare loops due to heat conduction and cooling effects. Therefore, the centroid of an X-ray looptop source at higher energy bands will be higher in altitude, for which we can define as normal temperature distribution. With observations made by {\\it RHESSI}, we analyzed 10 M- or X-class flares (9 limb flares). For all these flares, the movement of looptop sources shows an obvious U-shaped trajectory, which we take as the signature of contraction-to-expansion of flare loops. We find that, for all these flares, normal temperature distribution does exist, but only along the path of expansion. The temperature distribution along the path of contraction is abnormal, showing no spatial order at all. The result suggests that magnetic reconnection processes in the contraction and expansion phases of these solar flares are different. ",
        "introduction": "It is widely accepted that solar flare energy comes from sudden release of free magnetic energy via magnetic reconnection. In commonly adopted flare models, the ever-ascending Y-type reconnection point in the solar corona results in expanding flare loops and separation motion of flare foot points (FPs) (Kopp \\& Pneuman 1976). The expansion of flare ribbons and flare loops is an important signature of progressive magnetic reconnection in the corona. However, the contraction of flare loops in the early phase of flares may suggest a different reconnection scenario.  The signature for the contraction of flare loops has been reported by several authors (Sui et al. 2003, 2004; Li \\& Gan 2005; Liu et al. 2004; Veronig et al. 2006; Ji et al. 2004b, 2006, 2007, 2008). Contraction motion includes two aspects: the converging motion of FPs and the correlated downward motion of a LT source, such as the M1.1 flare of 2004 November 1 (Ji et al. 2006). The contraction picture is different from the shrinkage of flare loops with rooted FPs being fixed or still in expansion. From our experience, the similar event like the M1.1 flare is rare, since HXR FPs are usually missing during the initial phase of a flare, such as the X3.9 flare of 2003 November 3 (Veronig et al. 2006). To investigate the converging motion of FPs, the most preferable wavelength is H$_\\alpha$ blue wing, at which H$_\\alpha$ emission is believed to be caused by nonthermal electrons (Canfield et al. 1984).  The classical flare model puts the energy releasing site above flare loops, from which we can expect that a higher temperature source will be located above a lower temperature source due to heat conduction and cooling effects. In this Letter, we will name this distribution as normal temperature distribution (NTD). For the X10 flare of 2003 October 29, Ji et al. (2008) reported that a HXR sigmoid structure contracts during the impulsive phase of the flare. The contraction is the result of reconnection between two highly-sheared flux ropes. For magnetic reconnection between flux ropes, the existence of NTD is not required, since the energy releases inside flux ropes. Thus, temperature distribution along the altitude is an important factor for testing the reconnection scenarios.  Based on above thinking, we re-analyzed the M2.1 flare of 2002 September 9. The flare is a well-observed sample showing the early converging and subsequent separation motion of the flare kernels. We found that the motion of the HXR LT source is well correlated with that of the flare kernels. Notably, the NTD occurs only during the expansion period of flare loops. During the contraction period, higher and lower temperature structures are mixed together. To find supporting evidences for this abnormal temperature distribution during the contraction period, we surveyed nearly all M-class and X-class limb flares from 2002 to 2005, which were well-observed by {\\it RHESSI} (Lin et al. 2002), we found at least 9 events showing the phenomenon.  In \\S 2 we present the result of the flare of 2002 September 9, in \\S 3 we give the other 9 supporting events. Discussions and conclusions are briefly given in \\S 4. ",
        "conclusions": "In this Letter, we report the finding of the early abnormal temperature distribution of flares' X-ray LT sources during contraction phase of ten solar flares. In a 2D framework for flare models, an energy releasing site is assumed to be above flare loops. Following the scenario, we can expect that a LT source at higher energies will be higher in altitude due to heat conduction and cooling effects. The results from the ten flares show that this expectation can be met, but only during the expansion phase of solar flares. During the contraction period, however, high energy and low energy LT sources are mixed, showing a kind of abnormal temperature distribution.  The physics for the contraction of flare loops is still not well-understood. From above ten events, we have seen that the turning points in the trajectories of the LT sources occur near flare peak times. Therefore, the contraction is related to magnetic reconnection in the impulsive phase of solar flares. Ji et al. (2008) reported that, during the contraction period of the X10 flare of 2003 October 29, HXR emissions at all energies share the similar sigmoidal configuration. The contraction corresponds to the shrinkage of the HXR sigmoid, which is the result of magnetic reconnection between highly-sheared flux ropes. We may propose that the abnormal temperature distribution is associated with magnetic reconnection between highly-sheared flux ropes. On the other hand, the downward and upward phases of the LT motion may reflect two regimes of reconnection: bursting and no-bursting (Karlick\\'{y} 2008: private communication). The abnormal temperature structure can be explained by a presence of small plasmoids in the current sheet just above contracting arcade. The plasmoids are formed in the bursting regime of the reconnection and move also downwards where they interact with the arcade. This interaction represents additional reconnection that changes the normal temperature structure (B\\'{a}rta et al. 2008; Koloma\\'{n}ski \\& Karlick\\'{y} 2007). In the phase of the upward LT motion the reconnection take place without these plasmoids (no-bursting regime of the reconnection). Therefore the temperature is in agreement with the prediction made by standard flare models."
    },
    "0808/0808.3983_arXiv.txt": {
        "abstract": "We present multi-color optical observations of long-duration $\\gamma$-ray bursts (GRBs) made over a three year period with the robotic Palomar 60\\,inch telescope (P60).  Our sample consists of all 29 events discovered by \\Swift\\ for which P60 began observations less than one hour after the burst trigger.  We were able to recover $80\\%$ of the optical afterglows from this prompt sample, and we attribute this high efficiency to our red coverage.  Like \\citet{mmk+08}, we find that a significant fraction ($\\approx 50\\%$) of \\Swift\\ events show a suppression of the optical flux with regards to the X-ray emission (so-called ``dark'' bursts).  Our multi-color photometry demonstrates this is likely due in large part to extinction in the host galaxy.  We argue that previous studies, by selecting only the brightest and best-sampled optical afterglows, have significantly underestimated the amount of dust present in typical GRB environments. ",
        "introduction": "\\label{sec:intro} The launch of the \\Swift\\ $\\gamma$-Ray Burst (GRB) Explorer \\citep{gcg+04} in 2004 November has ushered in a new era in the study of GRB afterglows.  \\Swift\\ offers a unique combination of event rate ($\\sim 100$\\,yr$^{-1}$; almost an order of magnitude increase over previous missions) and precise localization ($\\sim 3\\arcmin$ radius error circles are distributed seconds after the burst, and refined to $\\sim 3\\arcsec$ minutes later). The on-board X-ray Telescope (XRT; \\citealt{bhn+05}) and UV-Optical Telescope (UVOT; \\citealt{rkm+05}), together with the rapid relay of these precise localizations to ground-based observers, has enabled an unprecedented glimpse into the time period immediately following the prompt emission over a broad frequency range.  Observations of X-ray afterglows with the XRT have generated particular interest in recent years.  In the pre-\\Swift\\ era, X-ray observations were limited to hours or days after the prompt emission, and were often poorly sampled compared with the optical and radio bandpasses.  Routine XRT observations of \\Swift\\ GRBs beginning at early times have revealed a central engine capable of injecting energy into the forward shock at times well beyond the duration of the prompt emission (e.g., \\citealt{brf+05,zfd+06}). This discovery has had a profound effect on our understanding of progenitor models.  While the X-ray afterglow is currently a well-explored phase space, comparatively few analogous studies have been performed in the optical bandpass.  \\citet{bkf+05} first suggested that \\Swift\\ optical afterglows were 1.8\\,mag fainter in the $R$-band than pre-\\Swift\\ events (at a common epoch of 12\\,hours after the burst).  Likewise, \\citet{rsf+06} found that only 6 of the first 19 \\Swift\\ bursts with prompt ($\\Delta t \\lesssim 100$\\,s) UVOT coverage yielded optical afterglow detections.  Since then, explaining the faintness of \\Swift\\ optical afterglows has remained one of the outstanding questions in the field.  One clear contributor is distance: the median redshift of \\Swift\\ events ($\\langle z_{Swift} \\rangle \\approx 2.0$)\\footnote{Calculated from J.~Greiner's compilation at http://www.mpe.mpg.de/\\~{}jcg/grbgen.html.} is significantly larger than the pre-\\Swift\\ sample ($\\langle z_{\\mathrm{pre-}Swift} \\rangle = 1.1$; \\citealt{bkf+05,jlf+06}).  In a comprehensive literature-based study of the brightest, best-studied \\Swift\\ afterglows, \\citet{kkz+07} find properties broadly similar to pre-\\Swift\\ events, after applying a cosmological k-correction.  On the other hand, \\citet{mmk+08} have recently presented a sample of 63 GRBs observed in the optical ($r^{\\prime}$-band) with the robotic 2\\,m Liverpool Telescope and Faulkes Telescopes (North and South).  The selection criteria for including a burst in their sample is never explicitly stated, and several non-\\Swift\\ bursts are included, making a direct comparison with the results of \\citet{kkz+07} difficult.  However, \\citet{mmk+08} do not exclude the significant fraction of events without optical detections from their analysis, providing a more unbiased look at optical afterglow properties. By measuring the ratio of optical to X-ray flux at a common time, these authors find that roughly half of the GRBs in their sample exhibit a relative suppression of the optical flux inconsistent with our standard picture of afterglow emission (e.g., \\citealt{spn98}), so-called ``dark'' bursts \\citep{jhf+04}.  This finding suggests that distance alone cannot explain the faintness of \\Swift\\ optical afterglows.  Several other possibilities have been suggested to explain optically dark GRB afterglows.  Undoubtedly some GRBs, like GRB\\,050904 \\citep{hnr+06,kka+06}, originate from such large redshifts ($z \\gtrsim 6$) that Ly-$\\alpha$ absorption in the inter-galactic medium (IGM) completely suppresses the optical flux \\citep{lr00}.  Alternatively, late-time energy injection from the central engine, manifested as bright X-ray flares and/or extended periods of shallow decay, may be artificially increasing the X-ray flux, leading to spurious claims of optically dark GRBs \\citep{mmk+08}.  One final possibility is extinction native to the GRB host galaxy.  As a population, long-duration GRB host galaxies exhibit extremely large neutral H column densities (e.g., \\citealt{hmg+03,bpc+06}), typically falling at $\\log N_{H} > 20.3$\\,cm$^{-2}$ (so-called Damped Ly-$\\alpha$, or DLA systems; \\citealt{wgp05}).  And within their hosts, GRBs trace the blue light from hot young stars in the disk even more closely than core-collapse supernovae \\citep{bkd02,fls+06}.  Both findings are consistent with the observed association between long-duration GRBs and massive star death (e.g., \\citealt{wb06}).  In spite of these expectations, relatively few GRB afterglows to date exhibit signs of large host galaxy extinction (e.g., \\citealt{cbm+07,rvw+07,tlr+08}). \\citet{kkz+07} find only a modest amount of dust ($\\langle A_{V} \\rangle = 0.20$\\,mag) for the 15 events in their ``golden'' sample, an identical value found from an analogous study of pre-\\Swift\\ afterglows \\citep{kkz06}. The primary drawback of such studies, however, is the large and uncertain role of selection effects: by including only the brightest, best-sampled optical afterglows, \\citet{kkz+07} may be preferentially selecting those events in low-extinction environments.  Understanding these selection effects is one of the primary goals of this work.  The Palomar 60\\,inch telescope (P60) is a robotic, queue-scheduled facility dedicated to rapid-response observations of GRBs and other transient events \\citep{cfm+06}.  With a response time of $\\Delta t \\lesssim 3$\\,min and a limiting magnitude of $R \\gtrsim 20.5$ (60\\,s exposure), the P60 aperture is well suited to detect most \\Swift\\ optical afterglows \\citep{as07}. In addition, with a broadband filter wheel providing coverage from the near-UV to the near-IR, P60 can also provide multi-color data on the afterglow evolution.  In this work, we present the P60-\\Swift\\ Early Optical Afterglow sample: 29 unambiguously long-duration GRBs detected by the \\Swift\\ Burst Alert Telescope (BAT; \\citealt{bbc+05}) with P60 observations beginning at most one hour after the burst trigger time.  This sample offers two distinct advantages over previous efforts to understand the optical afterglow emission from GRBs.  First and foremost, our study enforces a strict selection criterion independent of the optical afterglow properties, and therefore will allow us to study the properties of the \\Swift\\ population in a relatively unbiased manner.  Secondly, nearly all events contain multi-color ($\\gp\\,\\Rc\\,\\ip\\,\\zp$) observations that allow us to evaluate the importance of host galaxy extinction for a fraction of our sample.  Altogether, we aim to discriminate between the competing hypotheses proffered to explain dark GRB afterglows in the \\Swift\\ era.  Throughout this work, we adopt a standard $\\Lambda$CDM cosmology with $h_{0}$ = 0.71\\,km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\mathrm{m}} = 0.27$, and $\\Omega_{\\Lambda} = 1 - \\Omega_{\\mathrm{m}} = 0.73$ \\citep{sbd+07}.  We define the flux density power-law temporal and spectral decay indices $\\alpha$ and $\\beta$ as $f_{\\nu} \\propto t^{-\\alpha} \\nu^{-\\beta}$ (e.g., \\citealt{spn98}). All errors quoted are 1 $\\sigma$ (i.e., $68\\%$) confidence intervals unless otherwise noted. ",
        "conclusions": "\\label{sec:discussion}   \\subsection{Anomalous P60 Detection Efficiency} \\label{sec:deteff:disc} We have demonstrated in \\S~\\ref{sec:deteff:obs} that P60 was able to detect optical afterglow emission from a large fraction ($\\sim 80\\%$) of events for which observations began within an hour of the burst trigger.  While the 1.5\\,m aperture is relatively large for a robotic facility, it would be nonetheless informative to understand systematic effects that affect our afterglow recovery rate.  The ultimate goal, of course, is to better inform future GRB follow-up campaigns.  The first lesson from this campaign is the importance of observing in redder filters.  We have shown in \\S~\\ref{sec:dark:obs} that typical \\Swift\\ events suffer from a non-negligible amount of host galaxy extinction (Tab.~\\ref{tab:early}). Coupled with the additional effect of Ly-$\\alpha$ absorption in the IGM from a median redshift of $\\langle z \\rangle \\approx 2$, it is clear that a large fraction of the low UVOT detection efficiency is caused by its blue observing bandpass.  The P60 automated follow-up sequence, consisting of alternating exposures in the \\Rc, \\ip, and \\zp\\ filters, while initially designed for identification of candidate high-$z$ events, is actually well-suited to maximize afterglow detection rates.  The large fraction of P60-detected bursts with spectroscopic redshifts, on the other hand, is almost certainly an artifact of the unequal longitudinal distribution of large optical telescopes.  With the exception of the South African Large Telescope (SALT), all optical telescopes with apertures larger than 8\\,m fall within six time zones (UT-4 to UT-10).  It is not entirely surprising then, that so many promptly discovered P60 optical afterglows have spectroscopic redshifts from immediate follow-up with the largest optical facilities.  While building the largest optical facilities is often prohibitively expensive for all but the largest collaborations, 1\\,m class facilities are much more feasible, both in terms of cost and construction time scale.  We wish here to echo the thoughts of many previous GRB observers (e.g., \\citealt{as07}) that future automated facilities be built at longitudes (and latitudes) not covered by current facilities.  NIR coverage is particularly crucial to detect the most extinguished events and provide tighter constraints on the afterglow SED and hence host galaxy extinction.  A longitudinally spaced ring of 1\\,m class facilities, as for example envisioned by the Las Cumbres Observatory Global Telescope\\footnote{See http://lcogt.net.} is well positioned in the future to recover the vast majority of GRB optical afterglows, assuming the follow-up is done in the reddest filters possible. Such coverage will be particularly important as we transition into the \\textit{Fermi} era, with its significantly decreased rate of precise GRB localizations.  \\subsection{Re-visiting Dark Bursts} \\label{sec:darkbursts:disc} We now turn our attention to the issue of dark bursts in the \\Swift\\ era. In \\S~\\ref{sec:dark:obs}, we demonstrated that a large fraction ($\\approx 50\\%$) of \\Swift\\ afterglows showed suppressed emission in the optical bandpass (relative to the X-ray), that was due in large part to extinction in the host galaxy.  Given the natural expectation that GRBs, since they are associated with massive stars, should form in relatively dusty environments, we wish to understand why our study of \\Swift\\ events yields such a dramatically different dark burst fraction than previous work on pre-\\Swift\\ GRBs \\citep{jhf+04}.  We believe selection effects are one large cause of this discrepancy. It is clear that previous studies of GRB host galaxy extinction, by selecting the brightest and best-sampled events, provide a strongly biased view.  Many, if not most, GRB hosts, appear to suffer from a significant amount of dust extinction ($A_{V} \\gtrsim 0.5$).  Even the study of \\citet{jhf+04}, though it included \\textit{all} pre-\\Swift\\ GRBs with an X-ray afterglow, could be biased towards unextinguished events as well.  Before \\Swift, target-of-opportunity X-ray observations often required the accurate localization provided by an optical (or radio) afterglow.  Thus those events with the brightest optical afterglows (assumed to have on average smaller extinction) were more likely to be observed in the X-ray, biasing the optical-to-X-ray spectral index to larger values of $\\beta_{OX}$.  Another, more subtle, effect, may also cause \\Swift\\ afterglows to appear darker than pre-\\Swift\\ afterglows, independent of host galaxy extinction. Because \\Swift\\ is a more sensitive instrument, it detects GRBs at a higher average redshift than any previous mission.  Consider a host frame extinction of $A_{V} = 0.1$\\,mag.  At $z = 1$, typical for pre-\\Swift\\ events, the observed \\Rc\\ filter corresponds to roughly to rest-frame $U$-band, and so an extinction of 0.17\\,mag (assuming a Milky Way-like extinction curve).  On the other hand, at $z = 3$, the observed \\Rc-band corresponds to a rest frame wavelength of $\\lambda = 1647$\\,\\AA.  So at high redshift, the same amount of dust will produce nearly twice as much extinction in the observed bandpass. Solely because of redshifts effects, similar environments will produce different observed spectral slopes.  This effect is exacerbated by the nature of dust grains in most GRB host galaxies, as the SMC extinction curve does not show the pronounced 2175\\,\\AA\\ bump seen from the Milky Way \\citep{p92}.  If GRBs do trace the cosmic star formation rate, our results suggest a significant fraction of star formation occurs in highly obscured environments. \\citet{kkz06} found a weak correlation between host reddening and sub-mm flux, and we believe a sensitive mid-infrared or sub-mm survey of GRB host galaxies would be an important confirmation of our results.  However, instead of focusing on the brightest, best studied afterglows, as has often been done in the past (e.g., \\citealt{tbb+04,mhc+08}), we instead suggest a survey of the host galaxies of the optically darkest GRB afterglows to see if these events really do exhibit signs of obscured star formation."
    },
    "0808/0808.2917_arXiv.txt": {
        "abstract": "Gamma-ray bursts (GRBs) are cosmologically distributed, very energetic and very transient sources detected in the $\\gamma$-ray domain. The identification of their $x$-ray and optical afterglows allowed so far the redshift measurement of  \\ngrbs\\ events, from $z=0.01$ to $z=6.29$. For about half of them, we have some knowledge of the properties of the parent galaxy.  At high redshift ($z>2$), absorption lines in the afterglow spectra give information on the cold interstellar medium in the host. At low redshift ($z<1.0$) multi-band optical-NIR photometry and integrated spectroscopy reveal the GRB host general properties. A redshift evolution of metallicity is not noticeable in the whole sample. The typical value is a few times lower than solar. The mean host stellar mass is similar to that of the Large Magellanic Cloud, but the mean star formation rate is five times higher. GRBs are discovered with $\\gamma$-ray, not optical or NIR, instruments. Their hosts do not suffer from the same selection biases of typical galaxy surveys. Therefore, they might represent a fair sample of the most common galaxies that existed in the past history of the universe, and can be used to better understand galaxy formation and evolution. ",
        "introduction": "\\begin{figure} \\centerline{\\includegraphics[scale = .4]{fig1.eps}} \\begin{center} \\caption{Histogram of GRBs with measured redshift (empty histogram \\ngrbs\\ objects). The high-$z$ and low-$z$ filled histograms are the subsample of GRBs studied using optical afterglow spectra (GRB-DLAs) and multiband photometry of the host galaxies (GRB hosts), respectively.} \\label{z_hist} \\end{center} \\end{figure}  More than 40 years have passed since the discovery of the first gamma-ray burst (GRB), the most energetic explosions in the universe, by the US military satellite Vela. It took 30 years to demonstrate their cosmological origin (\\cite[Metzger et al.\\ 1997]{metzger97}). Their $\\gamma$-ray energies (typically $10^{51}$ ergs, emitted in less than a couple of minutes) emerge from a collimated jet in a core-collapse supernovae, or the merger of two compact objects (neutron stars or black holes).  Due to the highly transient nature of GRBs, their redshift, measured from the optical afterglow or the host galaxy, is known today for \\ngrbs\\ objects only (Figure~\\ref{z_hist}). Although GRBs are very rare (a rate of 1 event every $10^5$ years in a galaxy is estimated, after correcting for the jet opening angle), they are so energetic that a few events are expected to be detectable from Earth every day. As shown by the discovery of high redshift events (the highest ever being GRB~050904 at $z=6.3$; Kawai et al.\\ 2006), GRBs offer the opportunity to explore the most remote universe, under extreme conditions, hard to observe using traditional tools.  The first scientific paper on GRBs dates from 6 years after the Vela discovery (Klebesadel, Strong \\& Olson 1973). In 1975, already 100 different theories where proposed, and today more than 5000 refereed papers on GRBs have been published.  Before the 1997 discovery, theories on the energetic emission ranked from the impact of comets onto neutron stars (Harwit \\& Salpeter 1973) to collisions of chunks of antimatter with normal stars (Sofia \\& van Horn 1974). However, the most likely hypothesis was already postulated years earlier, when nothing was really known about GRBs: Colgate (1968) predicted $\\gamma$-ray emission from supernovae in distant galaxies.  The curse and blessing of GRBs is their fast fading. The light curve is extremely steep and therefore very hard to catch. The emitted energy is so immense that it cannot last long. However, it allows us to see a hidden universe. One extreme case is the recent event GRB~080319B at $z=0.937$ (Bloom et al.\\ 2008), visible for a short time by naked eye (pick optical magnitude $m=5.6$). This GRB was already hard to observe spectroscopically with the largest telescopes 7 hours after its discovery.  \\begin{table} \\begin{center} \\caption{Overview of host diagnostics with GRB observations.} \\label{t1} {\\scriptsize \\begin{tabular}{|l|c|c|c|c|}\\hline & {\\bf GRB-DLA} & {\\bf GRB host} & {\\bf GRB host} \\\\ {\\bf Diagnosis}    &  {\\bf (AG spectroscopy)}  & {\\bf (Int. spectroscopy)} & {\\bf (Photometry)} \\\\ \\hline SFR & $\\times$ & $\\surd\\surd$ & $\\surd$  \\\\  \\hline Metallicity & $\\surd\\surd\\surd$ & $\\surd\\surd$ & $\\times$  \\\\ \\hline Dust extinction & $\\surd$ & $\\surd\\surd$& $\\times$ \\\\ \\hline Dust depletion & $\\surd\\surd\\surd$ & $\\times$ & $\\times$ \\\\ \\hline Stellar mass & $\\times$ & $\\times$ & $\\surd\\surd\\surd$   \\\\ \\hline Age & $\\times$ & $\\surd$ & $\\surd$ \\\\ \\hline \\end{tabular} } \\end{center} \\end{table} ",
        "conclusions": "The main, still unsolved, question about GRB hosts is whether they represent a fair sample of the whole star-forming galaxy population, or they are a distinct population of galaxies. GRB hosts are generally small, star-forming galaxies detected at any redshift, up to $z=6.3$. The mean stellar mass (measured mainly for the low-$z$ subsample) is a few times above $10^9$ M$_\\odot$. That is the stellar mass of the Large Magellanic Cloud. The mean SFR is 5 times higher than the LMC (Savaglio et al.\\ 2008). Metallicities of the hosts, measured from absorption lines in the afterglow spectra at $z>2$ or from emission lines in the host spectra at $z<1.0$, do not indicate a clear redshift evolution (Figure~\\ref{fig4}), with values mainly between solar and 1/10 solar.  This suggests that the chemical enrichment is not a parameter characterizing the GRB host population.  It is clear that small star-forming galaxies are the most common galaxies that existed in the entire history of the universe. For instance, the star formation history of the universe, studied for different galaxy stellar-mass bins, indicates that small galaxies dominated for redshift $z<2$, where massive galaxies experienced a fast decline (Juneau et al.\\ 2005). GRB hosts are low stellar-mass star-forming galaxies and can provide a very efficient way to understand galaxy formation and evolution in the most active phase of the universe. Future missions will give the possibility to study them in much larger quantities.  \\begin{acknowledgement} We thank the conference organizers for the outstanding organization. The author thanks Karl Glazebrook and Damien Le Borgne for the very fruitful and long-lasting collaboration.  \\end{acknowledgement}"
    },
    "0808/0808.1639_arXiv.txt": {
        "abstract": "Orbital periods in AM Her stars (\\emph{polars}) are synchronized with spin periods of white dwarf by its high magnetic field. Since the last study of $P_{orb}$ distribution of these systems, the number of known objects of such type has more than doubled. This challenged us to compile a new updated catalogue of cataclysmic variables with highly magnetic white dwarfs (polars) and to study their $P_{orb}$ distribution. In this paper we also discus if \"spike\" is reliable feature in the distribution. (\"Spike\" is a concentration of polars in the distribution of their orbital periods near $P_{orb}$ = 114 min and was previously discussed by Ritter \\& Kolb (1992) and Shahbaz \\& Wood (1996).) ",
        "introduction": "AM Her stars or Polars - is a subtype of cataclysmic variables, where binary stellar system contains highly magnetic white dwarf. The strength of magnetic field allows to control accretion in such system (mass is transferring directly on the white dwarf pole, without formation of accretion disc). As a distinct from intermediate polars, AM Her stars orbital period is synchronized with white dwarf spin period.  Last study of $P_{orb}$ distribution were made by Shahbaz \\& Wood (1996) in 1996 (just 43 systems were known by that time). They have noted that discovery of one more AM Her-type star with the orbital period outside of the spike will decrease its significance below 99\\% level. Since that time the number of known systems has more then doubled. This incentivised us to compile a new updated catalogue of cataclysmic variables with highly magnetic white dwarfs and to study distribution of their orbital periods and to re-calculate significance of spike in the similar way, as it was done in previous work. ",
        "conclusions": "\\begin{itemize} \\item Most of known polars have orbital periods below \"period gap\" (50 of 91 known polars). \\item We obtained a grater statistic (91 systems in total) and were impelled to make the spike ranges some broader: $\\Delta P_{spike}$ = 4 min ($\\Delta P_{spike}$ = 2 min in Ritter \\& Kolb (1992) - 17 polars in total were known, $\\Delta P_{spike}$ = 3 min in Shahbaz \\& Wood (1996) - 43 polars were known). \\item Though spike remains to be significant feature, its significance has decreased and now its value is about $99.6 \\%$, instead of  the 99.9\\%, as it was in the work of T.Shahbaz \\& Janet H.Wood (1996). \\item The observed number of systems in the gap is inconsistent with the period distribution being uniform at the 96.3\\% level.  So we have a lowering of this value, just like as for the spike significance.  \\end{itemize}"
    },
    "0808/0808.1580_arXiv.txt": {
        "abstract": "We generalize the Bekenstein-Sandvik-Barrow-Magueijo (BSBM) model for the variation of the fine structure 'constant', $\\alpha ,$ to include an exponential or inverse power-law self-potential for the scalar field $% \\varphi $ which drives the time variation of $\\alpha $, and consider the dynamics of $\\varphi $ in such models. We find solutions for the evolution of $\\varphi $ or $\\alpha $ in matter-, radiation- and dark-energy-dominated cosmic eras. In general, the evolution of $\\varphi $ is well determined solely by either the self-potential or the coupling to matter, depending on the model parameters. The results are general and applicable to other models where the evolution of a scalar field is governed by a matter coupling and a self-potential. We find that the existing astronomical data stringently constrains the possible evolution of $\\alpha $ between redshifts $z\\simeq 1-3.5$ and the present, and this leads to very strong limit on the allowed deviation of the potential from that of a pure cosmological constant. ",
        "introduction": "\\label{sect:Introduction}  For the first time there is a body of detailed astronomical evidence consistent with the time variation of a traditional constant of Nature. The observational programme of Webb \\emph{et al}. \\cite{webb1,webb2} has completed detailed analyses of 128 Keck-HIRES quasar absorption line systems at redshifts $0.5<z<3$ using the many-multiplet method to compare separations between line separations affected by special-relativistic effects and found evidence consistent with the fine structure constant at redshift $z$, $\\alpha (z)$, having been \\textit{smaller} in the past, at redshifts $z\\simeq 1-3.5.$ The shift in the value of $\\alpha $ between its value $\\alpha (z)$ at redshift $z$ and its present-day value $\\alpha (0)$, for all the data is given provisionally by \\begin{equation*} \\Delta \\alpha /\\alpha \\equiv \\lbrack \\alpha (z)-\\alpha (0)]/\\alpha (0)=(-0.57\\pm 0.10)\\times 10^{-5}. \\end{equation*} Subsequent reduction of another data set of 23 VLT-UVES quasar absorption systems $0.4\\leq z\\leq 2.3$ by Chand \\emph{et al}., \\cite{chand1,chand2} using a partial version of the many-multiplet method at first produced a result consistent with no variation in $\\alpha $, with an unusually small uncertainty, $\\Delta \\alpha /\\alpha =(-0.06\\pm 0.06)\\times 10^{-5}.$ However, the data reduction did not allow $\\Delta \\alpha $ to be a free parameter in the data fitting, and a reanalysis of the same data set by Murphy \\emph{et al}. \\cite{murph} using the full many-multiplet method increases the uncertainties sixfold, and leads to a revised bound of% \\begin{equation*} \\Delta \\alpha /\\alpha =(-0.64\\pm 0.36)\\times 10^{-5}. \\end{equation*} Any present-day variation of $\\alpha $ can also be constrained by direct laboratory comparisons of clocks based on different atomic frequency standards over a period of months or years. Until recently, the most stringent atomic clock constraints on any current temporal variation of $% \\alpha $ were \\begin{equation*} \\dot{\\alpha}/\\alpha =(-3.3\\pm 3.0)\\times 10^{-16}\\,yr^{-1}, \\end{equation*}% which arose by combining measurements of the frequencies of Sr \\cite{blatt}, Hg \\cite{fortier}, Yb \\cite{peiknew}, and H \\cite{fischer} relative to Caesium; Cing\\\"{o}z \\emph{et al}. \\cite{cingoz} have also recently reported a less stringent limit of $\\dot{\\alpha}/\\alpha =-(2.7\\pm 2.6)\\times 10^{-15}\\,yr^{-1}$. If the systematic errors can be fully understood, an ultimate sensitivity of $10^{-18}\\,yr^{-1}$ may be possible with this method \\cite{nguyen}. If a linear variation in $\\alpha $ is assumed then the Murphy \\emph{et~al}. quasar measurements equate to $\\dot{\\alpha}/\\alpha =(6.4\\pm 1.4)\\times 10^{-16}\\,yr^{-1}$ \\cite{webb1, webb2}. If the variation is due to a light scalar field described by a theory like that of Bekenstein and Sandvik, Barrow and Magueijo (BSBM) \\cite{bek, bsbm}, then the rate of change in the constants is exponentially damped during the recent dark-energy-dominated era of accelerated expansion, and one typically predicts a present-day value of \\begin{equation*} \\emph{\\ }\\dot{\\alpha}/\\alpha =1.1\\pm 0.3\\times 10^{-16}\\,yr^{-1}\\emph{\\ } \\end{equation*}% by direct extrapolation from the Murphy \\emph{et al}. data \\cite{webb1, webb2}. This is not ruled out by the atomic clock constraints mentioned above. For comparison, the Oklo natural reactor constraints, which are based on the need for the $\\mathrm{Sm}^{149}+n\\rightarrow \\mathrm{Sm}^{147}+\\gamma $ neutron capture resonance at $97.3~\\mathrm{MeV}$ to have been present $% 1.8-2~\\mathrm{Gyr}$ ago at $z=0.15$, as first pointed out by Shlyakhter \\cite% {sh}, are currently \\cite{fuj} $\\Delta \\alpha /\\alpha =(-0.8\\pm 1.0)\\times 10^{-8}$ or $(8.8\\pm 0.7)\\times 10^{-8}$ (because of the double-valued character of the neutron capture cross-section with reactor temperature) and \\cite{lam} $\\Delta \\alpha /\\alpha >4.5\\times 10^{-8}$ $(6\\sigma ),$ when the non-thermal neutron spectrum is taken into account. However, there remain significant environmental uncertainties regarding the reactor's early history and the relationship between changes in the resonance energy level and those in the values of any underlying constants. For reviews of the wider issue of varying constants in addition to $\\alpha $, see the reviews in refs.~\\cite{revs}, and for some implications of the unification of fundamental forces see refs.~\\cite{unify}.  Recently, Rosenband \\emph{et al}. \\cite{Rosenband} measured the ratio of aluminium and mercury single-ion optical clock frequencies, $f_{\\mathrm{Al+}% }/f_{\\mathrm{Hg+}}$, at intervals over a period of about a year. From these measurements, the linear rate of change in this ratio was found to be $% (-5.3\\pm 7.9)\\times 10^{-17}\\,yr^{-1}$ (but see ref. \\cite{bs} for some refinements). These measurements provides the strongest limit yet on any temporal drift in the value of $\\alpha $: \\begin{equation*} \\ \\dot{\\alpha}/\\alpha =(-1.6\\pm 2.3)\\times 10^{-17}\\,yr^{-1}.\\ \\end{equation*}% This limit is strong enough to exclude theoretical explanations of the change in $\\alpha $ reported by Webb \\emph{et al}. \\cite{webb1, webb2} based on the slow variation of an effectively massless scalar field \\cite{bsbm}, even allowing for the damping by cosmological acceleration, unless there is a significant effect that slows the locally observed effects of changing $% \\alpha $ on cosmological scales (for a detailed analysis of global-local coupling of variations in constants (see Refs.~\\cite{shawb}).  Theories in which $\\alpha $ varies will in general lead to violations of the weak equivalence principle (WEP). This is because the $\\alpha $ variation is carried by a scalar field, $\\varphi ,$ and this couples differently to different nuclei because they contain different numbers of electrically charged particles (protons). The theory discussed here has the interesting consequence of leading to a relative acceleration of order $10^{-13}$ \\cite% {bmswep} if the free coupling parameter is fixed to the value given in Eq.~(% \\ref{om}) using a best fit of the theories cosmological model to the quasar observations of refs.~\\cite{webb1, webb2}. Other predictions of WEP violations have also been made in refs. \\cite{poly, zal, dam}. The observational upper bound on this parameter from direct experiment is just an order of magnitude larger, at $10^{-12},$ and limits from the motion of the Moon are of similar order, \\cite{nord}, but space-based tests planned for the STEP mission are expected to achieve a sensitivity of order $% 10^{-18} $ and will provide a completely independent check on theories of time-varying $e$ and $\\alpha $ \\cite{wep, step}.  In view of this tension between direct local measurements and astronomical measurements of the fine structure 'constant' it is important to explore the widest possible range of self-consistent theoretical models for the time-evolution of $\\alpha $ so as to understand the possible evolutions of $% \\Delta \\alpha /\\alpha $ over the range $0<z<6$ that spans the astronomical, geochemical, and laboratory measurements. In the remainder of this paper we will present cosmological extensions to the simple BSBM scalar field models for varying $\\alpha $ that include a non-zero self-interaction potential, $% V(\\varphi )$ for the scalar field, $\\varphi $, carrying the spacetime evolution of $\\alpha $. We will consider two representative theories, where $% V$ has exponential and power-law variation, respectively, and determine the solutions for the cosmological evolution and the time-variation of $\\alpha $ during the radiation, dust, and dark-energy dominated eras of the universe.  The organization of this paper is as follows: in \\S ~\\ref{sect:varying_alpha} and \\S ~\\ref{sect:cosmo_eqns} we present the theory of varying $\\alpha $ based on the coupling of a scalar field to the electromagnetically charged matter and list the relativistic equations for the investigations of this theory. In \\S ~\\ref{sect:BSBM} and \\S ~\\ref{sect:pot_only} we review, respectively, the cosmological evolutions of the scalar field $\\varphi $ for the model with no scalar field self-interaction potential (just coupling with matter) and for quintessence models with exponential and inverse power-law self-potentials. The \\S ~\\ref{sect:pot_and_coup} is devoted to an investigation of how the scalar field $\\varphi $ evolves if both the matter coupling term and the bare self-potential are non-zero, which is supplemented by the numerical examples shown in \\S ~\\ref{sect:numerics}. Finally, conclusions are drawn in \\S ~\\ref{sect:conclusion}. ",
        "conclusions": "\\label{sect:conclusion}  In this paper we have considered the dynamics of the varying-$\\alpha $ theories which arise when the original BSBM theory is generalised by introducing an exponential or inverse power-law self-potential for the scalar field driving the variation of $\\alpha $. These two representative potentials capture the essential ingredients of general potentials without minima. There are two situations to distinguish and analyse separately: according as to whether or not the scalar-field potential comes to dominate the late-time dynamics of the universe. In combination with the coupling with matter, the additional bare potential forms an effective total potential $V_{eff}$ for the scalar field $\\varphi $ which governs the allowed time variation of $\\alpha $. We have presented the solutions to the scalar-field equation of motion in cases where $V_{eff}$ is dominated solely by the coupling term or the bare potential, respectively, in different cosmic eras. In most cases the bare-potential-dominated solution differs from the coupling-dominated one; the contributions of these two terms to $% V_{eff}$ vary with time, and it is possible for there to be a transition from one solution-type to another. The numerical results show that the transition between solutions types can be very smooth, and for most of the time the solution for the scalar field tracks either the bare-potential-dominated or the coupling-dominated solution. These features ensure that the evolution of $\\varphi $ under $V_{eff}$ has a very simple pattern. The main results are briefly summarized in Table 1, and these results are quite general, not depending on whether the parameters defining the potential ($\\lambda $ and $\\gamma $) are extremely large or not.  The consequences for the time evolution of the fine structure 'constant' of adding potentials $V(\\varphi )$ \\ to the BSBM theory are summarised as follows. In light of the observational constraints on how much variation in $% \\alpha $ there can be over redshifts $z<6$, we find that the restrictions on the interaction potential parameters ($\\lambda $ in the exponential potential and $\\gamma $ in the inverse power-law potential) must be very strong, in order to prevent the bare potentials from becoming unacceptably dominant and driving unacceptable fast time variation of $\\alpha $. For example, in the case of an exponential potential, $V=V_{0}\\exp (-\\lambda \\varphi ),$ we must have $\\lambda \\gtrsim 10^{6}\\sim 10^{7}$ with $|\\zeta |/\\omega \\sim \\mathcal{O}(10^{-6})$, and for the inverse power-law potential, $V=V_{\\ast }\\varphi ^{-\\gamma },$ we need $\\gamma \\gtrsim 10^{6}\\sim 10^{7}$ with $|\\zeta |/\\omega \\sim \\mathcal{O}(10^{-6})$ -- the exact constraint, of course, depends also on the initial conditions). This means that the $\\lambda $ and $\\gamma $ are constrained to be so large that if they appear in quintessence models then the scalar field is practically indistinguishable from a cosmological constant, which has no dynamics at all. The total variation of $\\alpha $ can be enhanced compared to the case of no bare potential (BSBM), and the variation commences much earlier. Finally, because with an exponential or inverse power-law potential the scalar field will not approach a constant even in the dark-energy-dominated era, the fine structure constant $\\alpha $ will continue to increase for all future time, until eventually there will no stable atoms in the Universe at all \\cite{bsm2}."
    },
    "0808/0808.3299_arXiv.txt": {
        "abstract": "We determine the most likely dark-matter fraction in the elliptical galaxy quadruply lensing the quasar \\pg\\ based on analyses of the X-ray fluxes of the individual images in 2000 and 2008.  Between the two epochs, the $A_2$ image of \\pg\\ brightened relative to the other images by a factor of six in X-rays.  We argue that the $A_2$ image had been highly demagnified in 2000 by stellar microlensing in the intervening galaxy and has recently crossed a caustic, thereby creating a new pair of micro-images and brightening in the process. Over the same period, the $A_2$ image has brightened by a factor of only 1.2 in the optical.  The most likely ratio of smooth material (dark matter) to clumpy material (stars) in the lensing galaxy to explain the observations is $\\sim$90\\% of the matter in a smooth dark-matter component and $\\sim$10\\% in stars. ",
        "introduction": "The theory of gravitational lensing is by now quite well understood \\citep[e.g.,\\,the review by][]{1999fsu..conf..360N}.  For the case of a quasar quadruply imaged by an intervening galaxy, a very simple model for the lensing potentials  --- a monopole plus a quadrupole --- usually succeeds in fitting the positions of quasar images at the 1--2\\% level.  However, it has become increasingly clear that these same models do considerably worse at fitting the {\\em relative fluxes} from quasar images \\citep[e.g.,][]{2004ApJ...610...69K,2002ApJ...567L...5M}. Such ``flux ratio anomalies'' are thought to be the product of small scale structure in the gravitational potentials of the lensing galaxies.  There are two leading explanations for this small scale structure.  One intriguing explanation is that we are seeing \\textit{milli}-lensing by dark matter condensations of sub-galactic mass \\citep{1995ApJ...443...18W,1998MNRAS.295..587M,2002ApJ...572...25D, 2001ApJ...563....9M, 2002ApJ...565...17C}, which are predicted in large numbers in $N$-body simulations. However, the much more likely explanation (and exciting for very different reasons) is that the anomalies are largely the result of \\textit{micro}-lensing by stars in the intervening galaxy \\citep{1995ApJ...443...18W, 2002ApJ...580..685S}.  If the flux ratio anomalies are due to {\\em milli}-lensing, i.e., $10^4$\\,--\\,$10^8$~\\Msun\\ dark matter condensations \\citep{1992ApJ...397L...1W}, then the Einstein radii of such masses, projected back to the quasar, are sufficiently large that we would expect the flux ratios to (i) be the same at all wavelengths, and (ii) remain constant with time (except for source variability).  In fact, neither of these expectations based on {\\em milli}-lensing is observed, and the results are overwhelmingly more compatible with stellar {\\em micro}-lensing.  If this is indeed the case, then (i) the stellar Einstein radii projected back to the quasar are more nearly comparable in size with the expected quasar emission regions, thereby allowing for a probe of the inner regions of their accretion disks; (ii) variations in the flux ratio anomalies with time in one source, or from source to source, can provide a direct measure of the dark-to-stellar matter ratios at projected radial distances of $\\sim$$2-6$ kpc in elliptical galaxies; and (iii) the flux ratio anomalies are expected to vary dramatically on time scales as short as a few years.  We note in passing that item (i) above works because microlensing, in effect, is the most powerful zoom lens in astronomy, probing angular sizes down to $\\sim$$10^{-6}$~arcsec, which is a factor of $\\sim$100 better than even VLBI.  Recently we have systematically analyzed ten quadruply-imaged quasars using \\cxo\\ archival data and \\hubble\\ visible images \\citep{2007ApJ...661...19P}. We find that the flux ratio anomalies in the X-ray images of quads are systematically larger than for the same quads imaged in the visible, by a factor of $\\sim$2.  As expected from the models \\citep[see][]{2002ApJ...580..685S, 2007ApJ...661...19P}, it is the highly magnified saddle-point image among the four images that is most susceptible to stellar microlensing.  \\citet{2007ApJ...661...19P} concluded that the extent of the quasar accretion disks in the optical (i.e., $r_{\\rm opt}$) must be comparable with (i.e., $\\gtrsim 1/3$) the Einstein radii, $r_{\\rm ein}$, of the stellar microlenses in order to reduce the flux ratio anomalies by the factor of $\\sim$2 observed.  This conflicts with values of $r_{\\rm opt}$ calculated from simple accretion-disk models, which are expected to be considerably smaller than the Einstein radii, with typical ratios of $r_{\\rm opt}/r_{\\rm ein}$ in the range of only $0.01 - 0.3$, with a median value of $0.04$ \\citep{2007ApJ...661...19P}.  Thus, the observationally inferred optical emission regions in quasars, based on microlensing, are much larger than anticipated.  This is an intriguing mystery to be pursued.  \\begin{figure}[t] \\centering \\includegraphics[width=0.235\\textwidth,height=0.235\\textwidth]{f1a.pdf}\\hglue0.015\\textwidth \\includegraphics[width=0.235\\textwidth,height=0.235\\textwidth]{f1b.pdf}\\vglue0.015\\textwidth \\includegraphics[width=0.235\\textwidth,height=0.235\\textwidth]{f1c.pdf}\\hglue0.015\\textwidth \\includegraphics[width=0.235\\textwidth,height=0.235\\textwidth]{f1d.pdf} \\caption{Images of \\pg.  Clockwise from upper left: Maximum-likelihood reconstructed \\chandra\\ image from 2000, smoothed by a Gaussian; maximum-likelihood reconstructed \\chandra\\ image from 2008, smoothed by a Gaussian; expected image predicted from singular isothermal sphere plus shear model of lensing galaxy; Magellan $i'$-band image. Each image is $4''\\times4''$.  The apparent lack of any X-ray emission from $A_2$ in 2000 is an artifact of the maximum-likelihood reconstruction.} \\label{fig:images} \\end{figure}  Of equal astrophysical significance, the same observations of the amplitudes and frequency of occurrence of X-ray flux ratio anomalies can also be used to infer the fraction of dark matter at distances from the center of the lensing elliptical galaxies corresponding to the impact parameter of the images \\citep[typically $\\sim$2--6 kpc;][]{2004IAUS..220..103S}. In this paper we pursue this latter line of investigation for the quad lens \\pg.  In particular, we describe a new \\chandra\\ observation of \\pg\\ in January 2008 which indicates that the $A_2$ image has dramatically brightened in X-rays compared to its state in 2000 (see Fig.\\,\\ref{fig:images}).  In \\S2.1 we review prior optical and X-ray observations of \\pg, while in \\S2.2 we present the new \\chandra\\ observations and describe the analysis by which we determined the flux ratios.  In \\S3 we describe how inferences about the dark-to-stellar matter ratio can be made from observations of microlensing.  Finally, in \\S4 we summarize our results. ",
        "conclusions": "We have observed a dramatic change in the X-ray flux of the $A_2$ image of \\pg\\ in \\chandra\\ observations that were separated by eight years.  The short timescale for the flux change clearly indicates the presence of stellar microlensing rather than milli-lensing due to dark matter substructure in the elliptical lensing galaxy.  The observations of the individual fluxes point toward a substantial dark matter fraction of $\\sim 80-95$\\%.  \\begin{figure}[t] \\centering \\includegraphics[width=0.49\\textwidth]{f9.pdf} \\caption{Probability density ({\\it top}) and cumulative probability ({\\it bottom}) of $A_2$ fractional flux in 2008 given the value in 2000.  The dashed lines mark the 1-$\\sigma$ confidence intervals of the measured value in 2008.} \\label{fig:bayes2} \\end{figure}   One particularly interesting aspect of the observed change in flux ratio between the $A_2$ and $A_1$ images is the very rapid timescale on which it occurred ($\\sim$8 years).  Figure \\ref{fig:Wambs} indicates how far bulk translation of the lensing galaxy is likely to move the quasar image with respect to the caustic pattern during an eight-year interval.  For typical expected transverse speeds of $\\sim$300 km s$^{-1}$ it seems unlikely that the image will start on a point of very low magnification and end up with near nominal magnification in just over eight years, as indicated in Fig.~\\ref{fig:bayes2}.  With an unexpectedly high speed of $\\sim$1000 km s$^{-1}$ the likelihood for the observed change in magnification of the $A_2$ image becomes considerably greater.  In this regard, we note that the recent report of an even larger change in a flux ratio in the quad lens RX\\, J1131$-$1231 over an interval of $\\sim$2 years \\citep{Chartas} indicates that the variation that we observed in \\pg\\ is perhaps not a rare occurrence.  The difference between the effect of bulk and random motion of the microlensing stars is probably not sufficient to account for these rapid variations, leaving us with an unresolved puzzle."
    },
    "0808/0808.2640_arXiv.txt": {
        "abstract": "Continuing our study of galaxy populations in the Antlia cluster, we present a photometric analysis of four galaxies classified as compact elliptical (cE) galaxies in the Ferguson \\& Sandage (1990, hereafter FS90) Antlia Group catalogue. Only 6 members of this rare type of galaxies are known until now. Using data in various photometric systems (Washington $C$, Kron-Cousins $R$, Bessel $V$ and $I$, HST $F814W$ and $F435W$), we measured brightness and colour profiles, as well as structural parameters. By comparing them with those of other galaxies in the Antlia cluster, as well as with confirmed cE galaxies from the literature, we  found that two of the FS90 cE candidates, albeit being spectroscopically confirmed Antlia members, are not cE galaxies. However, one of these objects presents strong ellipticity and position angle variations that resemble those already reported for M32, leading us to speculate about this kind of objects being progenitors of cE galaxies. The other two FS90 cE candidates, for which radial velocities are not available, match some features typical of cE galaxies like being close in projection to a larger galaxy, displaying flat colour profiles, and having a high degree of compactness. Only one of the remaining cE candidates shows a high central surface brightness, two components in its brightness profile, distinct changes in ellipticity and position angle where the outer component begins to dominate, and seems to follow the same trend as other confirmed cE galaxies in a luminosity versus mean effective surface brightness diagram. Moreover, it shows a distorted inner structure with similar characteristics to those found by simulations of interacting galaxies, and an extremely faint structure that seems to link this object with one of the Antlia dominant galaxies, has been detected in MOSAIC-CTIO, FORS1-VLT, and ACS-HST images. The cE nature of this galaxy as well as the possible interaction with its bright companion, still have to be confirmed through spectroscopy. ",
        "introduction": "\\label{intro}  Early-type dwarf galaxies are the most frequent galaxy type in nearby groups and clusters of galaxies \\citep[e.g.][]{FS90}. Although the information related with their evolution in such environments is codified in their spatial distribution, chemical composition and kinematics, a single scenario including all their properties in a consistent frame is still missing. The main difficulty resides in the fact that they do not constitute a homogeneous class of objects. Among them, there are nucleated and non-nucleated galaxies, as well as systems that were found to harbour discs and spiral structure \\citep*{J00,CB05,L06b}. Moreover, central star formation has also been reported in some objects \\citep{L06a}.  The so called compact elliptical \\citep[cE, e.g.][]{N87}, compact dwarf elliptical \\citep[cdE,][]{D01} or M32-like \\citep[e.g.][]{ZB98} galaxies are members of the low-luminosity galaxy family, but instead of having low surface brightness like the most common early-type dwarfs, they have notorious high surface brightness \\citep[e.g.][]{N87}. They constitute a very rare group, as there are many galaxies classified as compact ellipticals  \\citep*[see, for instance,][]{B85}, but only five objects have been confirmed as such, in addition to the prototype M32.  The known examples are all companions of larger galaxies. They are M32 itself, a satellite of the M31 (Andromeda) spiral galaxy, NGC\\,4486B close to M87 in the Virgo cluster \\citep{SB84,D91}, NGC\\,5846A close to the giant elliptical NGC\\,5846 (\\citealp{D91}; \\citealp*{Ma05}), A496cE close to the central cluster dominant (cD) galaxy of the cluster Abell 496 \\citep{Ch07}, and two objects in the Abell 1689 cluster \\citep[CG$_{A1689,1}$ and CG$_{A1689,2}$,][]{M05}. It should be noted that object CG$_{A1689,1}$ from \\citep{M05}, has a quite deviating radial velocity with respect to its closest projected giant elliptical.  In particular, M32 presents the following characteristics: (a) it is a satellite of a larger galaxy; (b) the brightness profile cannot be accurately fitted with a single S\\'ersic law \\citep*{G02,choi02}; (c) high surface brightness in comparison with ellipticals of the same luminosity \\citep{N87} and correspondingly a small size \\citep{G02,choi02}; (d) a radial change in both age and metallicity, leading however to a quite flat colour profile, its stellar population being younger and more metal-rich at the centre \\citep{R05}. The other confirmed cE galaxies show similar properties with respect to surface brightness, compactness and projected location close to brighter galaxies. However, there might be some differences in their colour gradients and brightness profiles (see, for example, \\citealp{L96} and \\citealp{F06}, regarding NGC\\,4486B) as well as in metal content and age \\citep[e.g.][]{SB06,Ch07}.  Some effort has been made to find more examples of these very rare objects in nearby clusters or groups, like Fornax \\citep{D01} and Leo \\citep{ZB98}. However, all the examined candidates have been rejected as cE galaxies. As the questions related to their origin and their role in the framework of galaxy evolution remain still open, it would be of great interest to find more objects and to study them in relation to their environment.  The Antlia cluster of galaxies is the third nearest well populated galaxy cluster after Virgo and Fornax. It exhibits a complex structure consisting of several subgroups, the most conspicuous ones being dominated by the giant elliptical galaxies NGC\\,3258 and NGC\\,3268. X-ray observations showed extended emissions around both subgroups \\citep*{ped97,nak00}. These emissions are concentrated towards the dominant galaxies, but extensions elongated in the direction to the other subgroup are also present, suggesting an ongoing merger. \\citet*{dir03} and \\citet*{B08} have shown that the globular cluster systems around NGC\\,3258 and NGC\\,3268 are elongated in the same direction as a connecting line between the two galaxies, resembling the X-ray results.  The photographic survey of \\cite{FS90} was the first and last major effort devoted to study the galaxy population of the Antlia cluster. They identified, by visual inspection, 375 galaxies that are listed in their Antlia Group Catalogue (hereafter labelled by the acronym FS90 plus the catalogue number). It gives, among other data, a membership status (1. definite member, 2. likely member, 3. possible member) and a morphological type for each galaxy. FS90 classified 10 objects as E(M32?) or S0(M32?), and 1 object as d:E(M32?),N.  Given the lack of an extensive analysis of Antlia's galaxy population, we initiated the Antlia Cluster Project with the aim at performing the first CCD-photometric and spectroscopic study of this cluster. In the first paper of this long-term project \\citep[][hereafter Paper I]{SC08}, we presented photometric properties of the early-type galaxy population. Among our results it was found that early-type members define a tight colour-magnitude relation, spanning 9 mag from giant ellipticals to dwarf galaxies, without a perceptible change of slope. This slope is similar to those found in other clusters like Virgo, Fornax, Perseus and Coma, which are dynamically different to Antlia, and it is also consistent with that displayed by the so called `blue tilt' of metal-poor globular clusters in NGC\\,4486.  Among the Paper I sample, there were four galaxies classified as cE candidates by \\citet{FS90}. Two of these objects are spectroscopically confirmed Antlia members that seem to be normal low-luminosity early-type galaxies. The other two are each one close in projection to either one of the dominant galaxies, and are separated from the locus of the early-type members in the luminosity-mean effective surface brightness diagram. They display some characteristics of cE galaxies, but radial velocities are not available.  In this paper, the second of the Antlia Cluster Project, we present a photometric analysis of four FS90 cE candidates (namely FS90\\,110, FS90\\,165, FS90\\,192, and FS90\\,208) located in the central region of the Antlia cluster. We aim at obtaining photometric evidence in favour of, or against to, these objects being genuine cE galaxies. Given the small number of cE galaxies known to date, any additional members of this class could give important clues on their evolutionary path. The central question is probably whether they have been dynamically transformed by the interaction with a massive galaxy.  Throughout this paper we will adopt $(m-M)=32.73$ as the Antlia distance modulus \\citep{dir03}, which corresponds to an Antlia distance of 35.2 Mpc. At this distance 1 arcsec subtends 170 pc. The paper is organized as follows. In Section\\,\\ref{data} we give information about our photometric data. In Section\\,\\ref{analysis} and \\ref{HST} we present the analysis of the data, and in Section\\,\\ref{conclusions}, a discussion and our conclusions. ",
        "conclusions": "\\label{conclusions}  \\subsection{FS90\\,165}  From the analysis performed in the previous section, we conclude that FS90\\,165, an spectroscopically confirmed Antlia member, cannot be considered as a cE galaxy. It follows fundamental relations of low surface brightness early-type Antlia galaxies (see also Paper I), is larger than any confirmed cE object, and presents no compact morphology. It is also located far away in projection from any bright galaxy, and from the central cluster region, a characteristic shared by all confirmed cE galaxies. Its $T_1$ brightness profile is well fitted by a single S\\'ersic law along its whole radius range, i.e. it does not show two components. However, the S0 morphological classification assigned to this galaxy by FS90 is in agreement with the B4 coefficient values that point to disky isophotes, and with the red inner disc found in its $(C-T_1)$ colour map and confirmed through unsharp masking. There are no images from VLT or HST for this object.  \\subsection{FS90\\,208}  Certainly, FS90\\,208 is not a cE galaxy either. It is a confirmed Antlia member that follows the same fundamental relations as FS90\\,165 and the rest of low surface brightness early-type galaxies, and it is not compact. Besides, it is not close in projection to any bright galaxy and is far from the centre of the cluster. However, it displays a two component brightness profile and presents important variations in ellipticity and position angle that resemble those shown by M32 \\citep{choi02}, albeit stronger. Curiously, these variations arise at the same galactocentric radii than those displayed by FS90\\,110, and behave in a similar way. FS90 have classified this object as an S0 galaxy, which is in agreement with the two components present in its brightness profile, as well as with the B4 coefficient which indicates that the isophotes are nearly disky. Moreover, a process of unsharp masking reveals what seems to be an inner bar structure, slightly detected with the subtraction of a fixed ELLIPSE model of the galaxy from the MOSAIC image. However, a triaxial object in which the axial ratios vary with radius, could also display isophote twisting \\citep{BM98}.  \\subsection{FS90\\,192}  There is no radial velocity available for FS90\\,192. It displays no perceptible colour gradients and, if it was an Antlia member, it would be as compact as other cE galaxies. Furthermore, its projected distance to NGC\\,3268 would be smaller than that between NGC\\,4486B and NGC\\,4486. Although it does not display two components in its brightness profile, it is worth noting that NGC\\,4486B in the Virgo cluster presents a single component profile as well and a similar n (i.e. 1/N) value.  Despite all the interesting features mentioned above, we cannot confirm nor rule out that this object is in the background. It does not follow the early-type Antlia members fundamental relations, as it appears shifted towards higher surface brightnesses, fainter magnitudes, or redder colours. M32 would hold a comparable position on the colour-magnitude diagram if it were placed at the Antlia distance. However, FS90\\,192 does not follow the relation between mean effective surface brightness and luminosity defined by the confirmed cEs. Under the hypothesis that it is a background object, FS90\\,192 would become an isolated elliptical galaxy of moderate luminosity (highly unlikely), as we have not found any concentration of galaxies in its neighbourhood.  \\subsection{FS90\\,110}   There is no radial velocity available for FS90\\,110 and, as stated in Sect.\\ref{K-corr}, it could be a background galaxy placed at $\\sim$120 Mpc. However, it seems to be the most firm candidate for a cE galaxy.  If it was an Antlia member, its projected separation from NGC\\,3258 would be of the same order as those of M32 from M31, and NGC\\,5486A from NGC\\,5486. Moreover, it would be much closer to the dominant galaxy than NGC\\,4486B to NGC\\,4486, and its effective radius would be smaller than that of NCG\\,5846A.  It presents a high central surface brightness in the brightness profile obtained from the ACS image, and its $(C-T_1)$ and $(V-I)$ colour profiles are quite flat, in agreement with what is found for M32. Its brightness profile shows two components in the equivalent radius range 2 -- 11.8 arcsec, that seems to be described by N $>$ 1 (i.e. n $<$ 1) S\\'ersic indices. It is tempting to suggest that such N index for the outer component might be linked to a truncation of the profile \\citep*{E08}, due to a possible interaction with NGC\\,3258. However, such hypothesis should be studied in the light of the confirmed membership status of FS90\\,110. Furthermore, it would be interesting to test if the outer component of confirmed cE galaxies displaying two components in their brightness profiles, could be fitted by a S\\'ersic law with N $>$ 1 (n $<$ 1).  The variation of the ellipticity and position angle vs. radius displayed by FS90\\,110, resemble those found in M32. These changes occur at a galactocentric radius at which the outer component seems to begin to dominate, in agreement with what \\citet{G02} found for M32.  FS90\\,110 is the only FS90 cE candidate that follows the luminosity vs. mean effective surface brightness relation defined by bright ellipticals, which corresponds to the \\citet{K77} scaling relation, towards fainter magnitudes, smaller radius, or higher mean effective surface brightness. It also seems to follow the Antlia members' colour-magnitude relation, though it is located at the red border.  The elongation and twisting of FS90\\,110's outermost isophotes in the direction of NGC\\,3258 are noticeable. That would be consistent with the extremely low surface brightness structure detected in the MOSAIC images, and confirmed with the FORS1 and ACS frames. Such kind of `bridge' that seems to link FS90\\,110 with its bright partner, would be in agreement with the fact that most confirmed cE galaxies are located in the vicinity of brighter companions. In that sense, it is also remarkable the warped inner structure displayed in the ACS colour map, and confirmed through unsharp masks of FORS1. Quite similar stellar structures are found in numerical simulations \\citep{B01,M06} as a consequence of galaxy interactions.  \\subsection{Final remarks}  A galaxy that is interacting with a more massive partner will feel tidal forces most strongly  in its outskirts, while its central region will be less affected \\citep[see, for example,][for a model of the interaction between M32 and M31]{B01}. As a consequence, it would not be surprising to observe asymmetric and `egg-shaped' outer isophotes, or/and detect low surface brightness stellar `tails', arising from the disruption of the outer regions of the satellite \\citep{M06}. In particular, these faint structures could be observed either like one or two symmetric `tails', depending on projection effects \\citep[see fig.\\,4 in][]{M06}. The distorted outermost isophotes of FS90\\,110, and the extremely faint structure detected in the HST, VLT and CTIO images fit quite well in such a scenario.  The similarities in the ellipticity and position angle variations against radius displayed by FS90\\,110 and FS90\\,208, which arise at a similar galactocentric radius of $\\sim 5.5$ arcsec (i.e. in their inner regions), are also remarkable. These kind of variations have already been detected in M32, although not so strongly. All these pieces of evidence lead us to speculate about the possibility that, an object similar to FS90\\,208 might be the progenitor of a cE galaxy. Thus, it is tempting to look for possible links between these two kind of objects.  We may think that a system with similar characteristics as FS90\\,208 may lose its outermost regions due to the interaction with a bright companion, then becoming compact. As a consequence, it could also experience a redistribution of its stellar content, that might produce a warped inner structure, a higher central surface brightness, and an attenuation of its ellipticity and position angle variations with radius. Moreover, the bulge-to-disc ratio in such object could increase after losing its outer parts due to the interaction. With regard to this point, it should be noted that the inner component of FS90\\,110 seems to be 3 times brighter than the outer one, while that of FS90\\,208 is just 1.6 times brighter in relation with the outer component.  As an aditional point, we recall that the Antlia cluster seems to be particularly rich in S0 galaxies and FS90\\,208 seems to be one of them. If the evolutionary path of cE galaxies include this kind of objects, S0 rich clusters would arise as favorable environments for the formation of compact galaxies.  Dynamical simulations and the spectroscopically confirmed membership status of FS90\\,110 and FS90\\,192, will help to test if any of the above statements is indeed plausible."
    },
    "0808/0808.0190_arXiv.txt": {
        "abstract": "We show that simple scalar field models can give rise to curvature singularities in the effective Friedmann dynamics of Loop Quantum Cosmology (LQC). We find singular solutions for spatially flat Friedmann-Robertson-Walker cosmologies with a canonical scalar field and a negative exponential potential, or with a phantom scalar field and a positive potential. While LQC avoids big bang or big rip type singularities, we find sudden singularities where the Hubble rate is bounded, but the Ricci curvature scalar diverges. We conclude that the effective equations of LQC are not in themselves sufficient to avoid the occurrence of curvature singularities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.1196_arXiv.txt": {
        "abstract": "Elucidating the processes that governed the assembly and evolution of galaxies over cosmic time is one of the main objectives of all of the proposed Extremely Large Telescopes (ELT). To make a leap forward in our understanding of these processes, an ELT will want to take advantage of Multi-Objects Adaptive Optics (MOAO) systems, which can substantially improve the natural seeing over a wide field of view. We have developed an end-to-end simulation to specify the science requirements of a MOAO-fed integral field spectrograph on either an 8m or 42m telescope. Our simulations re-scales observations of local galaxies or results from numerical simulations of disk or interacting galaxies. The code is flexible in that it allows us to explore a wide range of instrumental parameters such as encircled energy (EE), pixel size, spectral resolution, etc. For the current analysis, we limit ourselves to a local disk galaxy which exhibits simple rotation and a simulation of a merger. While the number of simulations is limited, we have attempted to generalize our results by introducing the simple concepts of ``PSF contrast'' which is the amount of light polluting adjacent spectra which we find drives the smallest EE at a given spatial scale. The choice of the spatial sampling is driven by the ``scale-coupling''. By scale-coupling we mean the relationship between the IFU pixel scale and the size of the features that need to be recovered by 3D spectroscopy in order to understand the nature of the galaxy and its substructure. Because the dynamical nature of galaxies are mostly reflected in their large-scale motions, a relatively coarse spatial resolution is enough to distinguish between a rotating disk and a major merger. Although we used a limited number of morpho-kinematic cases, our simulations suggest that, on a 42m telescope, the choice of an IFU pixel scale of 50-75 mas seems to be sufficient. Such a coarse sampling has the benefit of lowering the exposure time to reach a specific signal-to-noise as well as relaxing the performance of the MOAO system. On the other hand, recovering the full 2D-kinematics of z$\\sim$4 galaxies requires high signal-to-noise and at least an EE of 34\\% in 150 mas (2 pixels of 75 mas). Finally, we carried out a similar study for a hypothetical galaxy/merger at z=1.6 with a MOAO-fed spectrograph for an 8m, and find that at least an EE of 30\\% at 0.25 arcsec spatial sampling is required to understand the nature of disks and mergers. ",
        "introduction": "Developing a coherent model for the mass assembly of galaxies over cosmic time is a complex and difficult task. It is difficult because developing such a coherent model involves highly non-linear physics (e.g., the cooling and collapse of baryons, or feedback from stars and super-massive blackholes which regulate both the growth of the stellar mass and black holes themselves) and all over a very wide range of physical scales (from large scale structure to star clusters to black holes). Our knowledge of galaxy formation  {\\it in situ} relies mostly on observations of integrated quantities,which is insufficient to understand the details of the process of formation, e.g., the interplay between the baryonic and dark matter, or the complex physics of baryons such as merging, star formation, feedback, various instabilities and heating/cooling mechanisms, etc. At the heart of our developing understanding of galaxy evolution is the relative importance of secular evolution through (quasi-)adiabatic accretion of mass, instabilities, and resonances (e.g., \\citealt{Semelin05}), versus more violent evolution through merging (e.g., \\citealt{Hammer05}), as a function of look-back time.  Because of the complexity of the processes involved, understanding the mass assembly of galaxies over cosmic time requires constraining a wide range of phenomenology for large samples of objects. In this respect integral field spectroscopy is particularly well suited in that it allows us to map the spatially resolved physical properties of galaxies (see, e.g., \\citealt{Puech06b}). Unfortunately, splitting the light into many channels in a integral field spectrograph requires long integration times, even on large telescopes, to detect low surface brightness emission. Thus, conducting large statistical studies necessary for understanding the processes driving galaxy evolution will demand a high multiplex capability to efficiently investigate sufficient numbers of objects.  Using current facilities on 8m class telescopes, it is already possible to map the physical properties of distant galaxies. For example, FLAMES/GIRAFFE on the VLT offers the unique ability to observe 15 distant galaxies simultaneously. The first (complete) 3D sample of galaxies at moderate redshifts has revealed a large fraction ($\\sim$40\\%) of all z$\\sim$0.6 galaxies with $M_B <$-19.5 have perturbed or complex kinematics (\\citealt{Flores06, Yang07}). This population of dynamically disturbed galaxies is most probably a result of a high prevalence of mergers and merger remnants (see also \\citealt{Puech06a, Puech08}). At higher redshift, integral field spectroscopy of several objects has been obtained using single object integral field spectrographs available on 8-10 meter class telescopes (e.g., \\citealt{Forster06}). However, many of these data sets were obtained with limited spatial resolution sampling galaxies on scales of a few kpc, and the true dynamical nature of many sources sampled at these spatial and spectral resolutions remains uncertain (e.g., \\citealt{Law07}).  Improving and extending this approach to the general understanding of the nature of faint galaxies and to galaxies in the early Universe will require a new generation of multi-object integral field spectrographs working in the near-IR, with increased sensitivity and better spatial resolution on even larger telescopes. Because of the complex interplay between spatial and spectral features and our limited knowledge of high redshift galaxies, it is helpful, perhaps necessary, to rely on numerical simulations for constraining the design parameters of this new kind of instrument. For this purpose, we have developed software that simulate end-to-end the emission line characteristics of local galaxies and numerical simulations of galaxies to show how they would appear in the distant Universe. By changing various parameters like the resolution, pixel scale, and point spread function it is possible to constrain the instrumental characteristics and performance against a set of galaxy characteristics (e.g., velocity field). This paper is the first in a series to investigate the scientific design requirements of possible future instruments, telescope aperture and design, and site characteristics appropriate for constrain how galaxies grew and evolved with cosmic time. In this paper in particular, we aim to present the details of our assumptions that go into the software, as well as a first attempt to constrain some of the high level specification for a hypothetical but realistic multi-object integral field spectrograph assisted by MOAO system. Because 2D kinematics (i.e., velocity field and velocity dispersion maps) are one of the most obvious and easiest to currently simulate, they provide the most straight forward way of constraining the high level science design requirements for this type of instrument. We emphasize that the goal of this paper is not to study the detailed scientific capability of such instruments. Instead, the present paper aims at examining a few scientifically motivated cases to derive relevant specifications for such types of instruments with a capable MOAO system. These specifications will then be adopted as a baseline for studying their scientific capabilities subsequent papers. This paper is organized as follows: in Section 2, we detail the goals of this study and our methodology; in Sect. 3 we present the new simulation pipeline and in Sect. 4 how MOAO PSFs are simulated; In Sect. 5 we describe the kinematic measurements; In Sect. 6 and 7, we present the simulations and their results, which are discussed in Sect. 8. In Sect. 9, we draw the conclusions of this study. ",
        "conclusions": "We have developed software capable of end-to-end simulations of integral field spectrograph fed by an MOAO system. We have used this software to investigate a limited number of simulations in H-band, in order to give insights into how such MOAO-fed systems should perform on the VLT and the E-ELT. By requiring that the instrument is able to distinguish between a rotating disk and a major merger, we were able to constrain the required Ensquared Energy to be able to recover their large-scale motions. Separating the progenitors of the major merger requires only modest EE (15-26\\% at z=4 in a sample of 100-150mas; 30\\% at z=1.6 in 0.25 arcsec), but long exposure times ($\\sim$24hr). Higher EEs are generally needed to recover the full 2D kinematics in our simulations, typically 35\\%. Provided that the total SNR is large enough over the galaxy (typically $\\sim$ 5), it is possible to use more sophisticated methods, such as the diagnostic diagram introduced here, to distinguish between both types of objects, even at relatively low integration time for such distant objects (8 hrs). To generalize these results beyond what we described here, we have related these specifications to the concept of ``PSF contrast'', which is the amount of polluting light from adjacent spectra and which drives the minimal EE at a given spatial scale. The choice of the latter is driven mostly by the ``scale-coupling'' which is the relative size of the IFU pixel and the size of the features that the observe wishes to recovered in the data cube. Distinguishing between a major merger and a disk, for example, can be largely done by investigating the large scale motions within a system, and a relatively coarse spatial resolution appears then to be sufficient. Given this situation, a pixel scale of 50-75 mas seems to be a relatively good choice, since it provides sufficient SNR in less time than a finer sampling and also has the advantage of relaxing the MOAO system requirements. More stringent requirements can be set by attempting to resolve and investigate structure in high redshift galaxies and that will be simulated in subsequent papers."
    },
    "0808/0808.0755.txt": {
        "abstract": "We present a spectroscopic analysis of five stellar streams (`A', `B', `Cr', `Cp' and `D') as well as  the extended star cluster, \\ec4, which lies within \\streamc, all discovered in the halo of M31 from our CFHT/MegaCam survey. These spectroscopic results were initially serendipitous, making use of our existing observations from the DEep Imaging Multi-Object Spectrograph mounted on the Keck~II telescope, and thereby emphasizing the ubiquity of tidal streams that account for $\\sim$70\\% of the M31 halo stars in the targeted fields. Subsequent spectroscopy was then procured in \\streamc\\ and \\streamd\\ to trace the velocity gradient along the streams. Nine metal-rich ([Fe/H]$\\sim$-0.7) stars at $v_{\\rm hel}=-349.5$~km/s, $\\sigma_{v, corr}\\sim5.1\\pm2.5$\\,km/s are proposed as a serendipitous detection of \\streamcr, with followup kinematic identification at a further point along the stream. Six metal-poor ([Fe/H]$\\sim$-1.3) stars confined to a narrow, 15~km/s velocity bin centered at $v_{\\rm hel}=-285.6$~km/s,   $\\sigma_{v,corr}=4.3^{+1.7}_{-1.4}$~km/s represent a kinematic detection of \\streamcp, again with followup kinematic identification further along the stream. For the cluster \\ec4, candidate member stars %consistent with the HST-constrained red giant branch with average [Fe/H]$\\sim$-1.4 ([Fe/H]$_{spec}$=-1.6), are found at $v_{\\rm hel}=-285$~km/s  suggesting it could be related to \\streamcp. No similarly obvious cold kinematic candidate is found for \\streamd, although candidates are proposed in both of two spectroscopic pointings along the stream (both at $\\sim-400$km/s). Spectroscopy near the edge of \\streamb\\ suggests a likely kinematic detection at $v_{\\rm hel}\\sim-330$~km/s, $\\sigma_{v, corr}\\sim6.9$\\,km/s, while a candidate kinematic detection of \\streama\\ is found (plausibly associated to M33 rather than M31) with $v_{\\rm hel}\\sim-170$~km/s, $\\sigma_{v, corr}=12.5$\\,km/s. The low dispersion of the streams in kinematics, physical thickness, and metallicity makes it hard to reconcile with a scenario whereby these stream structures as an ensemble are related to the giant southern stream. We conclude that the M31 stellar halo is largely made up of multiple kinematically cold streams. ",
        "introduction": "\\footnotetext[1]{The data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation.}  %-giant streams and dwarfs intro  %{\\bf PERHAPS still to add: mass measurements of the streams; chemical analysis of the stream stars on average relative to the background halo population in these fields - no differences photometrically? what about spectral features?}  Stellar streams represent the visible debris of small galaxies being cannibalized by large galaxies, memorials to the merging process by which the halos of galaxies are built up. The best known examples of streams in the Milky Way (MW) have recently been mapped far more extensively by the SDSS-DR5 by Belokurov et al.\\ (2006, 2007): the tidally stripped stars and globular clusters associated with the Sagittarius dwarf spheroidal and the Low Latitude stream, along with a newly discovered ``Orphan Stream'' so named for its lack of obvious progenitor. In M31, thanks to our ability to efficiently map vast regions of the halo, the number of discovered giant streams already outnumbers that of the MW (Ibata et al.\\ 2007). However, of the eight stellar streams that have been identified in the halo of M31, only one has plausibly been identified to a dwarf satellite: the loop connecting to NGC~205 (McConnachie et al.\\ 2005). This suggests that the other streams could represent an additional seven `uncataloged' satellites, although some of the streams might be produced by a common progenitor, as suggested by models of Fardal et al.\\ (2007, 2008) for the M31 Giant Southern Stream, or in a similar fashion to the Sgr dwarf and its numerous wraps around the Milky Way. It is important to characterize their orbits, metallicities and masses, to understand what their progenitors must have been.  % These streams are still recognizable as ...  The existence of stellar streams tells us that a progenitor galaxy has undergone significant mass loss. This is due to a combination of its orbit and its phase of evolution -- the amount of dark matter mass the galaxy has lost (Pe\\~narrubia et al.\\ 2008a,b). % %Dwarfs which are on more circular orbits are in principle less likely to %be in a phase of catastrophic mass loss (although interactions with the disk  can lead to their disruption -- Pe\\~narrubia et al.\\ 2006), but Satellites on circular orbits are harder to disrupt, but if they are massive enough (i.e.\\ of the order of the LMC), dynamical friction will bring them close to the host galaxy centre, where the interactions with the disc will lead to their tidal disruption -- e.g., Pe\\~narrubia et al.\\ 2007). In addition, to form a stream one has to remove most of the dark matter halo ($\\sim$90--99\\%). The most important parameter that controls the mass loss rate of a dSph is the pericentre distance (the orbital eccentricity is of second order). On the other hand dwarfs with highly elliptical orbits spend a lot of time near apocentre where they are unlikely to be disrupted by the host. It is therefore not immediately obvious that streams represent preferred types of orbits on average. However, a number of theoretical studies have shown that significant information about the orbital properties of the progenitor galaxy can be derived from the streams (e.g, Ghigna et al.\\ 1998; Helmi et al.\\ 1999a).  Streams can be much more informative to study than dwarfs because their orbits can be directly traced and constrained. Fellhauer et al.\\ (2006) were able to accurately constrain the shape of the Galactic potential through the bifurcation of Sagittarius streams in Belokurov et al.\\ (2006). Understanding the range of orbits of satellites to large galaxies will help us to understand how the halos of these galaxies formed. This is especially interesting in light of M31's huge stellar halo reflecting the dark matter dominated halo out to $\\simgt$150~kpc  (e.g., Irwin et al.\\ 2005, Gilbert et al.\\ 2006,  Ibata et al.\\ 2007). However streams can also be much harder to analyse observationally: the distances are problematic, there's a much lower spatial density and they have a larger extent so that observational sampling is not trivial. There is also the difficulty to infer the membership of different stream pieces, especially if we expect different chemical signatures due to metallicity gradients in the progenitor system (e.g., Ibata et al.\\ 2007).   %M31 has 16 dwarf galaxy satellites, along with M33 which is likely %orbiting M31 as well. %Importantly, the faint dSph satellites of the MW and M31 begin %to depart significantly from the established correlations between %mass and metallicity, and M/L versus light.  Streamy/blobby structures %(other than the giant stream) are individually interesting and constraining for the halo formation. They can represent the only traceable product of long disintegrated progenitors, yet retain a coherent body for statistical analysis. Streams can provide important clues on the structure of the progenitors (e.g.\\ metallicity gradients, mass to light -- M/L) as well as on the shape of the host dark matter halo (e.g.\\ prolate versus oblate) (Martinez-Delgado et al.\\ 2008). Future study of these structures will be able to put them in a much better \"near-field cosmology\" context, eventually understanding their ages and chemical histories. However it is important to uncover and study them now, even in limited capacities necessitated by the small numbers of spectroscopically identifiable stars and HR-diagram depths, so we can build our models on the most complete context.  We have initiated a spectroscopic survey of the new streams found in M31's halo using the DEep Imaging Multi-Object Spectrograph on Keck~II to derive radial velocities and metallicities of red giant branch (RGB) stars. In this contribution, we discuss spectroscopic pointings in each of streams `C' and  `D' which we obtained by serendipity, since the spectroscopy was taken prior to knowledge of the photometrically discovered streams, as well as followup spectroscopic pointings in both of these streams. We also analyze spectroscopic data from Koch et al.\\ (2008) lying within the Ibata et al.\\ (2007) streams `A' and `B'.  %{\\bf %shows how streams really must be everywhere? quantify how lucky we were?}  %summarize the characteristics of the stream in the %intro (from Ibata et al. 07) %-extent, metallicities etc. ",
        "conclusions": "In conclusion, we have conducted a Keck/DEIMOS spectroscopic survey of five stellar streams, recently uncovered through deep imaging observations of the halo.  $\\bullet$ We have uncovered a kinematic substructure at \\vhel=-349.5$\\pm1.8$~km/s from a spectroscopic field lying in the Ibata et al.\\ (2007) \\streamc. The cold component has $\\sigma_{v_r}=5.1\\pm2.5$ and a narrow range in [Fe/H]=-0.7$\\pm$0.2, which we propose represents a metal-rich component, \\streamcr.  $\\bullet$ We have uncovered a second kinematic substructure in the same field as \\streamcr\\ at \\vhel=-285.6$\\pm1.2$~km/s with $\\sigma_{v_r}=4.3^{+1.7}_{-1.4}$~km/s (non-zero at $>$3$\\sigma$ confidence interval) and a narrow range in [Fe/H]=-1.3$\\pm$0.2, which we propose represents a metal-poor stream, \\streamcp. We demonstrated that this kinematic \\streamcp\\  has a counterpart in a spatially offset metal-poor region of \\streamc\\ in Ibata et al.\\ (2007).  $\\bullet$ We plausibly detect both \\streamcr\\  and \\streamcp\\ at a position $\\sim$30kpc further north along the structure, with no detectable velocity gradiant for \\streamcr, and a measured velocity gradient of $\\sim$40~km/s for \\streamcp. %the majority of stars even in 305TaS must give rise to the photometric contrast, and thus % it is likely the stars of similar Fe/H in the -400km//s range are the continuation of this stream.  $\\bullet$ We were unable to identify kinematic substructure unambiguously associated to \\streamd\\  from our serendipitous spectroscopic pointing, however subsequent spectroscopy well centered in the \\streamd\\ identifies a likely cold kinematic structure which has a viable counterpart in the serendipitous pointing. We propose a kinematic detection of \\streamd\\ at \\vhel=-390.5~km/s with $\\sigma_{v_r}=4.2$~km/s.  $\\bullet$ Spectroscopy near the edges of \\streama\\ and \\streamb\\ suggest a likely kinematic detection for \\streamb\\ with $v_{\\rm hel}\\sim-330$~km/s, $\\sigma_{v, corr}\\sim6.9$\\,km/s, and a kinematic detection of \\streama\\ at $v_{\\rm hel}\\sim-172$~km/s, $\\sigma_{v_r}\\sim12.5$\\,km/s. Neither spectroscopic pointing in these streams is ideally placed, and additional spectroscopic observations are well motivated to further constrain the kinematics of these structures.  $\\bullet$ The extended cluster EC4 lies in the \\streamc\\ region, with kinematics (v$_{\\rm hel}$=-285\\,km/s) %, $\\sigma_{v, corr}\\sim3$~km/s), and metallicity ([Fe/H]=-1.4) which suggest it is related to the more metal-poor stream \\streamcp. %The resolved velocity dispersion of \\ec4\\ suggests the possible presence of a sizable dark matter component  (Collins et al.\\ 2008a). %with M/L$\\sim$35, although the results are clearly preliminary based on only 8 member stars. \\ec4\\ could be the progenitor of the metal-poor \\streamcp\\ (somewhat unlikely given the apparent stellar mass difference between the stream and \\ec4), or it may simply be a structure carried along by the disrupted stream progenitor. In this case, and if \\ec4\\ has a sizable dark matter component, we have in fact detected the very first {\\it sub-sub-halo} (i.e.\\ a galaxy that was bound to a satellite galaxy), possibly explaining its small (r$_c$=30~pc) size.   $\\bullet$ By explicitly removing stars belonging to the streams by their kinematics we can assess the underlying M31 stellar halo density and metallicity on the minor axis. This contrasts the purely photometric approach where Galactic contamination and stars belonging to stream substructures can only be removed statistically. Our resulting halo measurement has so few stars as to be highly uncertain statistically, however, it does reveal the general power of kinematic analysis of the halo population for future endeavors. The fraction of background halo stars in these stream fields suggests the conclusion that stellar halos are largely made up of multiple kinematically cold streams."
    },
    "0808/0808.1369_arXiv.txt": {
        "abstract": "We suggest a scenario where the three light quark flavors are sequentially deconfined under increasing pressure in cold asymmetric nuclear matter as\\eg in neutron stars. The basis for our analysis is a chiral quark matter model of Nambu--Jona-Lasinio (NJL) type with diquark pairing in the spin-1 single flavor (CSL), spin-0 two flavor (2SC) and three flavor (CFL) channels. We find that nucleon dissociation sets in at about the saturation density, $n_0$, when the down-quark Fermi sea is populated (d-quark dripline) due to the flavor asymmetry induced by $\\beta$-equilibrium and charge neutrality. At about $3n_0$ u-quarks appear and a two-flavor color superconducting (2SC) phase is formed. The s-quark Fermi sea is populated only at still higher baryon density, when the quark chemical potential is of the order of the dynamically generated strange quark mass. We construct two different hybrid equations of state (EoS) using the Dirac-Brueckner Hartree-Fock (DBHF) approach and the EoS by Shen \\etal in the nuclear matter sector. The corresponding hybrid star sequences have maximum masses of, respectively, 2.1 and 2.0 M$_\\odot$. Two- and three-flavor quark-matter phases exist only in gravitationally unstable hybrid star solutions in the DBHF case, while the Shen-based EoS produce stable configurations with a 2SC phase component in the core of massive stars. Nucleon dissociation due to d-quark drip at the crust-core boundary fulfills basic criteria for a deep crustal heating process which is required to explain superbusts as well as cooling of X-ray transients. ",
        "introduction": "The phenomenology of compact stars is intimately connected to the EoS of matter at densities well beyond the nuclear saturation density, $n_0=0.16$ fm$^{-3}$. Compact stars are therefore natural laboratories for the exploration of baryonic matter under extreme conditions, complementary to those created in terrestrial experiments with atomic nuclei and heavy-ion collisions. Recent results derived from observations of compact stars provide serious constraints on the nuclear EoS, see \\cite{Klahn:2006ir} and references therein. A stiff EoS at high density is needed to explain the high compact-star masses, $M\\sim 2.0$~M$_\\odot$, reported for some low-mass X-ray binaries (LMXBs)\\eg 4U 1636-536 \\cite{Barret:2005wd}, and the large radius, $R > 12$~km, of the isolated neutron star RX J1856.5-3754 (shorthand: RX J1856) \\cite{Trumper:2003we}. Another example is EXO 0748-676, an LMXB for which the compact-star mass {\\it and} radius have been constrained to $M\\ge 2.10\\pm 0.28$~M$_\\odot$ and $R \\ge 13.8 \\pm 0.18$~km \\cite{Ozel:2006km}. However, the status of the results for the latter object is unclear, because the gravitational redshift $z=0.35$ observed in the X-ray burst spectra \\cite{Cottam:2002} has not been confirmed, despite numerous attempts. While compact-star phenomenology apparently points towards a stiff EoS at high density, heavy-ion collision data for kaon production \\cite{Fuchs:2005zg} and elliptic flow \\cite{Danielewicz:2002pu} set an upper limit on the stiffness of the EoS \\cite{Klahn:2006ir}.  A key question regarding the structure of matter at high density is whether a phase transition to quark matter occurs inside compact stars, and whether it is accompanied by unambiguous observable signatures. It has been argued  that the observation of a compact star with high mass and large radius, likewise reported for EXO 0748-676, would be incompatible with a quark core \\cite{Ozel:2006km}, because the softening of  the EoS due to quark deconfinement lowers the maximum mass and the radius in comparison to the nuclear matter case. However, Alford \\etal \\cite{Alford:2006vz} have demonstrated with a few counter examples that quark matter cannot be excluded by this argument. In particular, for a recently developed hybrid star EoS \\cite{Klahn:2006iw}, based on the DBHF approach in the nuclear sector and a three-flavor chiral quark model \\cite{Blaschke:2005uj}, stable hybrid stars in the mass range from 1.2 M$_\\odot$ up to 2.1 M$_\\odot$ have been obtained, in accordance with modern mass-radius constraints, see also \\cite{Blaschke:2007ri}. In that model, a sufficiently low critical density for quark deconfinement was achieved with a strong diquark coupling, while a sufficient stiffness for a high maximum mass of the compact star sequence was obtained with a repulsive vector meanfield in the quark matter sector. The corresponding hybrid EoS for symmetric matter was shown to fulfill the constraints from elliptic flow data in heavy-ion collisions. In the present work we want to discuss a new scenario of quark deconfinement, which could play an important role in asymmetric matter, in particular for the phenomenology of compact stars.  Chiral quark models of the NJL type with dynamical chiral symmetry breaking have the property that the restoration of this symmetry (and the related quark deconfinement) at zero temperature is flavor specific. When solving the gap and charge-neutrality equations selfconsistently one finds that the chiral symmetry restoration for a given flavor occurs when the chemical potential of that flavor reaches a critical value that is approximately equal to the dynamically generated quark mass, $\\mu_{f}=\\mu_c \\approx m_{f}$, where $f=u, d, s$. In asymmetric matter the quark chemical potentials are different. Consequently, the NJL model behavior suggests that the critical density of deconfinement is flavor dependent, see Fig. \\ref{f:phases}. \\begin{figure} % \\begin{tabular}{ll} \\includegraphics[angle=0,width=0.5\\textwidth]{d-drip2.eps}& \\includegraphics[angle=0,width=0.4\\textwidth,clip=]{CSL-gaps.eps} \\end{tabular} \\caption{Left panel: Chemical potentials of up and down quarks (strange quark sector not shown). With increasing quark chemical potential $\\mu_q=(\\mu_u+\\mu_d)/2$ in isospin asymmetric matter the quark flavors pass sequentially the threshold ($\\mu_c$) for chiral symmetry restoration (deconfinement), which entails nucleon dissociation. Right panel: Solution of the NJL gap equations under isospin asymmetry.} \\label{f:phases} \\end{figure} In this approach the down quark flavor is the first to drip out of nucleons as the density increases, followed by the up quark flavor and eventually also by strange quarks. This behavior is absent in simple and commonly applied thermodynamic bag model equations of state since they are essentially flavor blind.  Under the $\\beta$-equilibrium condition in compact stars the chemical potentials of quarks and electrons are related by $\\mu_d=\\mu_s$ and $\\mu_d=\\mu_u+\\mu_e$. The mass difference between the strange and the light quark flavors $m_s \\gg m_u, m_d$ has two consequences: (1) the down and strange quark densities are different, so charge neutrality requires a finite electron density and, consequently, (2) $\\mu_d>\\mu_u$. When increasing the baryochemical potential, the d-quark chemical potential is therefore the first to reach the critical value $\\mu_c$ where the chiral symmetry gets (approximately) restored in a first-order transition and a finite density of d-quarks appears. Due to the finite value of $\\mu_e$, the u-quark chemical potential is still below $\\mu_c$ while the s-quark density is zero due to the high s-quark mass. A {\\it single-flavor} d-quark phase therefore forms in co-existence with the positively charged nuclear-matter medium.  Why has this interesting scenario been left unnoticed? On the one hand, bag models, which are commonly used to describe quark matter in compact star interiors cannot address sequential deconfinement. On the other hand, the single-flavor d-quark phase is negatively charged and cannot be neutralized in a purely leptonic background. This was a reason to disregard it in dynamical approaches like the NJL models. In the following we discuss the single-flavor phase for the first time under the natural assumption that the neutralizing background is nuclear matter. Since nucleons are bound states of quarks, the physical context in which such a mixed phase of nucleons and free d-quarks occurs is that of the dissociation of nucleonic bound states of quarks (Mott effect). ",
        "conclusions": "In this contribution we have suggested a new quark-nuclear hybrid EoS for compact star applications that fulfills modern observational constraints from compact stars. Due to isospin asymmetry, down-quarks may ``drip out'' from nucleons and form a single-flavor color superconducting (CSL) phase that is mixed with nuclear matter already at the crust-core boundary in compact stars. The CSL phase has interesting cooling and transport properties that are in accordance with constraints from the thermal and rotational evolution of compact stars. It remains to be investigated whether this new compact star composition could lead to unambiguous observational consequences. We conjecture that the d-quark drip may serve as an effective deep crustal heating mechanism for the explanation of the puzzling superburst phenomenon and the cooling of X-ray transients."
    },
    "0808/0808.3464_arXiv.txt": {
        "abstract": "{ We present the AGILE gamma-ray observations of the field containing the puzzling gamma-ray source 3EG J1835+5918. This source is one of the most remarkable unidentified EGRET sources. } { An unprecedentedly long AGILE monitoring of this source yields important information on the positional error box, flux evolution, and spectrum. }{3EG J1835+5918 has been in the AGILE field of view several times in 2007 and 2008 for a total observing time of 138 days from 2007 Sept 04 to 2008 June 30 encompassing several weeks of continuous coverage.}{With an exposure time approximately twice that of EGRET, AGILE confirms the existence of a prominent gamma-ray source (\\sagile) at a position consistent with that of EGRET, although with a remarkably lower  average flux value for photon energies greater than 100 MeV. {A 5-day bin temporal analysis of the whole data set of \\sagile~ shows some evidence for variability of the gamma-ray flux. } The source spectrum between 100 MeV and 1 GeV can be fitted with a power law with photon index in the range 1.6-1.7, fully consistent with the EGRET value.} { The faint X-ray source RX J1836.2+5925 that has been proposed as a possible counterpart of \\s$\\;$ is well within the AGILE error box. Future continuous monitoring (both by AGILE and GLAST) is needed to confirm the gamma-ray flux variability and to unveil the source origin, a subject that is  currently  being pursued through a multiwavelength search for counterparts. } ",
        "introduction": " ",
        "conclusions": "The AGILE observations and monitoring of \\sagile$\\;$ adds relevant information about this puzzling source. Postponing an account of our multifrequency observations of the region to a forthcoming paper, we briefly outline here a few important points regarding the search for a counterpart. We confirm a point already noticed by other authors, i.e., the absence of a blazar or any other relatively bright radio  sources  in the field containing the AGILE (and EGRET) error box. The analysis of  Mattox et al. (2001) shows a 100 mJy source (B1834+5904) positioned at the 99.5$\\%$ probability contour ($12'.7$ from the EGRET centroid position and $21'.8$ from the AGILE centroid position). By using the NED, SIMBAD, and  the  Massaro et al. (2007) catalogues, we confirm the absence of a radio-loud blazar  in the revised AGILE error box.  Outside this error box, we notice the existence of  RGB J1841+591, a BL Lac type object positioned at $43'.5$ from  the AGILE centroid position. Even though occasional contributions from this object cannot be excluded in our analysis of \\agile, we emphasise that the AGILE integrated flux positioning at the 95\\% contour level is clearly not consistent with any substantial contribution from RGB~J1841+591.  We also notice that the faint X-ray source RX J1836.2+5925 that was proposed as a possible counterpart is well within the AGILE error box of \\sagile. A hint of variability of this source was noticed when comparing two HRI ROSAT observations taken almost three years apart \\citep{Mirabal_2001}, prompting a claim for the discovery of a new class of compact gamma-ray sources \\citep{Mirabal_2000}. However, subsequent Chandra X-ray observations showed a practically constant X-ray flux of RX J1836.2+5925, and weakened the variability claim in favour of a scenario encompassing a constant Geminga-like source.  We note here that our group observed the region containing RX J1836.2+5925 with the SWIFT XRT and XMM-Newton during the period May-June 2008. Analysis of these data and a full discussion of the \\sagile$\\;$ counterpart problem will be presented elsewhere.  Future continuous monitoring (both by AGILE and GLAST) is needed to confirm the gamma-ray flux variability and to unveil the source origin."
    },
    "0808/0808.1461_arXiv.txt": {
        "abstract": "Magnetic fields in the early universe can significantly alter the thermal evolution and the ionization history during the dark ages. This is reflected in the $21$ cm line of atomic hydrogen, which is coupled to the gas temperature through collisions at high redshifts, and through the Wouthuysen-Field effect at low redshifts. We present a semi-analytic model for star formation and the build-up of a Lyman $\\alpha$ background in the presence of magnetic fields, and calculate the evolution of the mean $21$ cm brightness temperature and its frequency gradient as a function of redshift. We further discuss the evolution of linear fluctuations in temperature and ionization in the presence of magnetic fields and calculate the effect on the $21$ cm power spectrum. At high redshifts, the signal is increased compared to the non-magnetic case due to the additional heat input into the IGM from ambipolar diffusion and the decay of MHD turbulence.  At lower redshifts, the formation of luminous objects and the build-up of a Lyman $\\alpha$ background can be delayed by a redshift interval of $10$ due to the strong increase of the filtering mass scale in the presence of magnetic fields. This tends to decrease the $21$ cm signal compared to the zero-field case. In summary, we find that $21$ cm observations may become a promising tool to constrain primordial magnetic fields. ",
        "introduction": "{Observations of the $21$ cm fine structure line of atomic hydrogen have the potential to become an important means of studying} the universe at early times, during and even before the epoch of reionization. This possibility was suggested originally by \\citet{Purcell}, and significant process in instrumentation and the development of radio telescopes has brought us close to the first observations from radio telescopes like LOFAR\\footnote{http://www.lofar.org}. While one of the main purposes is to increase our understanding of cosmological reionization \\citep{Bruyn}, a number of further exciting applications have been suggested in the mean time. \\citet{Loeb} demonstrated how $21$ cm measurements can probe the thermal evolution of the IGM at a much earlier time, at redshifts of $z\\sim200$. \\citet{Barkana} suggested a method that separates physical and astrophysical effects and thus allows to probe the physics of the early universe. \\citet{FurlanettoDM} showed how the effect of dark matter annihilation and decay would be reflected in the $21$ cm line, and effects of primordial magnetic fields have been considered by \\citet{Tashiro}.\\\\ \\\\  Indeed, primordial magnetic fields can affect the early universe in various ways. The thermal evolution is significantly altered by ambipolar diffusion heating and decaying MHD turbulence \\citep{Sethi05, Sethi08, Schleicher}. \\citet{Kim} calculated the effect of the Lorentz force on structure formation and showed that additional power is present on small scales in the presence of primordial magnetic fields. It was thus suggested that reionization occurs earlier in the presence of primordial magnetic fields \\citep{Sethi05, TashiroReion}. However, as pointed out by \\citet{Gnedin}, the characteristic mass scale of star forming halos, the so-called filtering mass, increases signficantly when the temperature is increased. For comoving field strengths of $\\sim1$ nG, we found that the filtering mass scale is shifted to scales where the power spectrum is essentially independent of the magnetic field \\citep{Schleicher}. Reionization is thus delayed in the presence of primordial magnetic fields. We further found upper limits of the order $1$ nG, based on the Thomson scattering optical depth measured by WMAP 5 \\citep{Komatsu, WMAPAngular} and the requirement that reionization ends at $z\\sim6$ \\citep{Becker}.   The presence of primordial magnetic fields can have interesting implications on first star formation as well. In the absence of magnetic fields, \\citet{Abel} and \\citet{Bromm} suggested that the first stars should be very massive, perhaps with $\\sim100$ solar masses.  \\citet{Clark} and \\citet{Omukai} argued that in more massive and perhaps metal-enriched galaxies, fragmentation should be more effective and lead to the formation of rather low-mass stars, due to a stage of efficient cooling \\citep{Omukai}. {Magnetic fields may change this picture and reduce the stellar mass by triggering jets and outflows \\citep{Silk}. However, simulations by \\citep{Machida} show that the change in mass is of the order $10\\%$.} \\\\ \\\\  $21$ cm measurements can try to adress {primordial magnetic fields} in two ways: During the dark ages of the universe, at redshifts $z\\sim200$ well before the formation of the first stars, the spin temperature of hydrogen is coupled to the gas temperature via collisional de-excitation by hydrogen atoms \\citep{Allison, Zygelman} and free electrons \\citep{Smith, FF07}, constituting a probe at very early times. While collisional de-excitation becomes inefficient due to the expansion of the universe, the first stars will build up a Lyman $\\alpha$ background that will cause a deviation of the spin temperature from the radiation temperature by the Wouthuysen-Field effect \\citep{Wouthuysen, Field}. As primordial magnetic fields may shift the onset of reionization, the onset of this coupling constitutes an important probe on the presence of such fields. We adress these possibilities in the following way: In \\S\\ref{IGM}, we review our treatment of the IGM in the presence of primordial magnetic fields. The evolution of the $21$ cm background and the role of Lyman $\\alpha$ photons is discussed in \\S\\ref{21cm}. The evolution of linear perturbations in temperature and ionization is calculated in \\ref{fluctuations}. The results for the power spectrum are given in \\S\\ref{power}, and the results are further discussed in \\S\\ref{discussion}. ",
        "conclusions": "\\label{discussion} In the previous sections, we have presented a semi-analytic model describing the post-reionization universe and reionization in the context of primordial magnetic fields, and calculated the consequences for the mean $21$ cm brightness fluctuation and the large-scale power spectrum.  formation of We identify two regimes in which primordial magnetic fields can influence effects measured with $21$ cm telescopes. At low redshifts, primordial magnetic fields tend to delay reionization and the build-up of a Lyman $\\alpha$ background, thus shifting the point where the signal is at its maximum, and changing the amplitude of the $21$ cm power spectrum. The first $21$ cm telescopes like LOFAR and others will focus mainly on the redshift of reionization and can thus probe the epoch when a significant Lyman $\\alpha$ background builds up. As our understanding of the first stars increases due to advances in theoretical modeling or due to better observational constraints, this may allow us  to determine whether primordial magnetic fields are needed to delay the build-up of Lyman $\\alpha$ photons or not. As demonstrated above, comoving field strengths of the order $1$ nG can delay the build-up of a Lyman $\\alpha$ background by $\\Delta z\\sim 10$, which is significantly stronger than other mechanisms that might delay the formation of luminous objects. Lyman Werner feedback is essentially self-regulated and never leads to a significant suppression of star formation \\citep{Johnson, JohnsonO}, X-ray feedback from miniquasars is strongly constrained from the observed soft X-ray background \\citep{Salvaterra}, and significant heating from dark matter decay or annihilation would be accompanied by a significant amount of secondary ionization, resulting in a too large optical depth. On the other hand, an important issue is the question regarding the first sources of light. As shown in \\citet{Schleicher}, massive Pop. III stars are needed to provide the correct reionization optical depth. On the other hand, an additional population of less massive stars with an IMF according to \\citep{Scalo,Kroupa,Chabrier} might be present. Such a population would emit more photons between the Lyman $\\alpha$ and $\\beta$ line per stellar baryon \\citep{Leitherer}, and could thus shift the build-up of a Lyman $\\alpha$ background to an earlier epoch. The onset of efficient coupling via the Wouthuysen-Field effect thus translates into a combined constraint on the stellar population and the strength of primordial magnetic fields.\\\\ \\\\  A further effect occurs at high redshifts, where the additional heat input from magnetic fields due to ambipolar diffusion and the decay of MHD turbulence increases the $21$ cm signal. As gas decouples earlier from the radiation field, the difference between gas and radiation temperature is larger and collisions are more effective in coupling the spin temperature to the gas temperature. The determination of the $21$ cm signal from this epoch is certainly challenging, as the foreground emission corresponds to temperatures which are higher than the expected $21$ cm brightness temperature by several orders of magnitude. However, as pointed out in other works \\citep{DiMatteo, OhMack, Zaldarriaga, Sethi, Yu}, the foreground emission is expected to be featureless in frequency, which may allow for a sufficiently accurate subtraction. In this context, it may also to help to focus on the frequency gradient of the mean brightness temperature, rather than the $21$ cm brightness temperature itself. Upcoming long-wavelength experiments such as LOFAR, 21CMA (former PAST)\\footnote{http://web.phys.cmu.edu/~past}, MWA\\footnote{http://web.haystack.mit.edu/arrays/MWA/index.html}, LWA\\footnote{http://lwa.unm.edu} and SKA\\footnote{http://www.skatelescope.org}  may thus detect the additional heat from primordial magnetic fields in the neutral gas, or otherwise set new upper limits on primordial magnetic fields, perhaps down to $B_0\\sim0.1$ nG. Like the 21 cm transition of hydrogen, rotational and ro-vibrational transitions of primordial molecules may create interesting signatures in the CMB as well\\citep{Schleicher2}, which may provide a further test of the thermal evolution during the dark ages."
    },
    "0808/0808.1182_arXiv.txt": {
        "abstract": "We introduce the {\\it SMC in space and time}, a large coordinated space and ground-based program to study star formation processes and history, as well as variable stars, structure, kinematics and chemical evolution of the whole SMC. Here, we present the Colour-Magnitude Diagrams(CMDs) resulting from HST/ACS photometry, aimed at deriving the star formation history (SFH) in six fields of the SMC. The fields are located in the central regions, in the stellar halo, and in the wing toward the LMC. The CMDs are very deep, well beyond the oldest Main Sequence Turn-Off, and will allow us to derive the SFH over the entire Hubble time. ",
        "introduction": "The Small Magellanic Cloud (SMC) is the closest late-type dwarf and has many properties similar to those of the vast majority of this common class of galaxies. Its current metallicity (Z$\\simeq$0.004 in mass fraction, as derived from HII regions and young stars) is typical of dwarf irregular and Blue Compact Dwarf (BCD) galaxies, the least evolved systems, hence the most similar to primeval galaxies. Its mass (between 1 and 5 $\\times 10^9 M_{\\odot}$, e.g. \\cite[Kallivayalil et al. 2006]{K06} and references therein) is at the upper limit of the range of masses typical of late-type dwarfs. These characteristics, combined with its proximity, make the SMC the natural benchmark to study the evolution of late-type dwarf galaxies. Moreover, its membership to a triple system allows detailed studies of interaction-driven modulations of the star formation activity.  A wealth of data on the SMC are available in the literature, although not as much as for its bigger companion, the LMC. Yet, much more are needed for a better understanding of how the SMC has formed and evolved. We have thus embarked on a long-term project to study the evolution of the SMC in space and time. Our project plans to exploit the high performances in depth, resolving power or large field of view of current and forthcoming, space and ground based, telescopes, such as HST, VLT, Spitzer, SALT and VST. ",
        "conclusions": ""
    },
    "0808/0808.2462_arXiv.txt": {
        "abstract": "{ In supersymmetric models with a long-lived stau being the lightest Standard Model superpartner, the stau abundance during primordial nucleosynthesis is tightly constrained. Considering the complete set of stau annihilation channels in the minimal supersymmetric Standard Model (MSSM) with real parameters for scenarios in which sparticle coannihilations are negligible, we calculate the decoupling of the lighter stau from the primordial plasma and identify processes which are capable to deplete the resulting stau abundance significantly. We find particularly efficient stau annihilation at the resonance of the heavy CP-even Higgs boson and for a lighter stau with a sizeable left--right mixing due to enhanced stau-Higgs couplings. Even within the constrained MSSM, we encounter both effects leading to exceptionally small values of the resulting stau abundance. Prospects for collider phenomenology are discussed and possible implications of our findings are addressed with emphasis on gravitino dark matter scenarios. }    \\begin{document} ",
        "introduction": "\\label{sec:introduction}  The appearance of the lighter stau $\\stauone$ as the lightest Standard Model superpartner---or lightest ordinary superpartner (LOSP)---is a commonplace occurrence even in supersymmetric (SUSY) models with restrictive assumptions on the SUSY breaking sector such as the constrained minimal supersymmetric Standard Model (CMSSM).  If the lightest supersymmetric particle (LSP) is assumed to be the LOSP, this parameter region is not considered because of severe upper limits on the abundance of massive stable charged particles~\\cite{Yao:2006px}. However, for example, in axino/gravitino LSP scenarios~\\cite{Pagels:1981ke,Borgani:1996ag,Bonometto:1989vh,Covi:1999ty,Steffen:2007sp} and in scenarios with broken R parity~\\cite{Allanach:2003eb,Allanach:2006st,Allanach:2007vi,Takayama:2000uz,Buchmuller:2007ui},% \\footnote{In this work we assume that R-parity is conserved.} the $\\stau$ LOSP becomes unstable and thereby a viable option. Indeed, supersymmetric models with a long-lived $\\stauone$ LOSP are particularly promising for collider phenomenology~\\cite{Drees:1990yw,Nisati:1997gb,Feng:1997zr,Ambrosanio:2000ik,Buchmuller:2004rq,Hamaguchi:2004df,Feng:2004yi,Brandenburg:2005he,DeRoeck:2005bw,Martyn:2006as,Steffen:2006hw,Ellis:2006vu,Hamaguchi:2006vu}: Since the $\\stauone$ LOSP could escape the collider detector as a quasi-stable muon-like particle, it can be associated with signatures that are very different from the excess in missing energy expected in neutralino LSP scenarios.  In the early Universe the negatively charged LOSP $\\stauone$'s and the associated positively charged anti-staus $\\stauone^*$'s were in thermal equilibrium for temperatures of $T>m_{\\stauone}/20\\gtrsim T_{\\freezeout}$. At $T_{\\freezeout}$, the annihilation rate of the (by then) non-relativistic $\\stauone$'s becomes smaller than the Hubble rate so that they decouple from the thermal plasma.  Thus, for $T\\lesssim T_{\\freezeout}$, their yield $Y_{\\stau}\\equiv(n_{\\stauone}+n_{\\stauone^*})/s$ is given approximately by $\\Ystau\\approx Y^{\\equil}_{\\stau}(T_{\\freezeout})$, where $n_{\\stau}^{(\\equil)}\\equiv n_{\\stauone}^{(\\equil)}+ n_{\\stauone^*}^{(\\equil)}$ is the (equilibrium) number density of both $\\stauone$ and $\\stauone^*$ and $s=2\\pi^2\\,g_{*S}\\,T^3/45$ the entropy density with \\gstarS\\ effective degrees of freedom. This thermal relic abundance $\\Ystau$ is subject to cosmological constraints in SUSY scenarios with a long-lived $\\stauone$ LOSP:  \\begin{itemize}  \\item In axino/gravitino LSP scenarios, $\\Ystau$ governs the non-thermally produced (NTP) relic density of axino/gravitino dark matter that originates from $\\stauone$ decays~\\cite{Borgani:1996ag,Covi:1999ty} $\\Omega_{\\widetilde{\\mathrm{a}}/\\gravitino}^{\\NTP} h^2 = m_{\\widetilde{\\mathrm{a}}/\\gravitino}\\, \\Ystau\\, s(T_0) h^2 / \\rho_{\\mathrm{c}}$ where $m_{\\widetilde{\\mathrm{a}}/\\gravitino}$ denotes the axino/gravitino LSP mass and $\\rho_c/[s(T_0)h^2]=3.6\\times 10^{-9}\\,\\GeV$~\\cite{Yao:2006px}. Thus, the dark matter density $\\Omega_{\\mathrm{dm}}$ which limits $\\Omega_{\\widetilde{\\mathrm{a}}/\\gravitino}^{\\NTP}$ from above implies an upper limit on $\\Ystau$ for a given $m_{\\widetilde{\\mathrm{a}}/\\gravitino}$. This limit can become particularly restrictive in the case of additional sizeable contributions to $\\Omega_{\\mathrm{dm}}$ such as the ones from thermal axino/gravitino production $\\Omega_{\\widetilde{\\mathrm{a}}/\\gravitino}^{\\TP}$~\\cite{Bolz:2000fu,Brandenburg:2004du,Pradler:2006qh,Rychkov:2007uq}. For example, for $m_{\\widetilde{\\mathrm{a}}/\\gravitino}=50~\\GeV$ and $\\Omega_{\\widetilde{\\mathrm{a}}/\\gravitino}^{\\TP}=0.99\\,\\Omega_{\\mathrm{dm}}$ ($0.9\\,\\Omega_{\\mathrm{dm}}$), one finds $\\Ystau<10^{-13}$ ($10^{-12}$); cf.\\ Fig.~13 of Ref.~\\cite{Steffen:2006hw}.  \\item For $\\stauone$ decays during/after big bang nucleosynthesis (BBN), the Standard Model particles emitted in addition to the axino/gravitino LSP can affect the abundances of the primordial light elements.  This leads to upper limits on $\\xi_{\\mathrm{em}/\\mathrm{had}}\\equiv \\epsilon_{\\mathrm{em}/\\mathrm{had}}\\,\\Ystau$ that depend on the stau lifetime $\\tau_{\\stauone}$~\\cite{Cyburt:2002uv,Kawasaki:2004qu,Jedamzik:2006xz}. Here $\\epsilon_{\\mathrm{em}/\\mathrm{had}}$ denotes the (average) electromagnetic/hadronic energy emitted in a single $\\stauone$ decay, which can be calculated with particle physics methods for a given model. Accordingly, the BBN constraints on $\\xi_{\\mathrm{em}/\\mathrm{had}}$ can be translated into upper limits on $\\Ystau$; cf.\\ Fig.~12 of Ref.~\\cite{Steffen:2006hw} (and Figs.~14 and~15 of Ref.~\\cite{Kawasaki:2008qe}) for associated $\\Ystau$ limits in gravitino LSP scenarios, which can be as restrictive as $\\Ystau<10^{-14}$ ($10^{-15}$).  \\item The mere presence of the negatively charged $\\stauone$'s at cosmic times of $t\\gtrsim 5\\times 10^3\\,\\seconds$ can lead to ($^4$He$\\,\\stauone$) and ($^8$Be$\\,\\stauone$) bound states and can thereby allow for catalyzed BBN (CBBN) of $^6$Li and $^9$Be to abundances far above the ones obtained in standard BBN (SBBN)~\\cite{Pospelov:2006sc,Bird:2007ge,Pospelov:2007js,Pospelov:2008ta}. Indeed, confronting the abundances obtained in CBBN with observationally inferred bounds on the primordial abundances of $^9$Be (and $^6$Li) imposes restrictive upper limits of $\\Ystau\\lesssim 2\\times 10^{-15}$ ($2\\times 10^{-15}$\\,--\\,$2\\times 10^{-16})$ for $\\tau_{\\stauone}\\gtrsim 10^5\\,\\seconds$; cf.\\ Fig.~5 in Ref.~\\cite{Pospelov:2008ta} for $n_{\\stauone}=n_{\\stauone^*}$.  \\end{itemize}  For example, in gravitino LSP scenarios with the $\\stauone$ LOSP being the next-to-lightest supersymmetric particle (NLSP) and conserved R-parity, the listed cosmological constraints have been confronted with representative values~\\cite{Asaka:2000zh,Fujii:2003nr}% \\begin{align} \\label{eq:yield-approx} \\Ystau \\simeq (0.4-1.5) \\times 10^{-13} \\left( \\frac{\\mstauone}{100~\\GeV} \\right) \\ , \\end{align} which are in good agreement with the curves in Fig.~1 of Ref.~\\cite{Asaka:2000zh} that have been obtained for the case of a purely `right-handed' $\\stauone\\simeq\\stauR$ NLSP and a bino-like lightest neutralino, $\\neutralino\\simeq\\Bino$, with a mass of $m_{\\Bino}=1.1\\,\\mstauone$.  Thereby, it has been found that the (C)BBN constraints impose the limit $\\taustau\\lesssim 5\\times 10^3\\,\\seconds$~\\cite{Pospelov:2006sc,Pospelov:2008ta} with severe implications in the collider-friendly region of $\\mstauone < 1~\\TeV$: (i)~The $\\tau_{\\stauone}$ limit disfavors the kinematical determination of $\\mgravitino$~\\cite{Steffen:2006wx} and thereby both the determination of the Planck scale at colliders~\\cite{Buchmuller:2004rq} and the method proposed to probe the maximum reheating temperature $\\TR$ at colliders~\\cite{Pradler:2006qh}. (ii)~Within the CMSSM, the $\\tau_{\\stauone}$ limit implies an upper limit on the reheating temperature of $\\TR\\lesssim 10^7~\\GeV$~\\cite{Pradler:2006hh,Pradler:2007is,Pradler:2007ar} that disfavors the viability of thermal leptogenesis with hierarchical heavy Majorana neutrinos~\\cite{Fukugita:1986hr,Davidson:2002qv,Buchmuller:2004nz,Blanchet:2006be,Antusch:2006gy}. (iii)~The $\\tau_{\\stauone}$ limit can point to a CMSSM mass spectrum which will be difficult to probe at the Large Hadron Collider (LHC)~\\cite{Cyburt:2006uv,Pradler:2006hh,Pradler:2007is,Pradler:2007ar}. Indeed, the $\\tau_{\\stauone}$ limit can be relaxed only with a significant reduction of~(\\ref{eq:yield-approx}) which has been presented explicitly so far only for non-standard cosmological scenarios with a very low value of~$\\TR$~\\cite{Takayama:2007du} or with late-time entropy production after $\\stauone$ decoupling and before BBN~\\cite{Pradler:2006hh,Hamaguchi:2007mp}.% \\footnote{Note that some implications of the $\\tau_{\\stauone}$ limit can be evaded not only by relaxing it but also by respecting it, e.g., R-parity violation can lead to $\\tau_{\\stauone}<5\\times 10^3\\,\\seconds$~\\cite{Buchmuller:2007ui}.}  In this work we calculate the decoupling of the lighter stau from the primordial plasma by taking into account the complete set of stau annihilation channels in the MSSM with real parameters for SUSY spectra for which sparticle coannihilation is negligible. Using our own code for the computation of the resulting thermal relic stau abundance $\\Ystau$, we examine explicitly (i)~the effect of left--right mixing of the lighter stau, (ii)~the effect of large stau--Higgs couplings, and (iii)~stau annihilation at the resonance of the heavy CP-even Higgs boson $H^0$.  We consider both the ``phenomenological MSSM'' (pMSSM) (see, e.g.,~\\cite{Djouadi:2002ze}) in which the soft SUSY breaking parameters can be set at the weak scale, and the CMSSM, in which the gaugino masses, the scalar masses, and the trilinear scalar couplings are assumed to take on the respective universal values $\\monetwo$, $\\mzero$, and $A_0$ at the scale of grand unification $\\mgut\\simeq 2\\times 10^{16}\\,\\GeV$. Within the framework of the pMSSM, we show examples in which $\\Ystau$ can be well below $10^{-15}$.  Even within the CMSSM, we encounter regions with exceptionally small values of $\\Ystau\\lesssim 2\\times 10^{-15}$. The implications of these findings are discussed for scenarios with the gravitino LSP and the stau NLSP.  We also address the viability of a $\\stauone$--$\\stauone^*$ asymmetry.  Remarkably, we find that key quantities for the significant $\\Ystau$ reduction could be probed at both the LHC and the International Linear Collider (ILC).  A calculation of the thermal relic abundance of long-lived staus has also been part of a recent thorough study~\\cite{Berger:2008ti} which focusses on gauge interactions and on the effect of Sommerfeld enhancement. In contrast, the most striking findings of our study---in which Sommerfeld enhancement is not taken into account---are related to the Higgs sector of the MSSM.  At this point, we should also stress that the \\texttt{micrOMEGAs} code~\\cite{Belanger:2001fz,Belanger:2004yn,Belanger:2007zz,Belanger:2008sj} allows for sophisticated calculations of the thermal relic stau abundance also in regions in which coannihilation effects become important.  In fact, \\texttt{micrOMEGAs} has already been applied in several studies to calculate $\\Ystau$~\\cite{Fujii:2003nr,Pradler:2006hh,Kersten:2007ab,Pradler:2007ar,Berger:2008ti}. In this paper, we also work with \\texttt{micrOMEGAs} to cross check the results of our own $\\Ystau$ calculation and to calculate $\\Ystau$ in parameter regions in which sparticle coannihilations become relevant.  The outline of this paper is as follows.  In the next section we review basic properties of the staus to introduce our notations and conventions for the stau mixing angle.  Section~\\ref{sec:prim-ann} explains the way in which we calculate $\\Ystau$ and provides the complete list of stau annihilation channels.  In Sect.~\\ref{sec:dependence} we analyze the dependence of the most relevant stau annihilation channels on the stau mixing angle.  Effects of large stau--Higgs couplings and stau annihilation at the $H^0$ resonance are studied in Sects.~\\ref{sec:enhanc-coupl-higgs} and~\\ref{sec:reson-annih}, respectively.  The viability of a $\\stauone$--$\\stauone^*$ asymmetry is addressed in Sect.~\\ref{sec:comm-stau-stau}.  In Sect.~\\ref{sec:annih-chann} we present exemplary parameter scans within the CMSSM that exhibit exceptionally small $\\Ystau$ values.  Potential collider phenomenology of the parameter regions associated with those exceptional relic abundances and potential implications for gravitino dark matter scenarios are discussed in Sects.~\\ref{sec:collider} and~\\ref{sec:gravitino}, respectively. ",
        "conclusions": "\\label{sec:conclusions}  Supersymmetric models with a long-lived stau $\\stauone$ being the lightest Standard Model superpartner are well-motivated and very attractive in light of potentially striking signatures at colliders. For a standard thermal history with primordial temperatures $T>\\mstauone/20>\\Tf$---which is the working hypothesis is this work---the long-lived $\\stauone$ becomes an electrically charged thermal relic whose abundance can be restricted by cosmological constraints.  We have carried out a thorough study of primordial stau annihilation and the associated thermal freeze out. Taking into account the complete set of stau annihilation channels within the MSSM with real parameters for cases with negligible sparticle coannihilation, the resulting thermal relic $\\stauone$ yield $\\Ystau$ has been examined systematically. While related earlier studies focussed mainly on the $\\stauone\\simeq\\stauR$ case~\\cite{Asaka:2000zh,Fujii:2003nr,Pradler:2006hh,Berger:2008ti}, we have investigated cases in which $\\stauone$ contains a significant admixture of $\\stauL$ including the maximal mixing case and $\\stauone\\simeq\\stauL$.  We find that the variation of the stau mixing angle $\\thetastau$ does affect the relative importance of the different annihilation channels significantly but not necessarily the resulting $\\Ystau$ value for relatively small values of $\\tanb$. By increasing $\\tanb$, however, we encounter a dramatic change of this picture for large absolute values of the Higgs-higgsino mass parameter $\\mu$ and/or of the trilinear coupling $\\Atau$, which are the dimensionful SUSY parameters that govern simultaneously stau left-right mixing and the stau--Higgs couplings: Stau annihilation into $\\hhiggs\\hhiggs$, $\\hhiggs\\Hhiggs$, and $\\Hhiggs\\Hhiggs$ can become very efficient (if kinematically allowed) so that $\\Ystau$ can decrease to values well below $10^{-15}$. The scalar nature of $\\stauone$ allows those parameters to enter directly into the annihilation cross sections. This mechanism has no analogue in calculations of the thermal relic density of the lightest neutralino $\\neutralino$.  The stau--Higgs couplings are crucial also for the second $\\Ystau$ reduction mechanism identified in this work: Even for moderate values of $\\tanb$, we find that staus can annihilate very efficiently into a $\\bquark\\antibquark$ pair via $s$-channel exchange of the heavy CP-even Higgs boson $\\Hhiggs$ provided the MSSM spectrum exhibits the resonance condition $2\\mstauone\\simeq\\mH$. We have shown explicitly that the associated $\\Ystau$ values can be below $10^{-15}$ as well. This mechanism is similar to the one that leads to the reduction of the $\\neutralino$ density in the Higgs funnel region in which neutralino annihilation proceeds at the resonance of the CP-odd Higgs boson $\\Ahiggs$.  We have worked with an effective low energy version of the MSSM to investigate the $\\thetastau$-dependence of $\\Ystau$ and the two $\\Ystau$-reduction mechanisms in a controlled way. In addition, we have shown that the considered effects can be accommodated also with restrictive assumptions on the soft-SUSY breaking sector at a high scale. Within the CMSSM, we encounter both mechanisms each of which leading to $\\Ystau\\simeq 2\\times 10^{-15}$ in two distinct regions of a single $(\\monetwo,\\,\\mzero)$ plane.  We have discussed possibilities to probe the viability of the presented $\\Ystau$-reduction mechanisms at colliders.  While a $\\mH$ measurement pointing to $\\mH\\simeq 2\\mstauone$ would support resonant primordial stau annihilation, studies of Higgs boson production in association with staus, $\\positron\\electron\\,(\\gamma\\gamma) \\to\\stauone\\stauone^*\\hhiggs,\\stauone\\stauone^*\\Hhiggs$ could allow for an experimental determination of the relevant stau--Higgs couplings, for example, at the ILC. Moreover, we have outlined that associated $\\bquark\\antibquark\\hhiggs/\\Hhiggs$ production with $\\hhiggs/\\Hhiggs\\to\\stauone\\stauone^*$ has the potential to allow for a determination of both $\\mh$ and $\\mH$ at the LHC if a SUSY scenario with large $\\tanb$ and large stau--Higgs couplings is realized.  With the obtained small $\\Ystau$ values, even the restrictive constraints associated with CBBN could be respected so that attractive gravitino dark matter scenarios could be revived to be cosmologically viable even for a standard cosmological history. Within this class of models, collider evidence for supergravity, for the gravitino being the LSP, and for high values of the reheating temperatures of up to $10^9\\,\\GeV$ is conceivable, which could thereby accommodate simultaneously the explanation of the cosmic baryon asymmetry provided by thermal leptogenesis and the hypothesis of thermally produced gravitinos being the dark matter in our Universe.  \\bigskip  {\\bf Acknowledgments}  -- We are grateful to T.~Hahn, J.-L.~Kneur, and A.~Pukhov for patient correspondence on their computer codes. Furthermore, we are grateful to T.~Plehn, M.~Pospelov, and Y.Y.Y.~Wong for valuable discussions. This research was supported in part by the DFG cluster of excellence ``Origin and Structure of the Universe.''  \\bigskip  {\\bf Note added} -- Ref.~\\cite{Ratz:2008qh}, in which the potential suppression in the stau yield $\\Ystau$ due to an enhanced annihilation into $\\hhiggs\\hhiggs$ final states is also studied, appeared as this work was being finalized. This paper provides analytic approximations for the stau annihilation cross section into $\\hhiggs\\hhiggs$ and for the associated yield. In addition, results of numerical studies within the CMSSM, the NUHM, and a scenario with non-universal gaugino masses are presented that exhibit parameter regions with extremely small $\\Ystau$ values. In our work also enhanced stau annihilation into $\\hhiggs\\Hhiggs$ and into $\\Hhiggs\\Hhiggs$ and stau annihilation at the $\\Hhiggs$ resonance, which were not considered in~\\cite{Ratz:2008qh}, are discussed. In addition, our work provides a systematic investigation based on a complete set of stau annihilation channels, an outline of the way in which the mechanisms leading to the suppression of $\\Ystau$ can be probed at collider experiments, and a thorough presentation of the potential implications for gravitino dark matter scenarios."
    },
    "0808/0808.0184_arXiv.txt": {
        "abstract": "TeV-blazars are known as prominent non-thermal emitters across the entire electromagnetic spectrum with their photon power peaking in the X-ray and TeV-band. If distant, absorption of $\\gamma$-ray photons by the extragalactic background light (EBL) alters the intrinsic TeV spectral shape, thereby affecting the overall interpretation. Suzaku observations for two of the more distant TeV-blazars known to date, 1ES~1101-232 and 1ES~1553+113, were carried out in May and July 2006, respectively, including a quasi-simultaneous coverage with the state of the art Cherenkov telescope facilities. We report on the resulting data sets with emphasis on the X-ray band, and set into context to their historical behaviour. During our campaign, we did not detect any significant X-ray or $\\gamma$--ray variability. 1ES~1101-232 was found in a quiescent state with the lowest X-ray flux ever measured. The combined XIS and HXD PIN data for 1ES~1101-232 and 1ES~1553+113 clearly indicate spectral curvature up to the highest hard X-ray data point ($\\sim 30$ keV), manifesting as softening with increasing energy. We describe this spectral shape by either a broken power law or a log-parabolic fit with equal statistical goodness of fits. The combined 1ES 1553+113 very high energy spectrum ($90-500$~GeV) did not show any significant changes with respect to earlier observations. The resulting contemporaneous broadband spectral energy distributions of both TeV-blazars are discussed in view of implications for intrinsic blazar parameter values, taking into account the $\\gamma$-ray absorption in the EBL. ",
        "introduction": "Blazars are among the most extreme sources in the high energy sky. They constitute a subclass of the jet-dominated radio-loud population of active galactic nuclei (AGN), which in turn are sub-divided into flat spectrum radio quasars (FSRQ) and BL Lac objects where the absence or depression of strong emission lines characterizes the latter. Their observational properties include non-thermal continuum emission and irregular variability across the whole electromagnetic band, ranging from subhours at gamma-ray energies to week-long time scales in the radio band, often high optical and radio polarization, and a core-dominated radio morphology. In some sources, superluminal motion has been detected, indicating a relativistically enhanced emission region close to the line-of-sight \\cite{Blandford78}. As a consequence of this, any jet emission is highly beamed.  The spectral energy distribution (SED) of blazars shows two broad peaks in the $\\nu F_\\nu$ representation. The low energy hump is generally agreed to stem from synchrotron radiation of relativistic electrons and/or positrons in the jet whereas the origin of the high energy one is still under debate.  Leptonic models explain the complete non-thermal SED as synchrotron and inverse Compton emission from relativistic electrons or pairs which upscatter their self-produced synchrotron radiation or external photon fields. Hadronic models produce the high energy peak either via interactions of relativistic protons with matter \\cite{Pohl00}, ambient photons \\cite[e.g.][]{Mannheim93}, magnetic fields \\cite{Aharonian00}, and/or both, magnetic fields and photons \\cite[e.g.][]{Muecke01,Muecke03}.  Almost all of the 19 blazars (except for BL Lacertae \\cite{Albert07c} and 3C~279 \\cite{Teshima07}) detected reliably to date at the TeV energy belong to the subclass of high-frequency peaked BL Lacs (HBLs), with their two $\\nu F_\\nu$ peaks typically at UV/X-rays and TeV-energies. Until recently blazars have been detected preferentially in their flare state, owing to instrumental limitations. With the contemporary Cherenkov telescopes (such as H.E.S.S., MAGIC, VERITAS) it is possible to study also low activity states (subsequently referred to \"quiet state\").   TeV-blazars have also been used successfully as a powerful tool to probe the extragalactic background light (EBL) at IR/optical energies via the absorption of $\\gamma$-rays in the cosmic diffuse radiation field. The most stringent constraints on the EBL density at $\\sim 2\\mu$m stem from observations with the H.E.S.S. telescope system of the high-redshift, hard-spectrum TeV-blazar 1ES~1101-232 \\cite{Aharonian06a}. The most distant (redshift $z>0.1$) TeV-blazars established to date are 1ES~1101-232 ($z=0.186$), 1ES~1218+304 ($z=0.182$), 1ES~1011+496 ($z=0.212$), 1ES~0347-121 ($z=0.185$), 1ES~0229+200 ($z=0.140$), H~2356-309 ($z=0.165$), H~1426+428 ($z=0.129$), PKS~2155-304 ($z=0.116$), 1ES 1553+113 \\cite[$z>0.25$:][]{Perlman96,Falomo90,Aharonian06b}.   Extreme HBLs, as is the case for most very high energy (VHE) blazars, are characterized by their spectral extension to extremely high particle energies. TeV-blazars have been regularly observed in the past by X-ray instruments, however, not many have been detected so far in the hard X-ray range where the highest energy particles leave their imprints. The sensitivity reached by the hard X-ray detector HXD onboard the Suzaku satellite allowed for the first time to study in detail possible new spectral features and curvature at $>10$~keV, even in a blazar quiet state, which may help understanding the properties and origin of the VHE emission.   We report here on our $\\sim 50$~ksec Suzaku observations on each of two more distant TeV-blazars known to date: 1ES~1101-232 and 1ES~1553+113. These observations were covered by quasi-simultaneous observations with the H.E.S.S. telescope system and the MAGIC telescope. None of these sources have shown in the past significant variability in the TeV-band \\cite{Aharonian06a,Aharonian06b}, likely owing to their low flux close to the detection limit of the TeV facilities.  The paper is organized as follows: We provide a brief overview of both TeV-blazars in Sect.~2. Sect.~3 is devoted to a technical description of our Suzaku data analysis and its results. In Sect.~4 we report on the quasi-simultanuous multiwavelength coverage during the Suzaku observations of 1ES 1101-232 and 1ES 1553+113. In Sect.~5 we discuss our results in particular in context to their historical behaviour. Sect.~6 summarizes the main results of this work. ",
        "conclusions": "\\subsection{Historical X-ray behaviour of 1ES 1101-232 and 1ES 1553+113}   1ES 1101-232 has been detected during the past 15 years by a range of X-ray instruments, both covering the soft (ROSAT, XMM) and extending into the hard X-ray range (BeppoSAX, RXTE, SWIFT) at various sensitivity levels. Except for the continuous monitoring by RXTE's ASM and the 19.6~ksec monitoring by XMM \\cite{Aharonian07a}, most observations were carried out in snapshot mode, with each observation encompassing typically 1-5 ksec. The continuous Suzaku observation presented here constitutes therefore the first long ($>30$~ksec) high-sensitivity X-ray data set of this source. None of the past observations showed significant short time variability ($<$hour). Although the continuous mode observations of Suzaku generally allow to probe intraday variability, we find no significant variability of this time scale during our observations. Year-by-year variations as derived from the historical X-ray light curve are less than a factor $\\sim 2$. The highest ever measured X-ray flux level has been reported for 2005 from the RXTE-PCA measurements \\cite{Aharonian07a}. There, in eleven consecutive nights the source was observed each night for $\\sim 10$~ksec. No hints for strong flux variability were found, and the AGN was categorized to be in a non-flaring state. If this is the case, 1ES 1101-232 was likely never observed in a flaring state during the past 15 years. Only 4 months later SWIFT saw the source in a smooth decline of flux level with no sub-hour variations \\cite[e.g.][]{Massa08}. Our Suzaku observations, made approximately $10$ months later, found the lowest ever measured X-ray flux from 1ES 1101-232.  So far, essentially all sensitive X-ray observations of 1ES 1101-232 prefer a curved or broken power-law representation instead a simple power-law \\cite[e.g.,][]{Perlman05}, independent of flux level and - as demonstrated in this work - extending into the hard X-ray range. If interpreted as signatures of the particle energy gain process, one may conclude that acceleration takes place at all flux levels observed so far. In all cases there were no indications for a significant amount of X-ray absorbing material in the line-of-sight beyond the Galactic absorption column density.  The average spectrum during our Suzaku observations was quite similar to the one from the 20\\% higher flux XMM observations in 2001 \\cite[$\\Gamma\\sim2.1-2.4$:][]{Perlman05}, where also no significant variability was noticed. The 1997/98 BeppoSAX observations \\cite{Wolter00} had a similar value of the break energy at a factor $\\sim 1.3-1.9$ higher flux level. The flux reported from the 2004 observations with XMM was comparable to the one measured by BeppoSAX in 1997, however, the spectral break shifted to lower energies by a factor $\\sim 1.2$. Because of the lacking soft X-ray sensitivity, the PCA coverage of 1ES 1101-232 in 2005 revealed a break energy at much higher energies ($\\sim 8$~keV). If the 1-8~keV X-ray spectral index is compared to the observed flux level at each observation, no statistically convincing indication for a flux-index correlation is found (Spearman rank correlation coefficient $R_s=-0.34$ with chance probability $P_c=0.33$, Kendall's $\\tau_K=-0.30$ with $P_c=0.23$). Comparing, e.g., the XMM measurements carried out in 2001 and 2004, no significant spectral changes within the uncertainties occured despite a factor $\\sim 1.6$ higher flux in 2004. On the other side, the BeppoSAX observations revealed a spectral softening from the higher flux state in 1997 to the lower flux state in 1998. When compared to the high flux RXTE data from 2005 a noticable softening occured despite the overall high flux level. If following the evolution of the spectrum versus flux, no clear spectral hysteresis could be found. In summary, it seems that spectral changes observed from annual visits of this source do not follow particular patterns or relations with respect to flux changes. The overall spectral changes lie within a rather narrow range between $\\sim 2$ and 2.5 for the 1-8~keV photon spectral index.   The measured so far comparably long X-ray variability time scale in 1ES 1101-232 may imply that this source was observed only in the quiescence over the last $\\sim 15$~years. If one assumes the measured activity states from one year to the other to be not causally related, the annual measurements may then be treated as an independent ensemble of measurements from an HBL. In particular, it seems reasonable to probe the behaviour in the peak luminosity-peak energy diagram within the available HBL range. For a robust determination of the peak energy, low energy data, preferably in the IR/optical/UV range, are necessary. Giommi et al. (2005) used optical data from the literature in combination with the BeppoSAX measurements, while both XMM pointings had the advantage of simultaneous optical coverage with the Optical Monitor (OM) onboard the satellite, and SWIFT the onboard UVOT instrument. We did not find any convincing anti-correlation between peak energy and peak flux within the available 6 ensemble members ($R_s=-0.6$ with $P_c=0.21$, $\\tau_K=-0.33$ with $P_c=0.35$).    1ES~1553+113 is known as a generally X-ray soft, bright HBL with only modest X-ray variability (factor $\\leq 3.5$). SWIFT in 2005 \\cite{Trama07,Massa08} achieved a detection of this source up to only 8~keV, despite its high flux level in October 2005 that was reached within half a  year from a 3.5 times lower flux level. BeppoSAX in 1998 \\cite{Donato05} saw it only with the LECS/MECS at a relatively low flux level, no detection with the PDS was possible.  With our Suzaku observation, 1ES~1553+113's spectrum could finally be measured up to 30~keV while descending from the 2005 high flux level. Like for 1ES 1101-232, past X-ray observations encompass mostly snapshot observations of typically 3-10~ksec each. RXTE's observation campaign in 2003 \\cite{Osterman06} involved visits to the source about 3 times per day for $\\sim 3$~ksec, for a total period of 21 days. During this time a very smooth rise in flux up to 3 times the lowest flux level was observed. No hints of subhour variations were recorded in any of the past observations. With our $\\sim 50$~ksec continuous Suzaku measurements we tested for intraday variability, but obtained a null result also at this intermediate flux level which was quite comparable to the 2001 flux recorded by XMM \\cite{Perlman05}.  No significant differences were found between the spectra reported from the 2001 and 2006 observations: the average spectrum during our Suzaku observations was, within the uncertainties, in good agreement to the one from the XMM observations in 2001 \\cite[$\\Gamma\\sim2.2-2.4$: ][]{Perlman05}. The surprising spectral stability of this source holds even for the low flux state recorded with RXTE in 2003 \\cite{Osterman06}, despite the factor $\\sim$~5 difference in flux. Even within the tripling of the flux during the PCA observations spectral changes were not significantly beyond their uncertainties. In particular, any systematic spectral trend with changing flux is not apparent, neither on a year-by-year time scale nor within the rising part of the RXTE monitoring data within their admittedly large uncertainties: as is the case for 1ES 1101-232, the photon index above $3$~keV typically lingered around 2 and 2.5 even during flux increases by a factor $\\sim 2$ or more. Those modest spectral changes together with the smoothness of the long-term light curve may suggest that explosive events are among the least likely scenarios to account for the observed flux variations here.  The 1998 BeppoSAX observations \\cite{Donato05} had a lower break energy ($\\sim1$~keV) at a factor $\\sim 2.5$ lower flux level as compared to the 2006 Suzaku observations. SWIFT detected 1ES 1553+113 in both, a rather low flux state (April 2005) and relatively high flux level (October 2005), but with a surprisingly stable peak energy of $\\sim 0.4-0.5$~keV in the XRT spectrum \\cite[e.g.][]{Massa08}, $\\sim 3$ times lower than the Suzaku data from 2006 imply. The flux reached during the April 2005 observations with SWIFT was comparable to the one measured by BeppoSAX in 1998, however the peak energy from the XRT data shifted to lower energies by a factor $\\sim 2.5$. No break energies were reported from the 2003 RXTE/PCA data \\cite{Osterman06}.   Following the argumentation above we probe here the dependence between peak luminosity and peak energy also for 1ES 1553+113, and utilize preferably data from the IR/optical/UV range to complement the X-ray band measurements. Among the past observations only the XMM and SWIFT pointings were quasi-simultaneous measurements with optical/UV available, namely OM, UVOT and ROTSE data, respectively. We use here our Suzaku observations with the data from the quasi-simultaneous optical coverage by the KVA optical telescope for a combined X-ray - optical spectral fit. A log-parabolic shape $dN/dE \\propto E^{[-\\Gamma-\\beta\\log(E)]}$ was used to determine the peak energy to $\\sim 0.01$~keV with parameters $\\beta=0.075\\pm 0.001$ and $\\Gamma=2.298\\pm 0.003$. Note, however, that this representation does only give a $\\chi_{\\rm red}^2 \\approx 5.7$ due to the large weight given to the only optical point (that has a small uncertainty) with respect to the many X-ray data points. More lower energy points would clarify this. We did not use the peak values given for the 2003 RXTE campaign for this undertaking to assure for a sample with equally determined peak energies: the cm radio data from Osterman et al. (2006) for determining the peak energy would give too much weight towards the longer wavelengths when compared to a procedure where only optical/UV in combination with the X-ray data are used. The peak energies from the UVOT measurements range from $\\sim 0.05$ to 0.02~keV from the high to the low flux state, respectively. The latter is in very good agreement with the values derived from the XMM 2001 observations.  Also for 1ES 1553+113 any convincing systematic correlation between peak energy and peak flux within the very limited sample of 4 ensemble members is lacking evidence ($R_s=0.39$ with $P_c=0.61$, $\\tau_K=0.55$ with $P_c=0.26$). This may indicate that particle acceleration up to the maximum energy is not preferentially limited by Compton losses if beaming does not change significantly. Indeed, because of the lack of sufficiently dense radiation fields in HBL-type AGN, synchrotron losses may dominate the radiative particle loss channel. The absence of a link between the peak luminosity and peak energy can then be understood if either an increase of the jet's synchrotron radiation field is not connected to an increase of the magnetic field strength alone on year time scales, or expansion losses and/or escape from the emission region sets a boundary to the maximum electron energy at the so far observed activity states in 1ES~1101-232 and 1ES~1553+113.  \\subsection{Spectral energy distributions}  The overall SED of 1ES~1101-232, including the Suzaku spectrum obtained by us and the simultaneous H.E.S.S. data is shown in Fig.~\\ref{fig3}. If corrected by the minimal EBL \\cite[P0.45 from][]{Aharonian06a}, comparable power is emitted at $\\gamma$-rays than at the synchrotron hump with a $\\gamma$-ray peak beyond TeVs at the time of the quasi-simultaneous Suzaku -- TeV observations. In the framework of a simple one-zone SSC model a (assumed uniform) field strength of sub-equipartition strength $B^2 \\propto (\\nu F_{\\nu,IC})/(\\nu F_{\\nu,syn})$ can then be deduced. Alternatively, further target photon fields in the vicinity of the $\\gamma$-ray emission zone can lead to an enhanced $\\gamma$-ray photon output. In hadronic blazar emission models TeV-emission is the result of either reprocessed cascade emission initiated by pairs and/or photons produced in photomeson interactions, or by proton and/or $\\pi^{\\pm}/\\mu^{\\pm}$-synchrotron radiation \\cite[e.g.][]{Muecke01,Muecke03}. In environments of weak ambient photon fields proton-photon interactions constitute a correspondingly sub-dominant contribution to the $\\gamma$-ray component, which then leaves proton synchrotron (or curvature) radiation as the main $\\gamma$-ray producer.  \\clearpage \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{f9.eps}} \\caption{Broadband SED of 1ES~1101-232. The observed H.E.S.S. data were de-absorbed using the minimal EBL \\cite[P0.45 from][]{Aharonian06a}. Present Suzaku data and the simultaneous de-absorbed July 2006 H.E.S.S. data are indicated in red. Blue points (XMM, RXTE, HESS) indicate data from 2004/05 \\cite{Aharonian07a}, while historical data collected from the literature are in black open symbols, and light grey \\cite[BeppoSAX:][]{Wolter00}. The SWIFT data \\cite[][not shown here]{Massa08} are compatible with the XMM flux level.} \\label{fig3} \\end{figure} \\clearpage  TeV observations of 1ES~1101-232 were used previously \\cite{Aharonian06a} to set the most stringent limits on the EBL at $\\sim 0.7-4\\mu$m using spectral considerations \\cite{Aharonian06a}. These limits could not be improved owing to the limited quality from the low flux TeV-data from 2006.   The overall SED of 1ES~1553+113 with the contemporaneous Suzaku, VHE (H.E.S.S., MAGIC) and optical (KVA) data is shown in Fig.~\\ref{fig10}, complemented with historical ones from the literature. Using a minimal EBL (P0.45) for de-absorption, approximately equal or more power is emitted at $\\gamma$-rays than at the synchrotron hump with a $\\gamma$-ray peak at sub-TeVs. In the framework of a simple one-zone SSC model, this ratio indicates sub-equipartition field strengths in the emission region. Furthermore, if the same electrons emit TeV- and X-ray photons co-spatially, then the intrinsic spectral index at VHEs must not be harder than at X-ray energies. For an EBL density at the galaxy counts flux level this is the case for $z \\geq 0.4$. Alternatively, external photon fields may contribute to the target photon fields for inverse Compton scattering, or hadronic emission models may be at work here.  \\clearpage \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{f10.eps}} \\caption{Broadband non-simultaneous SED of 1ES~1553+113. The quasi-simultaneous data points (KVA, Suzaku, H.E.S.S., MAGIC) data are indicated as (red) squares. The green/grey line respresents the highest flux level observed by SWIFT \\cite{Massa08} of this source. The $\\gamma$-ray data have been de-absorbed for a source redshift of $z=0.3$ using the minimal EBL \\cite[P0.45 from][]{Aharonian06a}.} \\label{fig10} \\end{figure} \\clearpage"
    },
    "0808/0808.3843_arXiv.txt": {
        "abstract": "Radio observations from decimetric to submillimetric wavelengths are now a basic tool for the investigation of comets.  Spectroscopic observations allow us i) to monitor the gas production rate of the comets, by directly observing the water molecule, or by observing secondary products (e.g., the OH radical) or minor species (e.g., HCN); ii) to investigate the chemical composition of comets; iii) to probe the physical conditions of cometary atmospheres: kinetic temperature and expansion velocity.  Continuum observations probe large-size dust particles and (for the largest objects) cometary nuclei.  Comets are classified from their orbital characteristics into two separate classes: i) nearly-isotropic, mainly long-period comets and ii) ecliptic, short-period comets, the so-called Jupiter-family comets.  These two classes apparently come from two different reservoirs, respectively the Oort cloud and the trans-Neptunian scattered disc.  Due to their different history and --- possibly --- their different origin, they may have different chemical and physical properties that are worth being investigated.  The present article reviews the contribution of radio observations to our knowledge of the Jupiter-family comets (JFCs).  The difficulty of such a study is the commonly low gas and dust productions of these comets.  Long-period, nearly-isotropic comets from the Oort cloud are better known from Earth-based observations.  On the other hand, Jupiter-family comets are more easily accessed by space missions.  However, unique opportunities to observe Jupiter-family comets are offered when these objects come by chance close to the Earth (like 73P/Schwassmann-Wachmann 3 in 2006), or when they exhibit unexpected outbursts (as did 17P/Holmes in 2007).  About a dozen JFCs were successfully observed by radio techniques up to now. Four to ten molecules were detected in five of them.  No obvious evidence for different properties between JFCs and other families of comets is found, as far as radio observations are concerned. ",
        "introduction": "\\label{sect-intro}  Jupiter-family comets (JFCs, aka \\textit{ecliptic comets}) are short-period, low-inclination comets likely to undergo orbital perturbations by Jupiter.  We adopt here the definition based on the Tisserand invariant with respect to Jupiter $T_J$: JFCs are comets for which $2 < T_J < 3$ \\citep{levi:1996}. Special cases are 2P/Encke which slightly exceeds the $T_J = 3$ limit and for this reason sometimes classified as a \\textit{Encke-type} comet, and 29P/Schwassmann-Wachmann 1 with $T_J = 2.99$, alternatively classified as a \\textit{Centaur}.  JFCs are believed to come from the trans-Neptunian scattered disc.  They contrast with the \\textit{nearly isotropic comets}, which comprise long-period comets as well as short-period comets (the so-called \\textit{Halley type comets}), supposed to come from the Oort cloud.  Comets certainly did not form in these trans-Neptunian reservoirs.  Their real sites of formation and their orbital evolution are still highly debated topics.  (For reviews on cometary families and their dynamical evolution, see \\citet{levi:1996, fern:2008, lowr+:2008, morb:2008, morb+:2008, mars:2008}).  Having different dynamical histories and --- possibly --- different origins, these distinct classes of comets may have different chemical and physical properties that are worth being investigated.   \\begin{table*} \\caption{Remote sensing observing conditions for a selection of comets} \\label{table-remote} \\begin{center} \\begin{tabular*}{\\hsize}[]{@{\\extracolsep{\\fill}}llcccc} \\hline comet & date & $r_h$ & $\\Delta$ & $Q_\\mathrm{H_2O}$ & $FM$ \\\\ &      & [AU]  & [AU]     & [$10^{28}$ s$^{-1}$] \\\\ \\hline \\multicolumn{5}{l}{\\textit{Nearly-isotropic comets}} \\\\ C/1986 P1 (Wilson)           & May 1987       & 1.3 & 1.0  &  12 &  12 \\\\ C/1989 X1 (Austin)           & April 1990     & 1.2 & 0.25 & 2.5 &  10 \\\\ C/1990 K1 (Levy)             & August 1990    & 1.3 & 0.45 & 25  &  55 \\\\ C/1996 B2 (Hyakutake)        & March 1996     & 1.1 & 0.11 & 25  & 225 \\\\ C/1995 O1 (Hale-Bopp)        & April 1997     & 0.9 & 1.4  &1000 & 700 \\\\ C/1999 S4 (LINEAR)           & July 2000      & 0.77& 0.38 & 10  &  25 \\\\ C/1999 T1 (McNaught-Hartley) & December 2000  & 1.2 & 1.6  & 10  &   6 \\\\ C/2001 A2 (LINEAR)           & June 2001      & 1.0 & 0.24 & 10  &  40 \\\\ C/2000 WM$_1$ (LINEAR)       & December 2001  & 1.2 & 0.32 &  4  &  12 \\\\ C/2001 Q4 (NEAT)             & May 2004       & 1.0 & 0.32 & 20  &  60 \\\\ C/2002 T7 (LINEAR)           & May 2004       & 0.8 & 0.27 & 10  &  25 \\\\ C/2003 K4 (LINEAR)           & December 2004  & 1.4 & 1.2  & 15  &  12 \\\\ C/2004 Q2 (Machholz)         & January 2005   & 1.2 & 0.35 & 25  &  70 \\\\ \\hline \\multicolumn{5}{l}{\\textit{Halley-type comets}} \\\\ 1P/Halley                    & January 1986   & 0.7 & 1.5  & 120 &  80 \\\\ 109P/Swift-Tuttle            & November 1992  & 1.0 & 1.3  & 50  &  40 \\\\ 153P/2002 C1 (Ikeya-Zhang)   & April 2002   & 1.0 & 0.40 & 25  &  60 \\\\ 8P/Tuttle                    & January 2008   & 1.1  & 0.25 & 3   & 12 \\\\ \\hline \\multicolumn{5}{l}{\\textit{Jupiter-family comets}} \\\\ 22P/Kopff                    & April 1996     & 1.7  & 0.9  & 3.5 & 4 \\\\ 21P/Giacobini-Zinner         & October 1998   & 1.2  & 1.0  & 3   & 3 \\\\ 19P/Borrelly                 & September 2001 & 1.36 & 1.47 & 3   & 2 \\\\ 2P/Encke                     & November 2003  & 1.0  & 0.25 & 0.5 & 2 \\\\ 9P/Tempel~1                  & July 2005      & 1.5  & 0.77 & 1   & 1.3 \\\\ 73P/Schwassmann-Wachmann~3   & May 2006       & 1.0  & 0.08 & 2   & 25  \\\\ 17P/Holmes                   & October 2007   & 2.4  & 1.6  & $>100$ & $>60$  \\\\ 67P/Churyumov-Gerasimenko    & Mars 2009      & 1.24 & 1.7  & 1   & 0.6 \\\\ 103P/Hartley~2               & October 2010   & 1.1  & 0.12 & 1.2 & 10 \\\\ 45P/Honda-Mrkos-Pajdu\\v{s}\\'akov\\'a & August 2011    & 1.0  & 0.06 & 0.5 &  8 \\\\ \\hline \\end{tabular*} \\end{center}  Date is for best observing conditions. \\\\ $r_h$ = distance to Sun and $\\Delta$ = distance to Earth at that date. \\\\ $Q_\\mathrm{H_2O}$ = water production rate at that date.  \\\\ $FM = Q_\\mathrm{H_2O}/\\Delta$ = \\textit{figure of merit}, roughly proportional to the signal of cometary molecules. \\end{table*}  JFCs and other comets are far from being equally well observed.  To show this for Earth-based observations, we will use the \\textit{figure of merit} parameter $FM = Q_\\mathrm{H_2O}/\\Delta$, where $Q_\\mathrm{H_2O}$ is the water production rate in units of $10^{28}$ s$^{-1}$ and $\\Delta$ is the distance to the observer in AU. (Note that this parameter differs from the \\textit{figure of merit} introduced by \\citet{mumm+:2002}, which includes a dependency on the distance to the Sun.)  This parameter is roughly proportional to the expected signal intensity (for comets at heliocentric distances of the order of 1~AU), and allows us to evaluate and compare the observability of comets. Table~\\ref{table-remote} gives the figures of merit for Earth-based observations of recent long-period comets, as well as for recent and future returns of short-period comets.  One can see that unexpected, long-period comets from the nearly isotropic class of comets and Halley-type comets offer much better opportunities than short-period comets.  Indeed, the two best comets in the last twenty years were C/1996 B2 (Hyakutake) and C/1995 O1 (Hale-Bopp); unprecedented spectroscopic observations leaded to the identification of many molecules for the first time in these two objects \\citep{bock+:2005}.  All JFCs are comets with low water production rates $Q_\\mathrm{H_2O}$ of at most a few 10$^{28}$~s$^{-1}$ at perihelion.  The best observing conditions occur for these comets which make a close approach to the Earth (i.e., for $\\Delta$ significantly smaller than 1~AU).  This was recently the case for 73P/Schwassmann-Wachmann 3 (minimum geocentric distance $\\Delta = 0.08$ AU in May 2006) and it will also happen for 103P/Hartley~2 in the near future ($\\Delta = 0.12$ AU in October 2010).  But the \\textit{figure of merit} of such JFCs is still lower than that of some long-period comets for which $FM$ could exceed 50. Other unique opportunities also occur when JFCs exhibit unexpected outbursts during which the gas production rate may be increased by several orders of magnitude.  This was recently the case for 17P/Holmes, but such events are rare.  On the other hand, short-period comets, with their predicted returns, are the only practicable targets for space missions, which are not yet versatile enough to accommodate unexpected comets.  Flybys at a low velocity and rendezvous are only possible for ecliptic comets, due to the energy limitations of current space technology.  1P/Halley is the only explored comet which does not belong to the Jupiter family: the price to pay was a very high flyby velocity ($\\approx 70$ km~s$^{-1}$).  Spectroscopy at radio wavelengths is now a basic tool for the investigation of comets.  It allows us:  \\begin{itemize}  \\item to monitor the gas production rate of the comets, by directly observing the water molecule, or by observing secondary products (e.g., the OH radical) or minor species (e.g., HCN);  \\item to investigate the chemical composition of comets;  \\item to probe the temperature of cometary atmospheres by observing simultaneously several rotational lines of the same molecule;  \\item to investigate the expansion velocity and kinematics of cometary atmospheres from the observation of line shapes.  \\end{itemize}  Continuum observations of comets in the radio spectral range are also well suited for studying the nucleus and the dust coma via their thermal radiation. Since millimetric and submillimetric radiation can only be efficiently radiated from large particles ($\\gtrsim 1$ mm) which comprise most of the dust mass, this technique is a useful probe of the total dust mass production and of the size distribution of the large grains.  The present article reviews the contribution of radio observations to our knowledge of JFCs.  The outcome of radar experiments \\citep{harm+:2005}, which probe the nucleus and large dust particles, is beyond the scope of this paper. ",
        "conclusions": "Radio observations of JFCs (see Table~\\ref{table-summary} for a summary) are still sparse.  Up to now, successful radio observations of JFCs have been limited to a dozen objects, mainly restricted to exceptional comets (comets which made a close approach to Earth and/or which showed an outburst), or to the targets of space missions for which a special effort was made.  New radio instruments are to be soon available for the observations of comets with an increased sensitivity:  \\begin{itemize}  \\item the Large Millimeter Telescope (LMT), under completion in Mexico \\citep{irvi-schl:2005}; with its 50-m antenna, it will be soon the largest millimetric telescope of its category;  \\item the Herschel Space Observatory, to be launched in 2009 \\citep{crov:2005};  \\item the Atacama Large Millimetre Array (ALMA) in Chili \\citep{bive:2005, bock:2008};  \\item the Square Kilometre Array (SKA), which will operate at decimetric-centimetric wavelengths \\citep{butl+:2004};  \\item the MIRO instrument on Rosetta, a 30-cm diameter radio telescope to observe dedicated molecular radio lines in situ in comet 67P/Churyumov-Gerasimenko, as well as the continuum emission of its nucleus \\citep{gulk+:2007}.  \\end{itemize}   Specific opportunities to observe JFCs with a high $FM$ will occur in the near future (Table~\\ref{table-remote}):  \\begin{itemize}  \\item 103P/Hartley~2 in October 2010 (with a flyby of the redirected Deep Impact spacecraft, renamed as the EPOXI mission);  \\item 45P/Honda-Mrkos-Pajdu\\v{s}\\'akov\\'a in August 2011.  \\end{itemize}  It will thus be possible to significantly increase the sample of JFCs with known molecular composition, and to pursue the comparative study of the different cometary families.  Isotopic measurements in cometary volatiles are also important diagnostics on the origin of cometary material.  Most measurements have been made using millimetre or submillimetre spectroscopy.  However, if one excepts the Jupiter-family comet 17P/Holmes (Section~\\ref{17P}), the capabilities of ground and space-based instrumentation have limited the investigations to a few bright long-period comets, as reviewed in \\citet{altw-bock:2003} and \\citet{bock+:2005} and reported in \\citet{bive+:2007-PASS}.  Isotopic ratios could be different in JFCs and Oort cloud comets if the two populations formed at different places or different times in the solar nebula.  Though the LMT promises to be very useful for the study of the molecular composition of JFCs, isotopic investigations with this telescope will be limited to exceptionally bright comets \\citep{irvi-schl:2002}.  The HIFI instrument of the Herschel Space Observatory is to provide the opportunity to detect the 1$_{10}$--1$_{01}$ line of HDO at 509.3~GHz in comet 103P/Hartley 2, and to observe simultaneously several H$_2$O lines for an accurate determination of the D/H ratio in water. ALMA will allow the detection of HDO at 465 and 894 GHz in short-period comets with $FM > 5$, and DCN in comets with $FM > 10$.   \\balance  \\footnotesize{"
    },
    "0808/0808.1594_arXiv.txt": {
        "abstract": "Using the recently upgraded Long Baseline Array, we have measured the trigonometric parallax of PSR J0437--4715 to better than 1\\% precision, the most precise pulsar distance determination made to date.  Comparing this VLBI distance measurement to the kinematic distance obtained from pulsar timing, which is calculated from the pulsar's proper motion and apparent rate of change of orbital period, gives a precise limit on the unmodeled relative acceleration between the Solar System and PSR J0437--4715, which can be used in a variety of applications. Firstly, it shows that Newton's gravitational constant $G$ is stable with time ($\\Gdot/G  = (-5 \\pm 26) \\times 10^{-13}$\\ yr$^{-1}$, 95\\% confidence). Secondly, if a stochastic gravitational wave background existed at the currently quoted limit, this null result would fail $\\sim\\!50$\\% of the time. Thirdly, it excludes Jupiter-mass planets within 226 AU of the Sun in 50\\% of the sky (95\\% confidence). Finally, the $\\sim\\!1$\\% agreement of the parallax and orbital period derivative distances provides a fundamental confirmation of the parallax distance method upon which all astronomical distances are based. ",
        "introduction": "PSR J0437--4715 has been observed intensively since its discovery by \\citet{johnston93a}. It is the brightest and nearest observed millisecond pulsar, and has also been studied in the optical \\citep{bell93a}, ultraviolet \\citep{kargaltsev04a}, and X--ray \\citep{zavlin02a} bands.  The high rotational stability and close proximity of this pulsar--white dwarf binary system make it an excellent probe of General Relativity (GR) and alternative forms of gravitational theories.  The measurement of its Shapiro delay by \\citet{van-straten01a}, where the radio waves from the pulsar are delayed as they pass through the gravitational potential of the companion, is one such test which has shown consistency with GR predictions.  The search for the low frequency stochastic gravitational wave background (GWB) using pulsar timing arrays \\citep[e.g.][]{jenet05a} is another test of GR which is facilitated in part by timing of J0437--4715.  Although variation of Newton's gravitational constant $G$\\ with time is forbidden in GR, this is not required in alternate formalisms of gravity \\cite[e.g.][]{brans61a}. \\citet{verbiest08a} have measured the rate of change of orbital period \\pbdot\\ in the J0437--4715 system to better than 2\\% precision and shown that the agreement of the derived kinematic distance with the parallax distance derived from timing limits the time variation of $G$\\ to less than 3 parts in $10^{11}$. In this Letter, we present a new Very Long Baseline Interferometry (VLBI) determination of the position, proper motion, and parallax of PSR J0437--4715 and show that this improves the previously published \\Gdot\\ limit derived from this system by an order of magnitude, approaching the best published limit ($\\Gdot/G  = (4\\pm9)\\times 10^{-13}$\\ yr$^{-1}$), which is derived from Lunar Laser Ranging \\citep[LLR;][]{williams04a}.  Our VLBI observations and results are presented in \\S\\ref{sec:vlbi} and \\S\\ref{sec:results} respectively and the limitations on apparent accelerations (due to \\Gdot\\ or other possible causes such as unseen massive planets) are derived in \\S\\ref{sec:anomacc}. In \\S\\ref{sec:gwb}, we investigate the impact of the stochastic GWB on our results and derive an independent limit on the GWB amplitude.  We summarize our conclusions in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have observed PSR J0437--4715 using the LBA and obtained the most precise pulsar distance determination to date, with an error $<1.5$\\,pc.  Combined with accurate timing data, this has enabled us to confirm that  $|\\Gdot/G| < 3.1 \\times 10^{-12}$\\ yr$^{-1}$, the most stringent limit not obtained though Solar System tests. Alternatively, assuming an unchanging gravitational force, the results can be interpreted as excluding any unseen Jupiter--mass planets within 226 AU of the Sun in 50\\% of the sky. Finally, the agreement between VLBI and timing results in this single case implies an upper limit to the stochastic GWB amplitude which is within an order of magnitude of the best limit derived from observations of ensembles of pulsars."
    },
    "0808/0808.3628_arXiv.txt": {
        "abstract": "We investigate the direct contribution of strong, sunspot-like magnetic fields to helioseismic wave travel-time shifts via two numerical forward models, a 3D ideal MHD solver and MHD ray theory. The simulated data cubes are analyzed using the traditional time-distance center-to-annulus measurement technique. We also isolate and analyze the direct contribution from purely thermal perturbations to the observed travel-time shifts, confirming some existing ideas and bring forth new ones: (i) that the observed travel-time shifts in the vicinity of sunspots are largely governed by MHD physics, (ii) the travel-time shifts are sensitively dependent on frequency and phase-speed filter parameters and the background power below the $p_1$ ridge, and finally, (iii) despite its seeming limitations, ray theory succeeds in capturing the essence of the travel-time variations as derived from the MHD simulations. ",
        "introduction": "Local helioseismic diagnostic methods such as time-distance helioseismology \\citep{duvalletal1993}, helioseismic holography \\citep{lb1997} and ring-diagram analysis \\citep{hill1988}, have over the years provided us with unprecedented views of the structures and flows under sunspots and active regions. However, a growing body of evidence appears to suggest that interpretations of the measured statistical changes in the properties of the wave-field may be rendered inaccurate by complexities associated with the observations and wave propagation physics. Incorporating the full MHD physics and understanding the contributions of phase and frequency filters, and differences in the line formation height, are thought to be central to future models of sunspots.  One of the earliest studies that highlighted the interaction of waves with sunspots was the Fourier-Hankel analysis of \\cite{bdl1987}, who found that sunspots can absorb up to half of the incident acoustic-wave power and shift the phases of interacting waves quite significantly (see also \\citealt{braun1995}). These results were echoed over the years by a steady steam of theoretical results (e.g. \\citealp{bogdanetal1996}; \\citealp{cb1997}; \\citealp{ccb2003}; \\citealp{crouchetal2005}; \\citealp{cally2007}) that have consistently emphasized the need for more sophisticated modeling and interpretation of wave propagation in strongly magnetized regions.  Important advances in our observational understanding of sunspots were also achieved by \\cite{duvalletal1996} and \\cite{zkd2001}, who inferred the presence of flows underneath sunspots, and \\cite{kds2000} who estimated the sub-surface wave-speed topology. However, while the inversion procedures applied to derive these results fail to directly account for the tensorial nature of magnetic field effects, the action of the field is mimicked via changes in the acoustic properties of the medium (the so-called wave speed). Recently however, numerical forward models of helioseismic wave (e.g. \\citealp{cgd2008}; \\citealp{hanasoge2008}) and ray \\citep{mc2008} propagation in magnetized atmospheres have been developed and are beginning to make inroads into this problem. In particular, the results of \\cite{mc2008} and Cameron (2008; private communication) strongly suggest that active-region magnetic fields play a substantial role in influencing the wave field, and that the complex interaction of magnetic fields with solar oscillations, as opposed changes in the wave-speed, are the major causes of observed travel-time inhomogeneities in sunspots.  We study the impact of strong magnetic fields on wave propagation and the consequences for time-distance helioseismology using two numerical forward models, a 3D ideal MHD solver and MHD ray theory. The simulated data cubes are analyzed using the traditional surface-focused center-to-annulus method frequently applied in the time-distance analyses of sunspots (e.g., \\citealp{cbk2006}). Furthermore, we apply the same method as \\cite{mc2008} to also isolate and analyze the thermal contribution to the observed travel-time shifts. ",
        "conclusions": "Incorporating the full MHD physics into the various forward models used in local helioseismology is essential for testing inferences made in regions of strong magnetic fields. By comparing numerical simulations of MHD wave-field and ray propagation in a model sunspot, we find that: i) the observed travel-time shifts in the vicinity of sunspots are strongly determined by MHD physics, although sub-surface thermal variations also appear to affect ray timings by modifying the acoustic cut-off frequency, ii) the time-distance travel-time shifts are strongly dependent on frequency, phase speed filter parameters and the background power below the $p_1$ ridge, and finally iii) MHD ray theory succeeds in capturing the essence of center-to-annulus travel-time variations as derived from the MHD simulations.  The most unsettling aspect about this analysis is that despite using a background stratification that differs substantially from Model S \\nocite{cdetal1996} (Christensen-Dalsgaard et al. 1996) and a flux tube that clearly lacks a penumbra, the time shifts still look remarkably similar (at least qualitatively) to observational time-distance analyses of sunspots. Preliminary tests conducted with different sunspot models (e.g.,  different field configurations, peak field strengths etc.) have also provided similar results. Given the self-consistency of these results, as derived from both forward models, it could imply that we are pushing current techniques of local helioseismology to their very limits. It would appear that accurate inferences of the internal constitution of sunspots await a clever combination of forward modeling, observations, and a further development of techniques of statistical wave-field analysis."
    },
    "0808/0808.3905_arXiv.txt": {
        "abstract": "With a self-similar magnetohydrodynamic (MHD) model of an exploding progenitor star and an outgoing rebound shock and with the thermal bremsstrahlung as the major radiation mechanism in X-ray bands, we reproduce the early X-ray light curve observed for the recent event of XRO 080109/SN 2008D association. The X-ray light curve consists of a fast rise, as the shock travels into the ``visible layer\" in the stellar envelope, and a subsequent power-law decay, as the plasma cools in a self-similar evolution. The observed spectral softening is naturally expected in our rebound MHD shock scenario. We propose to attribute the ``non-thermal spectrum\" observed to be a superposition of different thermal spectra produced at different layers of the stellar envelope. ",
        "introduction": "SN 2008D, the best type Ibc supernova detected so far, is preceded by a X-Ray Outburst (XRO) captured by SWIFT satellite on 2008 January 9, and this XRO is interpreted as a shock breakout of a Wolf-Rayet (WR) progenitor with a radius of $\\sim 10^{11}$ cm \\citep {Soderberg2008}. The isotropic X-ray energy is estimated to be $\\sim 2\\times10^{46}$ erg, and there seems no collimation detected so the event is not regarded as a GRB. This XRO showed a rapid rise, peaked at $\\sim 63$ s, and a decay modelled to be exponential with an e-folding time of $\\sim 129$ s \\cite{Soderberg2008}. The follow-up optical and ultraviolet observations indicate a total supernova kinetic energy of $\\sim 2-4\\times10^{51}$ erg and a mass of SN ejecta to be $\\sim 3-5$ M$_{\\odot}$ \\cite{Soderberg2008}. Some authors estimate from a detailed spectral analysis that SN 2008D, originally a $\\sim 30$ M$_{\\odot}$ star, has a spherical symmetric explosion energy of $\\sim 6\\times10^{51}$ erg and an ejected mass $\\sim 7$ M$_{\\odot}$ \\cite{Mazzali}. The evolution of optical spectra of XRO-SN 2008D resembles that of XRO-SN 2006aj, whose progenitor is also believed to be a WR star \\cite{Campana2006}.  The production of $\\gamma$-rays and X-rays by shock breakouts has been proposed earlier \\citep{Colgate, Chevalier}. This XRO and the associated SN present an unprecedented case to be investigated in details, especially on interpretations for the rise and decay times of the X-ray light curve. The claim of an exponential decay may be premature given a fairly large scatter, and it may have concealed valuable physical clues offered by this XRO. During the XRO, the observed spectroscopic softening still lacks a convincing explanation. Here, we advance a self-similar MHD rebound shock model in an attempt to reproduce the observed X-ray light curve. The next section contains an overall description of the self-similar MHD model and the procedure of analysis; in the third section, we compare our model results with data; and conclusions are summed up in the last section. ",
        "conclusions": "With a self-similar MHD void shock model and the thermal bremsstrahlung as the main radiation loss, we obtain a fairly good fit to the X-ray light curve observed and confirm that XRO 080109 is most likely a shock breakout event. We identify that the decay in the X-ray light curve follows a power law (instead of an exponential law) in time since the core collapse and rebounce, which occurred $\\sim 552$ s before the observation of X-ray emissions. Meanwhile, the spectral softening is expected qualitatively. In this work, we use the most simplified radiation transfer presumption that the radiation produced in the `visible layer' can be totally observed. Actually, the optical depth varies with radius and should be treated in a more elaborate manner. Additionally, we presume that the boundaries of the ``visible layer\" $r_{in}$ and $r_{out}$ do not vary with time. Despite all these idealizations, our self-similar dynamic approach appears to be suitable to couple with the radiation process and to model X-ray outbursts in supernova as observed.  Regarding the X-ray spectra observed, they cannot be fitted with a simple blackbody profile, and a nonthermal power law profile was suggested \\cite{Soderberg2008}. Several mechanisms have been proposed to explain such X-ray spectra, for example the bulk comptonization by scatterings of the photons between the ejecta and a dense circumstellar medium \\cite{Soderberg2008}, or diluted thermal spectra which require the thermalization occurs at a considerable depth in the supernova \\cite{Chevalier2008}. We propose that the power-law profile might be a natural result of multi-colour superposition of blackbody spectra. Based on our scenario outlined here, X-ray emissions come from different layers within the radiative layer around the stellar surface, and the radiative layer has different temperatures at different depth. As a result, the observed X-ray spectra are the superposition of thermal blackbody components with different temperatures. We suggest that this might resolve issues of spectral profile and evolution.  During breakouts of rebound shocks and in the presence of MHD shock accelerated relativistic electrons usually presumed with a power-law energy spectrum within a certain electron energy range, we could also compute synchrotron emissions associated with such kind of SN shock breakouts. Among others, it is then possible to follow the evolution of magnetic field strength associated with SN explosions \\cite{Lou94, LouWang07} and estimate the effectiveness of accelerating relativistic particles (i.e., high-energy cosmic rays \\cite{Science06}). There is the freedom of choosing a few parameters to fit the data at a certain epoch. It is then possible to test the hypothesis of a self-similar shock evolution by further observations. In a more general perspective and on the basis of our dynamic models for rebound MHD shocks, we hope to further develop radiative diagnostics for shock breakouts of supernovae and thus for SN related GRBs.   \\begin{theacknowledgments} This research has been partially supported by Tsinghua Center for Astrophysics (THCA), by the National Natural Science Foundation of China (NSFC) grants 10373009 and 10533020 and by the National Basic Science Talent Training Foundation (NSFC J0630317) at the Tsinghua University, and by the SRFDP 20050003088 and the Yangtze Endowment from the Ministry of Education at Tsinghua University. \\end{theacknowledgments}"
    },
    "0808/0808.0137_arXiv.txt": {
        "abstract": "The concept of pyramid wavefront sensors (PWFS) has been around about a decade by now. However there is still a great lack of characterizing measurements that allow the best operation of such a system under real life conditions at an astronomical telescope.\\newline In this article we, therefore, investigate the behavior and robustness of the pyramid infrared wavefront sensor PYRAMIR mounted at the 3.5 m telescope at the Calar Alto Observatory under the influence of different error sources both intrinsic to the sensor, and arising in the preceding optical system. The intrinsic errors include diffraction effects on the pyramid edges and detector read out noise. \\newline The external imperfections consist of a Gaussian profile in the intensity distribution in the pupil plane during calibration, the effect of an optically resolved reference source, and noncommon-path aberrations. We investigated the effect of three differently sized reference sources on the calibration of the PWFS.  For the noncommon-path aberrations the quality of the response of the system is quantified in terms of modal cross talk and aliasing.  We investigate the special behavior of the system regarding tip-tilt control. \\newline From our measurements we derive the method to optimize the calibration procedure and the setup of a PWFS adaptive optics (AO) system. We also calculate the total wavefront error arising from aliasing, modal cross talk, measurement error, and fitting error in order to optimize the number of calibrated modes for on-sky operations. These measurements result in a prediction of on-sky performance for various conditions. ",
        "introduction": "\\label{sec:Intr} Within the context of AO systems pyramid wavefront sensors (PWFS) are a relatively novel concept. The special interest in this concept arises from the prediction of a gain in sensitivity -- and thus in limiting magnitude -- for a nonmodulated PWFS over a Shack-Hartmann sensor (SHS) \\cite{rag99} in closed-loop conditions for a well-corrected point source. The definition of this gain is that in order to achieve the same correction quality the PWFS needs less signal than the SHS.  This gain in sensitivity results basically from the fact that the accuracy of the measurement and the resulting reconstruction error $\\sigma^2_{SH}$ of the SHS depends on $\\lambda\\over{d_{sub}}$ with $\\lambda$ = sensing wavelength, $d_{sub}$ = subaperture size typically chosen on the order of the Fried-parameter $r_0$ for typical site seeing at science wave-length. In the case of the PWFS the measurement error $\\sigma^2_{P}$ depends on $\\lambda\\over{D}$ with $D$ being the telescope diameter ($D >> r_0$). From these formulas the difference in stellar magnitudes that are needed to achieve the same Strehl ratio for the exemplary case of tip-tilt only can be derived for both sensors as $\\Delta m = -2.5 log\\left({\\sigma_P^2\\over{\\sigma_{SH}^2}}\\right)\\approx -2.5 log\\left({r_0^2\\over{D^2}}\\right)$. Simulations show , for instance, \\cite{rag99} that the gain in sensitivity for a 4\\,m class telescope and a seeing of $0.5''$ ($r_0=20$\\,cm) is predicted to be 2 magnitudes. \\subsection{The Pyramid Principle}\\label{subsec:PP} Wavefront sensing based on the pyramid principle has its origin in the Foucault knife-edge test. The historical development of the PWFS is described in \\cite{Camp2006}. \\begin{figure}[h!] \\centering \\includegraphics[width=12cm]{f1.jpg} \\caption{(Courtesy of S. Egner) The pyramid principle. An example of corresponding pixels is marked. The gradient in this position of the pupil is calculated from the intensity differences in between these pixels.} \\label{im:f1} \\end{figure} The optical setup of a pyramid sensor is shown in Figure \\ref{im:f1}.  The transmissive, four-sided pyramid prism is placed in the focal plane.  The focus is placed on the tip of the pyramid. After the pyramid, a relay lens images the pupils onto the detector.\\newline The signal a four-sided AO PWFS system uses are the intensities inside four pupil images. The illumination of these images depends on the aberrations of the wavefront. The signal $S$ one extracts is the difference in intensities $I_{1,2,3,4}$ between corresponding pixels in the four pupils ,i.e., the pixels at the same optical position in the pupils as shown in Fig.\\ref{im:f1}:  \\begin{equation}S_x(x,y)={{I_1(x,y)+I_3(x,y)-[I_2(x,y)+I_4(x,y)]}\\over{I_1(x,y)+I_2(x,y)+I_3(x,y)+I_4(x,y)}}\\end{equation} \\begin{equation}S_y(x,y)={{I_1(x,y)+I_2(x,y)-[I_3(x,y)+I_4(x,y)]}\\over{I_1(x,y)+I_2(x,y)+I_3(x,y)+I_4(x,y)}}\\end{equation} for the x- and y-direction, respectively. These signals $S$ are what we will refer to as gradients, even if they might not exactly represent wavefront gradients. In the limit of small perturbations and a telescope with infinite aperture the frequency spectrum of the signal $S_x$ of our sensor is given by \\begin{equation}\\widetilde{S_x}=i{\\rm sgn}(u)\\widetilde{\\phi}(u,v).\\end{equation} Here $\\widetilde{()}$ means Fourier transform, $\\phi(u,v)$ the phase of the electromagnetic wave with $u$ and $v$ the coordinates in Fourier space, and sgn the sign-function (see \\cite{Costa2003}). Thus the sensor is working as a phase sensor in this regime. \\newline The principle of the PWFS implies some limitations that either do not occur in other sensor types, or have a much more \"dramatic\" impact on PWFSs than on other sensor types.  In this class fall the structure of the pyramid edges that cause both diffraction and scattering, the read out noise of the system, the goodness of centering the beam on the pyramid tip, noncommon-path aberrations, and the homogeneity of the pupil illumination, the latter being important especially during calibration.  It is these limitations, that may ultimately influence the choice of this or another sensor type, that will be examined in the course of this article. \\newline The structure in the remaining part of the article is as follows: In the next section, the PYRAMIR system is described in detail. The optical path and the possible detector read out modes are explained. Section 3 accounts for the calibration procedure of the system. The peculiarities of tip-tilt calibration and flattening the wavefront are presented.  Section 4 shows fundamental limitations to the performance of a pyramid system. ``Fundamental'' in this case is not meant to be based on natural laws and constants, but on always-present aberrations and nonideal conditions. Thus, the fundamental limitations discussed here can be eased by careful alignment and set up of the system, but they can never entirely be removed. In this context, we explore the effect of static aberrations, different calibration light sources, and diffraction and scattering on the pyramid edges to the response of the system. In the next section the implications of the modal cross talk of the system aliasing and measurement error to the number of modes to be calibrated, and the residual wavefront error are calculated. We end the section with a prediction of the on-sky performance under different seeing conditions. The last section concludes the results of our measurements and the implications for any pyramid wavefront sensor.\\newline On-sky results will be presented in a following paper \\cite{Peter2008}. ",
        "conclusions": "\\label{sec:conclusion} In this article we presented laboratory measurements performed with the pyramid wavefront sensor PYRAMIR.  After a short introduction into the history of PWFS and their working principles, we presented the PYRAMIR system as infrared PWFS and explained the details of the system especially the possible read out modes of the detector. This can reach a speed of about 300 Hz with a RON of 20 $e^-$ rms.  The calibration procedure of the system was described in detail. The pitfalls of TT-calibration were discussed in detail leading to the conclusion that a static part in TT will reduce the limiting magnitude and enhance TT-jitter and cannot be tolerated. Therefore we included a static TT-part in the bias pattern of the deformable mirror (dmBias) to perfectly center the beam. Also it has to be guaranteed that the amplitudes of calibration are the same in both axis in order to gain similar performance in both directions.\\newline Some of the fundamental sources of reduced performance were examined. We found that the amount of light diffracted out of the pupil images is 50\\% for a flat wavefront decreasing nonlinearly up to 2 $rad$ wavefront error where it becomes almost stable at 20\\%. A comparison with simulations shows that in our case this loss results from diffraction at the pyramid edges and imperfections in the optics. The latter outweighing for aberrations stronger than about 1.5 $rad$.\\newline The effects of a Gaussian illumination of the pupils, as well as the effect of an extended calibration light source, were investigated.  Both have strong influence on the sensitivity of the system and should be avoided.  On the other hand an extended target during the measurement does only slightly effect the performance. Thus the calibration light source should be as point like as possible for all applications.\\newline The effect of RON on the limiting magnitude of the guide star has been widely discussed. In the case of infrared detectors this noise is quite high. In our case 20 $e^-$ rms per pixel. This reduces the limiting magnitude by 5 mag.  We investigated into the best possible mode set between the eigen modes of PYRAMIR, of the DM and the KL modes. The last two mode sets seem promising for the use on-sky. The latter has a larger linear regime than the eigen modes of the DM. The difference is small and, therefore, the correction error will not vary by much. Still the larger linear regime might help to close the loop under bad seeing conditions. \\newline We tested the best treatment of noncommon-path aberrations by applying artificial modes to the DM. The best treatment turned out to direct these aberrations completely into the path of the sensor. A possible better way using two calibrations, one for modes with positive amplitude one for those with negative amplitude was proposed to reduce the error of the mode with static aberrations, but put aside due to the fact that it was not running stably during the testing on-sky.\\newline The importance of a small calibration amplitude to minimize the resulting reconstruction error was shown. \\newline After the measurement of these fundamental properties, the subject of the best number of modes to calibrate was addressed. To solve this we investigated into the behavior of modal cross talk, aliasing, and measurement error dependence on the number of modes calibrated. We found that the averaged aliasing coefficient varies only slightly with the number of modes whereas the average modal cross talk coefficient decreases inversely linear with this number and the measurement error rises linearly with this number. Including the (theoretical) contribution of the KL modes in the wavefront error on sky we could show that the contribution of aliasing rises like $N^{(2-\\sqrt{3})\\over 2}$; the contribution of cross talk becomes constant for larger $N$ and the measurement error rises linearly with $N$. In CL only the error due to modal cross talk changes with respect to the open loop measurement. It will decrease because the modes that are corrected will contribute less to the modal cross talk than in open loop. The error of the residual wavefront decreases like $N^{-{\\sqrt{3}\\over 2}}$. Altogether we could show that at the border of the limiting magnitude the fitting error surpasses all other errors for a low number of corrected modes but will be overpowered by the measurement error at about the place of the optimum number of modes to be corrected. For the PYRAMIR system this is 5 modes.  The other errors will become important for even slightly brighter stars. Here cross talk and aliasing error are almost identical in strength.  In the case of noncommon-path aberrations in the system the modal cross talk and aliasing errors will rise. Modal cross talk increases linearly, aliasing error stays constant until the static aberration reaches the border of the linear regime of the sensor, then it increases nonlinearly. \\newline From the entire error-budged we could predict the performance on-sky for various seeing conditions. For a seeing of 1'' and a wind speed on ground level of 5$ms^{-1}$ we can achieve a 39\\% Strehl ratio in $K^\\prime$-band, but we have to run at about maximum frame rate (300 Hz). The limiting S/N per subaperture on the detector is about 0.25 or 6.0 mag in $K^\\prime$.  For good seeing conditions and a moderate loop band width, the predicted Strehl ratio will be about 60\\% for bright stars. Again the limiting S/No per subaperture on the detector is about 0.25 or 7.2 mag in $K^\\prime$."
    },
    "0808/0808.2074_arXiv.txt": {
        "abstract": "Due to the high efficiency of planet detections, current microlensing planet searches focus on high-magnification events. High-magnification events are sensitive to remote binary companions as well and thus a sample of wide-separation binaries are expected to be collected as a byproduct.  In this paper, we show that characterizing binaries for a portion of this sample will be difficult due to the degeneracy of the binary-lensing parameters.  This degeneracy arises because the perturbation induced by the binary companion is well approximated by the Chang-Refsdal lensing for binaries with separations greater than a certain limit.  For binaries composed of equal mass lenses, we find that the lens binarity can be noticed up to the separations of $\\sim 60$ times of the Einstein radius corresponding to the mass of each lens. Among these binaries, however, we find that the lensing parameters can be determined only for a portion of binaries with separations less than $\\sim 20$ times of the Einstein radius. ",
        "introduction": "Searches for extrasolar planets by using microlensing are being carried out by observing stars located toward the Galactic bulge field.  The lensing signal of a planet is a short-duration perturbation to the smooth standard light curve of the primary-induced lensing event occurring on a background star \\citep{mao91, gould92}.  For the detections of the short-duration planetary lensing signals, these searches are using a combination of survey and follow-up observations, where the survey observations (e.g., OGLE, \\citet{udalski03}; MOA, \\citet{bond02}) aim to maximize the lensing event rate by monitoring a large area of sky and the follow-up observations (e.g., PLANET, \\citet{albrow01}; MicroFUN, \\citet{dong06}) intensively monitor the events alerted by the survey observations.  However, the limited number of telescopes restricts the number of events that can be followed at any given time and thus priority is given to those events that will maximize the planetary detection probability.  Currently, the highest priority is given to high-magnification events because the source trajectories of these events always pass close to the perturbation region around the planet-induced caustic located near the primary lens \\citep{griest98}. In addition, follow-up observations can be prepared for these events because the perturbation typically occurs near the peak of the event, which can be predicted from the data on the rising part of the event light curve.  As a result, six (OGLE-2005-BLG-071Lb, OGLE-2005-BLG-169Lb, OGLE-2006-BLG-109Lb,c, MOA-2007-BLG-400b, MOA-2007-BLG-192) of the eight reported microlensing planets were detected through the channel of high-magnification events \\citep{bond04, udalski05, beaulieu06, gould06, gaudi08, dong08, bennett08}.   In addition to planets, high-magnification events are sensitive to wide-separation binary companions as well.  Similar to the planetary case, the companion of a wide-separation binary induces a small caustic close to the primary lens and thus can produce a short-duration perturbation near the peak of a high-magnification event.  Due to the different nature of the companions, however, the perturbations induced by a wide-separation binary companion and a planet can be distinguished \\citep{albrow02, han08}. Therefore, under the current planetary lensing strategy focusing on high-magnification events, a considerable number of wide-separation binaries are expected to be detected as a byproduct and this sample might provide useful information about the physical distribution of binaries.   Unlike this expectation, however, we find that characterizing binaries for a significant portion of the binary lens sample will be difficult due to the degeneracy of the binary-lensing parameters.  This degeneracy arises because the perturbation induced by the binary companion is well approximated by the Chang-Refsdal lensing (hereafter C-R lensing) for binaries with separations greater than a certain value.  The C-R lensing represents single-mass lensing superposed on a uniform background shear.  For a wide-separation binary, the shear results from the combination of the binary-lens parameters and thus the individual parameters cannot be separately determined.   The paper is organized as follows.  In \\S\\ 2, we briefly describe the properties of binary and C-R lensing. In \\S\\ 3, we demonstrate the proximity between the binary and C-R lensing for binaries with separations beyond a certain limit.  We then set the range in the binary-lensing parameter space where the degeneracy of binary-lensing parameters occurs.  We discuss the meaning of this degeneracy in the studies of binaries by using microlensing. We summarize the results and conclude in \\S\\ 4. ",
        "conclusions": "By investigating the patterns of central perturbations produced by wide-separation binary companions, we found that the perturbation can be noticed and thus lens binarity can be identified for binaries with considerable separations.  However, we also found that perturbations for a significant portion of these binaries can be well approximated by the C-R lensing, implying that the individual binary-lensing parameters cannot be separately determined and thus characterization of the binary is difficult.  We set the range in the binary-lensing parameter space where this degeneracy occurs.  From this, we found that the lens binarity can be noticed up to the separations of $\\sim 60$ times of the Einstein radius corresponding to the mass of each lens of a binary composed of equal mass lenses.  Among these binaries, however, we found that the lensing parameters can be determined only for a portion of binaries with separations less than $\\sim 20$ times of the Einstein radius."
    },
    "0808/0808.0889_arXiv.txt": {
        "abstract": "We report the detection of very high-energy $\\gamma$-ray emission from the intermediate-frequency-peaked BL\\,Lacertae object W\\,Comae ($z = 0.102$) by VERITAS, an array of four imaging atmospheric-Cherenkov telescopes. The source was observed between January and April 2008. A strong outburst of $\\gamma$-ray emission was measured in the middle of March, lasting for only four days. The energy spectrum measured during the two highest flare nights is fit by a power-law and is found to be very steep, with a differential photon spectral index of $\\Gamma = 3.81 \\pm 0.35_{\\rm{stat}} \\pm 0.34_{\\rm{syst}}$. The integral photon flux above $200 \\, \\rm{GeV}$ during those two nights corresponds to roughly $9\\%$ of the flux from the Crab Nebula. Quasi-simultaneous Swift observations at X-ray energies were triggered by the VERITAS observations. The spectral energy distribution of the flare data can be described by synchrotron-self-Compton (SSC) or external-Compton (EC) leptonic jet models, with the latter offering a more natural set of parameters to fit the data. ",
        "introduction": "The blazars detected at very high energies (VHE, $E > 100 \\, \\rm{GeV}$) by ground-based imaging atmospheric-Cherenkov telescopes (IACTs) are extreme objects in the active galactic nuclei (AGN) population. Typically these sources show core-dominated emission, and they are characterized by rapid variability and strong broadband continuum emission ranging from the radio band to the X-ray band. Multi-wavelength data on blazars reveal that their spectral energy distribution (SED) is characterized by two broad, well-separated ``humps'' arising from synchrotron (low-energy) and inverse-Compton (IC) or hadronic emission (high-energy). Blazars are categorized into different sub-classes based on the frequencies at which these emission components reach a maximum. Flat-spectrum radio quasars (FSRQs) and low-frequency-peaked BL\\,Lacs (LBLs) are generally seen to have low-frequency, synchrotron peaks in the IR/optical regime, whereas high-frequency-peaked BL\\,Lacs (HBLs) exhibit peaks in the X-ray band, in several cases at energies of $\\sim$100~keV or higher. Intermediate-frequency-peaked BL\\,Lacs (IBLs) bridge the gap between LBLs and HBLs. The properties of the broad sub-classes of blazars, the luminosity-versus-frequency trends and possible physical explanations are discussed by \\cite{ghi08}.  Gamma rays are an important component of the SED of blazars; the integral power in the $\\gamma$-ray waveband is comparable or higher than that in the rest of the electromagnetic spectrum (from radio to X-rays). There are $65$ blazars detected at MeV/GeV energies by the EGRET instrument on board the {\\sl Compton Gamma Ray Observatory (CGRO)} \\citep{mat01, har99} and several of the other EGRET sources also have likely blazar counterparts. Ground-based IACTs have established $\\sim$20 blazars as emitters of VHE $\\gamma$-radiation; see for example \\citet{wak07}. While the blazars detected at MeV/GeV energies tend to be largely FSRQs (and some LBLs), almost all VHE blazars belong to the class of HBLs, the only exceptions being the LBLs BL\\,Lacertae \\citep{alb07}, S5\\,0716+71 \\citep{tes08} and the FSRQ 3C\\,279 \\citep{tes07}.  The IBL W\\,Comae (W\\,Com) at a redshift of $z = 0.102$ has long been an object of interest for VHE observatories. W\\,Com was discovered at radio frequencies \\citep{bir71} and later detected at X-ray energies by the {\\sl Einstein} Imaging Proportional Counter in June 1980 \\citep{wor90}. Data taken with the {\\sl BeppoSAX} satellite in 1998 \\citep{tag00} clearly showed that the transition between the low- and high-energy peaks in the SED occurs around $\\sim$4~keV. In April-May 1998, an exceptional optical outburst was detected from W\\,Com showing rapid variability on time scales of hours \\citep{mas99}.  At $\\gamma$-ray energies, W\\,Com was detected by EGRET in the $100 \\, \\rm{MeV} - 10 \\, \\rm{GeV}$ band \\citep{har99} and in a re-analysis of the data up to $25 \\, \\rm{GeV}$ \\citep{din01}. Due to its rather hard EGRET spectrum (photon spectral index $\\alpha=1.73\\pm 0.18$), with no sign of spectral cut-off \\citep{har99}, the source became even more interesting for VHE observations. However, W\\,Com was not detected by the Whipple IACT above $300 \\, \\rm{GeV}$ in 1993/94 \\citep{ker95} and 1995/96/98 \\citep{hor04}, nor by the solar heliostat Cherenkov telescope STACEE \\citep{sca04}. In this paper we report the discovery of VHE $\\gamma$-ray emission from W\\,Com with VERITAS. ",
        "conclusions": "VERITAS detected VHE $\\gamma$-ray emission from W\\,Com with a statistical significance of $4.9$~standard deviations for the entire data set (January to April 2008). A strong outburst was observed in March 2008 with a statistical significance of $>8$~standard deviations, that lasted for only four days. In addition to W\\,Com, a second extragalctic source (the VHE blazar 1ES\\,1218+304) is detected in the same field of view -- for the first time in VHE $\\gamma$-ray astronomy.  W\\,Com is the first VHE-detected blazar of the IBL class. The extension of the VHE catalog to the FRSQ, LBL and IBL classes will play a major role in our understanding of blazar populations and dynamics. The quasi-simultaneous SED of W\\,Com at the time of the VHE outburst can be modeled with a simple one-zone SSC model. However, an unusually low magnetic field of $B = 0.007 \\, \\rm{G}$ (more than an order of magnitude lower than typically found in the modeling of other BL\\,Lac-type blazars) and a small ratio of the magnetic field to electron energy density of $\\sigma = 1.3 \\cdot 10^{-3}$ are required. An EC model with more natural parameters ($B=0.36 \\, \\rm{G}$ and $\\sigma = 1$) provides a good fit and could account for shorter variability time scales. Our model results agree with the expectation that for IBLs (and LBLs) the higher optical luminosity plays an important role in providing the seed population for IC scattering.  The IBL W\\,Com will be an excellent target for future observations at GeV energies with GLAST and in the VHE regime with IACTs, including correlated GeV/TeV variability studies."
    },
    "0808/0808.2248_arXiv.txt": {
        "abstract": "(1270) Datura is the largest member of a very young asteroid cluster that was thought to be broken-up 0.45 Myr ago. The light-curve and the rotation-resolved reflectance spectra (0.6 $\\mu$m -- 1.0 $\\mu$m) were observed in order to find ``fresh'' surface. Our data show no significant spectral variation along the rotation phase. The depth of the 0.95 $\\mu$m absorption band, which indicates the degree of space weathering, was similar to that of an old S-type asteroid. This suggests that the reflectance spectrum in this wavelength range changes rapidly and saturates the depth of the 0.95 $\\mu$m absorption in less than 0.45 Myr in the main belt environment. ",
        "introduction": "A parent body of an ordinary chondrite, one of the most common stony meteorites, is believed to be a S-type asteroid that is commonly distributed in the inner main belt region and is thought to consist of silicate rocks covered with regolith. However, there are discrepancies between the reflectance spectra of ordinary chondrites and S-type asteroids: S-type asteroids show redder spectral slope in the visible wavelength and the depth of the absorption band at $\\sim0.95$ $\\mu$m is shallower than that of ordinary chondrites \\citep{cha73}.  These spectral discrepancies are attributed to ``space weathering,'' which alter a fresh, chondritic spectra to the S-type asteroid spectra (e.g., Clark et al. 2002). Although details of the physical process of the space weathering have not been fully understood, it is thought that irradiation of high energy particles (solar wind or cosmic ray) and/or bombardment of micro-meteorites makes nanophase iron in the surface of regolith that cases the spectral change \\citep{ada71, hap00, pie00, sas01}.  One of the unsolved issues on space weathering is to determine the time scale of the weathering process. Dynamical investigations of asteroid families are enabling us to create a time scale of the history of asteroid evolution (e.g., Nesvorn\\'y et al. 2005). \\citet{nes05} related the family age and the spectral slope, and \\citet{wil08} revised their results and found that the time scale of the spectral slope change by space weathering was $570\\pm220$ Myr. However, \\citet{cha07} and \\citet{ver07} observed (832) Karin, whose age was estimated to be 5.8 Myr, and found that the depth of 0.95 $\\mu$m absorption band was as shallow as a typical S-type asteroid. These observations suggest that the depth of 0.95 $\\mu$m absorption band changes faster than the reddening of the spectral slope.  Recently \\citet{nes06a} found a very young asteroid cluster, Datura cluster, whose age was estimated to be $0.45\\pm0.05$ Myr. Datura cluster has been exposed for only $1/10$ of the age of the Karin cluster. Thus this cluster is a good target to examine the alteration speed by space weathering at the 0.95 $\\mu$m absorption band.  (1270) Datura has a diameter of about 11 km and is orbiting in the inner main belt region ($a=2.2$ AU). I have observed (1270) Datura, which is the largest member of the Datura cluster, expecting to find a ``fresh'' region on its surface because the space-weathered old surface was removed by the cluster-forming break-up and a fragment might expose a fresh, interior surface on the whole area or at least in some area. In this letter I will report the rotation-resolved reflectance spectra of (1270) Datura and its implication for the space weathering. ",
        "conclusions": "The rotation curve shows fairly large amplitude that means we are not seeing (1270) Datura from above its pole, thus we see most of its surface when observing nearly a full rotational period. The fact that there is no significant difference among the spectra obtained at different rotation phases indicates that the surface of (1270) Datura has no significant large scale spatial variation in composition and in the degree of space weathering. The rotation speed of (1270) Datura is slow enough, and the size is large enough to be covered with regolith that must be newly created at the cluster-forming collision at 0.45 Myr ago and then deposited onto the surface. Since the deposition time scale might be longer compared to the rotation period of (1270) Datura, the dusts/gravels covered evenly over the surface. Thus it is natural that the surface of (1270) Datura is almost homogeneous when observed in a large spatial scale as discussed by \\citet{nes06b} and \\citet{ver07} for (832) Karin.  To see the degree of space weathering of the surface of (1270) Datura, the phase-averaged spectrum is compared with that of an old S-type asteroid (15) Eunomia (estimated age is 2.5 Gyr \\citep{nes05}) and some of the ordinary chondrites as shown in Figure~\\ref{fig_comp}. Contrary to our expectation, the spectrum of (1270) Datura does not resemble those of ordinary chondrites, but is close to an old S-type asteroid (15) Eunomia. This means that the surface of (1270) Datura was already altered in the same degree with an old S-type asteroid in respect to the depth of 0.95 $\\mu$m absorption band.  Because the surface was renewed 0.45 Myr ago by a cluster-forming collision, the change of the reflectance spectrum (from ordinary-chondrite-like one to that of S-type asteroids) must be completed in less than 0.45 Myr, if the age estimate by \\citet{nes06a} and the concept of space weathering are correct.   \\citet{bin96, bin01} found that many of the near-earth objects (NEOs) are Sq- or Q-type, which connects S-type spectra to ordinary chondrite spectra. If NEOs have been space weathered in the same manner as (1270) Datura, the surface of those objects have been exposed much less than $0.45$ Myr. If the surface was exposed by a recent collision in the main belt, we can estimate the current collision rate in the main belt. However, we must be careful that the environment of space weathering in NEO region differs from the inner main belt. Also the establishment and stability of regolith might be different since the size of the objects of \\citet{bin96, bin01} are one order-of-magnitude smaller than (1270) Datura."
    },
    "0808/0808.0417_arXiv.txt": {
        "abstract": "{In this  contribution  we discuss some   of the main problems in high energy astrophysics, and the perspectives to solve them using  different  types of  ``messengers'': cosmic  rays, photons  and neutrinos.}  \\normalsize\\baselineskip=15pt ",
        "introduction": "The  birth  of high  energy astrophysics  can  be traced to nearly a  century ago, when the balloon flights of Victor  Hess established  that  a form  of ionizing radiation, that  was soon given the name of  ``cosmic rays'', was   arriving from  outer space. Soon it was  demonstrated  that these ``rays''  are in fact ultrarelativistic  charged particles, mostly protons  and fully ionized nuclei\\footnote{ Electrons  have a  steeper  energy spectrum  and  constitute a fraction  of only few percent of  the  flux;   small  quantities of positrons  and  anti--protons  are also present.}, with  a  spectrum that  extends to  extraordinary  high energies. The existence of a large  flux of ultrarelativistic  particles arrived  completely  unexpected, and its origin  remained  a  mistery  that only now  is beginning  to be clarified. The  main reason   for this   very long   delay  in developing an  understanding of this  important physical phenomenon is that cosmic  rays (CR) do not point to their production  sites,  because their trajectory  is bent by the magnetic  fields   that permeates  both interstellar and intergalactic space.  Today   we know  that  our universe contains  several  classes of astrophysical  objects   where  non--thermal  processes are  capable to accelerate  charged particles, both  leptons  ($e^\\mp$)  or  hadrons (protons and nuclei), to very high energy.  These   relativistic particles, interacting inside or near  their  sources, can  produce   photons, and (in case of  hadrons) neutrinos  that  then  travel in  straight lines allowing  the imaging of the sources. At  the  highest energy also the  magnetic  bending of charged   particles can  become  sufficiently small to allow source imaging. A  detailed  study of the ``high energy universe''  can therefore in principle be performed  using three   different   ``messengers'': photons, neutrinos  and the cosmic ray themselves. This  ``multi--messenger'' approach is still in its infancy, but    has  the  potential to  give  us deep  insights about  the sites and physical  mechanisms  that  produce these very  high energy particles.  Many of  the  proposed  (or  detected)  acceleration sites are also associated with the violent acceleration of large macroscopic masses  (one example is  Gamma  Ray Bursts), and therefore a  fourth ``messenger'':  gravitational waves   has the potential to  give  very important and complementary information about these astrophysical environments. ",
        "conclusions": ""
    },
    "0808/0808.4002_arXiv.txt": {
        "abstract": "The main aim of this study is the comparison of gravitational waveforms obtained from numerical simulations which employ different numerical evolution approaches and different wave-extraction techniques. For this purpose, we evolve an oscillating, nonrotating, polytropic neutron-star model with two different approaches: a full nonlinear relativistic simulation (in three dimensions) and a linear simulation based on perturbation theory. The extraction of the gravitational-wave signal is performed via three methods: the gauge-invariant {\\it curvature-perturbation} theory based on the Newman-Penrose scalar $\\psi_4$; the gauge-invariant Regge-Wheeler-Zerilli-Moncrief {\\it metric-perturbation} theory of a Schwarzschild space-time; some generalization of the quadrupole emission formula. ",
        "introduction": "\\label{sec:intro}  The computation of the gravitational-wave emission from compact sources like supernova explosions, neutron-star oscillations and the inspiral and merger of two compact objects (like neutron stars or black holes) is one of the most lively subjects of current research in gravitational-wave astrophysics. This goal may be pursued using different numerical approaches. That is, (i) solving the {\\it full set} of coupled Einstein and matter equations; (ii) solving the {\\it linearized} Einstein and matter equations around a fixed background, when such an approximation is valid. In the latter case, with the additional condition of spherical symmetry, the formalism we employ is based on a multipolar expansion and the computation of the gravitational waves directly follows from the knowledge of the perturbative metric multipoles $k_\\lm$, $\\chi_\\lm$ and $\\psi_\\lm$. On the other hand, extracting gravitational waveforms from a space-time computed numerically in a given coordinate system is a highly nontrivial problem that has been addressed in various ways in the literature. In general, two routes have proven successful: (i) the gauge-invariant {\\it curvature-perturbation} theory based on the Newman-Penrose~\\cite{NP62} scalar $\\psi_4$, and (ii) the Regge and Wheeler~\\cite{RW57}, Zerilli~\\cite{Zerilli:1970se} theory of {\\it metric-perturbations} of a Schwarzschild space-time, recast in a gauge-invariant framework following the work of Moncrief~\\cite{Moncrief:1974am}.  The aim of our study is the computation of the gravitational waveforms emitted by the very controlled system constituted by a nonrotating polytropic relativistic star that oscillates nonisotropically around its spherically symmetric equilibrium configuration because of an axisymmetric perturbation. Our aim is to follow two (complementary) calculation procedures. On one hand, we perform a full 3+1 numerical simulation of the system, \\ie we compute a numerical solution of the Einstein equations without approximations except those of the numerical method itself. Because of its generality, this approach allows us to analyze different physical regimes, in particular, the case in which the ``perturbation'' is not small and nonlinear effects can play a relevant role with important consequences on the waveforms. On the other hand, we follow a perturbative approach based on the assumption that the perturbation is ``small''. If this is the case, one can (i) expand the metric around a fixed background (\\ie the Tolman-Oppenheimer-Volkoff solution), (ii) retain only the linear term of this expansion and (iii) solve the {\\it linearized} Einstein equations. In addition, since the star is nonrotating, one can factorize the angular dependence by means of a spherical-harmonic decomposition of the metric and matter fields, and, thus, only a 1+1 system of partial differential equations must be solved.  The present work has much in common with Refs.~\\cite{Shibata:2003aw,Pazos:2006kz}, where a comparison of different extraction techniques has been performed. Following the same inspiration of Ref.~\\cite{Pazos:2006kz}, we exploit perturbative computations to obtain ``exact'' waveforms to compare with the numerical-relativity--generated ones. As done in Ref.~\\cite{Shibata:2003aw}, we use an oscillating neutron star as a test-bed system, but we consider a wider range of possible wave-extraction techniques. Since there is a copious literature dealing with the problem of gravitational-wave extraction in numerical relativity, we prefer not to mention here the main bibliographic references, but rather to address the reader to the references in Refs.~\\cite{Shibata:2003aw,Pazos:2006kz} and to the citations in the following text.  The article is organized as follows. In Sec.~\\ref{sec:nsetp} we describe the numerical time-evolution methods and the gravitational-wave extraction techniques adopted. In Sec.~\\ref{sbsc:init} we introduce our choice of initial data and Sec.~\\ref{sec:res} is devoted to the presentation of our results. Conclusions that can be drawn from our results are discussed in Sec.~\\ref{sec:concl}.  Standard dimensionless units $c=G=M_\\odot=1$ and a spacelike signature $(-,+,+,+)$ are used. Greek indices are taken to run from $0$ to $3$, Latin indices from $1$ to $3$ and we adopt the standard convention for the summation over repeated indices. ",
        "conclusions": "\\label{sec:concl}  We have compared various gravitational-wave--extraction methods that are nowadays very popular in numerical-relativity simulations: (i) the Abrahams-Price~\\cite{Abrahams:1995gn} technique based on the gauge-invariant Regge-Wheeler-Zerilli-Moncrief perturbation theory of a Schwarzschild space-time; (ii) the extraction method based on Weyl curvature scalars, notably the $\\psi_4$ function; (iii) some (variations of) quadrupole-type formulas. We have applied these methods to extract gravitational radiation from 3D numerical-relativity simulations of the very controlled system represented by a neutron star (with polytropic EOS), that is oscillating nonradially due to an initial pressure perturbation. The simulations have been performed via the {\\tt Cactus-CCATIE-Carpet-Whisky} general-relativistic nonlinear code. This code evolves the full set of Einstein equations in full generality in the three spatial dimensions. The accuracy of the waveforms extracted from the simulations, using the three methods recalled above, has been assessed (for small perturbations) via a comparison with waveforms (assumed to be exact) computed by means of the {\\tt PerBACCo} perturbative code. This code is designed to evolve, in the time domain, the Einstein equations linearized around a TOV background. It is 1+1-dimensional (\\ie one temporal and one spatial dimension) and adopts a constrained evolution scheme. This latter choice allows for the computation of very long and very accurate time series and similarly accurate waveforms.  The initial pressure perturbation $\\delta p$ is given as an ``approximate'' eigenfunction of the star, whose maximum is a fraction of the central TOV pressure $p_c$. We focused only on $\\l=2$, $m=0$, quadrupolar deformations, but we analyzed four values of the perturbation in order to cover the transition from the linear to the nonlinear oscillatory regimes. We have first presented results of simulations done using only the 1D {\\tt PerBACCo} code to assess the accuracy of our exact waveforms. We have performed very long (about 1~s) and accurate simulations to extract both mode frequencies and damping times. We have analyzed finite-radius effects, finding that observers should be placed at extraction radius $r>200M$ in order to have amplitude errors below 1.6\\%.  In doing 3D simulations in the perturbative regime, ($10^{-3}\\lesssim {\\rm max}(\\delta p/p_c)\\lesssim 10^{-2}$), we have found that {\\it both} metric and curvature wave-extraction techniques generate waveforms that are consistent, both in amplitude and phasing, with the perturbative results. Each method, however, was found to have drawbacks. On one hand, the Zerilli-Moncrief function presents an unphysical burst in the early part of the waveform; on the other hand, the $\\psi_4$ scalar requires a polynomial correction to obtain the corresponding metric multipole. Our conclusion is that, in our setup, one needs both extraction methods to end up with accurate waveforms.  For larger values of the initial perturbation amplitude, nonlinear effects in the 3D general-relativistic simulations are clearly present. The effective relative amplitude of the main modes of the extracted gravitational wave is smaller for larger amplitudes of the initial perturbation, because of mode couplings. The Fourier spectra of the rest-mass--density projections [see Eq.~\\eqref{eq:rholm}] highlight that couplings between radial and quadrupolar fluid modes are present. Our study represents the first confirmation, in fully general-relativistic simulations, of the results of Ref.~\\cite{Passamonti:2007tm}, obtained via a perturbative approach.  In addition, we have shown that the (non-gauge-invariant) generalizations of the standard Newtonian quadrupole formula that we have considered can be useful tools to obtain accurate estimates of the frequency of oscillation. By contrast, amplitudes are always significantly under/overestimated, consistently with precedent observations of Refs.~\\cite{Shibata:2003aw,Nagar:2005cj}.   Finally, we discussed in detail some systematic errors that occur in the early part of the waveform extracted \\'a la Abrahams-Price. These errors show up, in the early part of the Zerilli-Moncrief function, in the form of a burst of {\\it junk} radiation whose amplitude grows linearly with the extraction radius. We have proposed some heuristic explanation of this fact and reproduced a similar behavior in low-accuracy perturbative simulations. Globally, our conclusion is that the extraction of the Zerilli-Moncrief function from a numerical-relativity simulation can be a delicate issue: small errors can conspire to give totally nonsensical results. Typically, these errors will show up as parts of the waveform whose amplitude grows with the observer's radius. We have also implemented the generalized wave-extraction approach based on the formalism of Refs.~\\cite{Sarbach:2001qq,Martel:2005ir,Pazos:2006kz,Korobkin:2008ji}, without any evident benefit. Note, however, that these kind of problems encountered with the Abrahams-Price wave-extraction procedure (as well as with its generalized version) seem to appear {\\it specifically} in the presence of matter. In binary black-hole coalescence simulations curvature and metric waveforms seem to be fully consistent~\\cite{Pollney:2007ss}. This last remark leads us to suggest that the Abrahams-Price wave-extraction technique, a ``standardized'' and very basic procedure and infrastructure that has been developed long ago (and tested at the time) for specific applications to black-hole physics, should be rethought and reanalyzed when the Einstein equations are coupled to matter. For this reason, in the presence of matter, since systematic errors could be hard to detect and are present already in the simplest cases, we strongly encourage the community to make use of {\\it both} wave-extraction techniques (curvature as well as metric perturbations) and to be always prepared to expect inaccuracies in the metric waveforms. In addition, concerning the many advantages related to extracting the metric waveforms directly from the space-time, we believe that it is also urgent and important for the community to have reliable implementations of the Abrahams-Price technique based on the Sarbach-Tiglio-Martel-Poisson~\\cite{Sarbach:2001qq,Martel:2005ir,Pazos:2006kz,Korobkin:2008ji} formalism."
    },
    "0808/0808.3414_arXiv.txt": {
        "abstract": "A set of 41 nearby stars (closer than 25 pc) is investigated which have very wide binary and common proper motion (CPM) companions at projected separations between $1000$ and $200\\,000$ AU. These companions are identified by astrometric positions and proper motions from the NOMAD catalog. Based mainly on measures of chromospheric and X-ray activity, age estimation is obtained for most of 85 identified companions. Color -- absolute magnitude diagrams are constructed to test if CPM companions are physically related to the primary nearby stars and have the same age. Our carefully selected sample includes three remote white dwarf companions to main sequence stars and two systems (55 Cnc and GJ 777A) of multiple planets and distant stellar companions. Ten new CPM companions, including three of extreme separations, are found. Multiple hierarchical systems are abundant; more than 25\\% of CPM components are spectroscopic or astrometric binaries or multiples themselves. Two new astrometric binaries are discovered among nearby CPM companions, GJ 264 and HIP 59000 and preliminary orbital solutions are presented. The Hyades kinematic group (or stream) is presented broadly in the sample, but we find few possible thick disk objects and none halo stars. It follows from our investigation that moderately young (age $\\lesssim 1$ Gyr) thin disk dwarfs are the dominating species in the near CPM systems, in general agreement with the premises of the dynamical survival paradigm. Some of the multiple stellar systems with remote CPM companions probably undergo the dynamical evolution on non-coplanar orbits, known as the Kozai cycle. ",
        "introduction": "\\label{firstpage} Components of wide stellar binaries and common proper motion pairs have been drawing considerable interest for many years. Despite the increasing accuracy of observations and the growing range of accessible wavelengths, the origin of very wide, weakly bound or unbound systems remains an open issue. \\citet{lep7} estimated that at least 9.5\\% of stars within 100 pc have companions with projected separations greater than 1000 AU. The renewed interest was boosted by the detection of a dearth of substellar mass companions in spectroscopic binaries, and by the attempts to account for the missing late-type members of the near solar neighborhood.  The main objective of this paper is to investigate a well-defined set of nearby stars in very wide CPM pairs and to discover new pairs, possibly with low-mass companions. The secondary goal of our investigation is to establish or estimate the age and evolutionary status of bona fide companions using a wide range of available astrometric and astrophysical data. The origin and status of wide CPM systems is still a mystery, because most of them are likely unbound or very weakly bound and are expected to be easily disrupted in dynamical interactions with other stars or molecular clouds (\\S~\\ref{surv.sec}). The nearest stars to the Sun, $\\alpha$ Cen and Proxima Cen, form a wide pair which may be on a hyperbolic orbit \\citep{ano}. It is expected that such systems should be mostly young, or belong to moving groups, remnant clusters or associations, but this has not yet been demonstrated on a representative sample. It is not known if the companions formed together and have the same age. We combine age-related parameters and data, including color-absolute magnitude diagrams (\\S~\\ref{hr.sec}), chromospheric activity indeces (\\S~\\ref{chr.sec}), coronal X-ray luminosity (\\S~\\ref{x.sec}), multiplicity parameters (\\S~\\ref{bin.sec}) and kinematics (\\S~\\ref{skg.sec}) to shed light on this problem.  Previous investigations in the field are too numerous to be listed, but a few papers in considerable overlap with this study should be mentioned. \\citet{pov} published a catalog of 305 nearby wide binary and 29 multiple systems. They discussed the importance of moving groups for separating different species of wide binaries and tentatively assigned 32 systems to the Hyades stream (called supercluster following Eggen's nomenclature), and 14 to the Sirius stream. \\citet{sal} undertook a comprehensive revision of the Luyten catalog for approximately 44\\% of the sky, drastically improving precision of epoch 2000 positions and proper motions, and supplying the stars with NIR magnitudes from 2MASS. This allowed \\citet{gou} to estimate, for the first time, trigonometric parallaxes of 424 common proper motion companions to Hipparcos stars with reliable parallaxes. This extrapolation of parallaxes to CPM companions is justified for high-proper motion pairs where the rate of chance alignments is small. There is significant overlap between the sample investigated in this paper and the catalog of \\citet{gou}, although we did not use the latter as a starting point for our selection. We are also employing the parallax extrapolation technique for dim companions not observed by Hipparcos when constructing color-magnitude diagrams in this paper. ",
        "conclusions": "\\label{dis.sec} One of the most interesting results of this paper is that we find little, if any, presence of thick disk or halo population in the local sample of very wide binaries. The only CPM system that may belong to the thick disk is the WD+dM4.5 pair HIP 65877 (DA3.5 white dwarf WD $1327-083$) and LHS 353, cf. \\citet{sil}. This shows that even the widest pairs at separations greater than $1000$ AU can survive for $\\simeq$ 1 Gyr staying constantly in the thin disk of the Galaxy, despite numerous encounter and dynamical interaction events. This observation does not refute the dynamical analysis presented in \\S~\\ref{surv.sec}, because thick disk and halo stars are very rare in the Solar vicinity, and our sample (based on bright stars in the Hipparcos catalog) is probably too small and incomplete to accommodate sufficient statistics. But if we boldly extrapolate this result to a wider part of the Galaxy, we conclude that normal thin disk, moderately young or very young, stars dominate wide CPM binary and multiple systems. Statistically, this is quite consistent with the estimation by \\citet{bart} on a larger sample of 804 Hipparcos visual systems, who found that $92\\%$ of systems belong to the thin disk (and are mostly young to middle-age), $7.6\\%$ to the thick disk, and much less than $1\\%$ to the halo. Further inroads in this study can be made by collecting a larger volume-limited sample of very wide binaries and a comparison with a representative set of nearby field stars. We consider this paper as an initial step in this direction.  Despite the considerable progress in recent years, chronology of solar-type stars is still in a rather pitiful state, and we find more evidence to this in the widely discrepant age estimates for a few CPM systems obtained with different methods. Although a significant fraction of CPM companions display enhanced chromospheric, X-ray and EUV activity, only few are patently in the pre-main-sequence stage of evolution (e.g., AT and AU Mic) where these signs of activity and the high rate of surface rotation can be attributed to a very young age. The origin of activity in most of our CPM systems lies in short-period binarity of their components, i.e., in hierarchical multiplicity. An interesting connection emerges between the presence of wide companions and the existence of short-period binaries. The reason for abundant multiple systems may partly be purely dynamical, in that the chances of survival are higher for systems with an internal binary because of the larger mass. An alternative astrophysical possibility is that the original fragmentation of a star-forming core takes place at various spatial scales and tends to produce multiple stellar systems, of which only hierarchical ones can survive for an appreciably long time.  Apparently, the time-scale of dynamical survival of wide companions (of order 1 Gyr) is sufficiently long compared to the time-scale of dynamical evolution of non-coplanar multiple systems (the Kozai cycle, \\S \\ref{kozai.sec}) for the latter to shape up the present-day systems. The existence of circularized spectroscopic binaries with periods less than a few days may be the direct consequence of the interaction with remote companions, followed by the tidal friction and loss of angular momentum \\citep{eggl}. Ultimately, the inner components will form a contact binary and then merge. The existence of CPM multiple systems in a wide range of ages and separations will allow us to investigate this process in detail as it unfolds. Indeed, even in our sample of modest size we find examples of inner pairs of intermediate periods and large eccentricities, which are apparently evolving toward the tidally circularized state. The Kozai-type mechanism can affect the dynamical stability and composition of planetary systems. We find two stars in our sample with multiple planets (55 Cnc and GJ 777 A), and both have interesting dynamical properties very much unlike our Solar system."
    },
    "0808/0808.0371_arXiv.txt": {
        "abstract": "We report high angular resolution observations of the HCN (3-2) line emission in the circumstellar envelope of the O-rich star W~Hya with the Submillimeter Array. The proximity of this star allows us to image its molecular envelope with a spatial resolution of just $\\sim 40$ AU, corresponding to about 10 times the stellar diameter. We resolve the HCN (3-2) emission and find that it is centrally peaked and has a roughly spherically symmetrical distribution. This shows that HCN is formed in the innermost region of the envelope (within $\\sim 10$ stellar radii), which is consistent with predictions from pulsation-driven shock chemistry models, and rules out the scenario in which HCN forms through photochemical reactions in the outer envelope. Our model suggests that the envelope decreases steeply in temperature and increases smoothly in velocity with radius, inconsistent with the standard model for mass-loss driven by radiative pressure on dust grains. We detect a velocity gradient of $\\sim 5$ km~s$^{-1}$ in the NW--SE direction over the central 40 AU. This velocity gradient is reminescent of that seen in OH maser lines, and could be caused by the rotation of the envelope or by a weak bipolar outflow. ",
        "introduction": "The detection at radio frequencies of various molecules in the circumstellar envelopes of evolved stars has considerably improved our understanding of circumstellar chemistry. This is particularly true for the prototype of carbon stars, IRC+10216, on which much attention has been focused. Several dozens of molecules have been detected in its circumstellar envelope (e.g., \\citealt{cer00}), reflecting its complex and rich chemistry. The inventory of molecules together with high angular resolution images of their spatial distribution have led to the construction of sophisticated chemical models for the circumstellar envelope of IRC+10216 (e.g., \\citealt{gla87,mil94}) that now form the basis of our understanding of the chemistry in C-rich envelopes.  The chemistry in O-rich envelopes, on the other hand, was not expected to be as rich because equilibrium chemistry calculations indicate that nearly all the carbon should be locked into CO and nitrogen into N$_2$ (\\citealt{tsu73}). Indeed, before 1985, the inventory of molecules detected around oxygen stars was sparse and did not include carbon-bearing molecules, except of CO. \\cite{deg85} reported the first detection of HCN towards O-rich stars. The list of HCN detections was later extended by \\cite{jew86}, \\cite{lin88}, \\cite{ner89}, \\cite{olo98} and \\cite{bie00}. Recently, \\cite{ziu07} also emphasized the chemical complexity in the envelope of the O-rich supergiant star VY~CMa, following the detection of various molecules such as HCO$^+$, CS, NaCl, PN and SiS.  Two competing models have been proposed to account for the formation of HCN in O-rich envelopes: photochemical reactions in the outer envelope (\\citealt{cha95}), or gas-phase non-equilibrium chemical reactions in the inner region close to the stellar photosphere (within few tens of AU from the star) due to shocks driven by stellar pulsation (\\citealt{dua99,dua00}). The two models predict very different HCN spatial distribution: in a hollow shell-like structure of radius 200 to 1000 AU depending on the mass loss for photochemical reactions, or a centrally concentrated distribution extending up to the photodissociation radius of HCN for the shock-driven chemical reactions.  Current observational data seem to favor pulsation-driven shock chemistry formation of HCN in O-rich stars. \\cite{bie00} conducted single-dish observations toward a sample of 16 O-rich stars in the HCN (3-2) and (4-3) lines. The detection of these high density tracers is inconsistent with photochemical production of HCN in the outer envelope. The detection of the HCN (3-2) (0,1$^{1c}$,0) and (8-7) transitions toward $\\chi$ Cyg gave further support for the formation of HCN in the innermost region of the envelope, within $\\sim 20$ stellar radii (i.e. $\\sim 30$ AU), based on excitation arguments (\\citealt{dua00}). \\cite{mar05} reported the first interferometric observations of HCN in the circumstellar envelopes of the O-rich stars IK~Tau and TX~Cam. In both stars, the HCN emission appears to be largely concentrated within $\\sim 800$ AU, as predicted by shock chemistry models.    If HCN indeed forms close to the stellar photosphere, it should be an excellent tracer of the inner envelope, and perhaps even of the wind acceleration zone where, in the standard model, dust particles form and radiation pressure on dust becomes most effective. Maser (e.g., H$_2$O or OH) emission in the inner envelope can be observed with very long baseline interferometry, but traces only a limited radial range. On the other hand, there are so far only few reports of high angular resolution observations of other thermal lines in the inner envelope, such as SiO (2-1) $\\nu = 0$ (\\citealt{luc92,sah93,sch04}), none of them with angular resolution higher than $\\sim 100$ AU.     In this letter, we report very high spatial resolution ($\\sim 40$ AU) observations of the HCN (3-2) emission toward W~Hya with the Submillimeter Array. W~Hya is one of the closest O-rich stars at a distance of 78 pc (based on Hipparcos data, \\citealt{per97}, revised by \\citealt{kna03}), and shows the brightest HCN (3-2) and (4-3) emission amongst the sample of O-rich stars observed by \\cite{bie00}. It is therefore an excellent target in which to investigate the distribution and kinematics of HCN in the circumstellar envelope of an O-rich star. ",
        "conclusions": ""
    },
    "0808/0808.0692_arXiv.txt": {
        "abstract": "We present {\\it Chandra} ACIS X-ray observations of the Galactic supernova remnant Cassiopeia A taken in December 2007. Combining these data with previous archival {\\it Chandra} observations taken in 2000, 2002, and 2004, we estimate the remnant's forward shock velocity at various points around the outermost shell to range between 4200 and 5200 $\\pm500 $ km s$^{-1}$.  Using these results together with previous analyses of Cas A's X-ray emission, we present a model for the evolution of Cas A and find that it's expansion is well fit by a $\\rho_{ej} \\propto r^{-(7-9)}$ ejecta profile running into a circumstellar wind. We further find that while the position of the reverse shock in this model is consistent with that measured in the X-rays, in order to match the forward shock velocity and radius we had to assume that $\\sim$ 30\\% of the explosion energy has gone into accelerating cosmic rays at the forward shock.  The new X-ray images also show that brightness variations can occur for some forward shock filaments like that seen for several nonthermal filaments seen projected in the interior of the remnant. Spectral fits to exterior forward shock filaments and interior nonthermal filaments show that they exhibit similar spectra. This together with similar flux variations suggests that interior nonthermal filaments might be simply forward shock filaments seen in projection and not located at the reverse shock as has been recently proposed. ",
        "introduction": "Cassiopeia A (Cas A) is one of the youngest known Galactic supernova remnants (SNR) with an estimated explosion date no earlier than $1681 \\pm19$ \\citep{fesen06}. Optical echoes of the supernova outburst have been recently detected \\citep{Rest08}, the spectra of which indicate Cas~A is the remnant of a Type IIb supernova event \\citep{Krause08} probably from a red supergiant in the mass range of 15--25 M$_{\\odot}$ that may have lost much its hydrogen envelope to a binary interaction \\citep{Young06}.  Viewed in X-rays, the remnant consists of a line emitting shell arising from reverse shocked ejecta rich in O, Si, Ar, Ca, and Fe \\citep{fabian80,markert83,vink96,hughes00,will02,will03,hwang03,laming03}. Exterior to this shell are faint X-ray filaments which mark the current position of the remnant's forward shock front. The emission found here is nonthermal X-ray synchrotron radiation as well as faint line emission from shocked circumstellar material (CSM).  \\citet{vink98} compared {\\sl Einstein} HRI to {\\sl ROSAT} HRI observations of Cas A to measure the expansion of the bright shell, finding an expansion age of $\\sim$ 500 yr, considerably less than the $\\sim$ 800 yr expansion age derived from 1.5 and 5.0 GHz radio observations \\citep{anderson95}, but similar to the 400--500 yr expansion age found by \\citet{agueros99} using data taken at 151 MHz.  More recently, \\citet{delaney03} using {\\sl Chandra} X-ray observations taken in 2000 and 2002 presented the first proper motion measurements of the forward blastwave velocity. Assuming a distance of 3.4 kpc \\citep{reed95}, they estimated a blast wave expansion velocity of $\\approx$ 5000 km s$^{-1}$.  Besides the outlying nonthermal emission filaments associated with the forward shock, some filamentary nonthermal X-ray emission is also seen in projection in the interior of the SNR \\citep{delaney04a}. Whether these interior filamentary emissions originate from a wrinkled forward shock seen in projection or arises from nonthermal emission mechanisms in the interior of the SNR is currently uncertain \\citep{laming01,uchiyama08,helder08}.  Comparisons of {\\it Chandra} observations taken in 2000, 2002, and 2004 revealed secular changes in several X-ray thermal knots and in one nonthermal filament projected in the remnant's interior \\citep{patnaude07}. \\citet{uchiyama08} using the same multi-epoch {\\it Chandra} observations found evidence for rapid variability in many more interior nonthermal X-ray emission filaments. Motivated by similar changes seen in RX J1713-3946 \\citep{uchiyama07}, they measured the time variability of selected filaments to determine the local magnetic field strength in the variable regions.  Their results suggest that the magnetic field in these regions is relatively high, B $\\sim$ 1 mG.  Such a high magnetic field strength would be consistent with equipartition field strengths inferred in observations of bright radio knots in the remnant \\citep{longair94,wright99}.  \\citet{uchiyama08} argue that their result points to a synchrotron origin for the emission coming from these knots, ruling out nonthermal bremsstrahlung from $\\sim$ 100 keV electrons \\citep{laming01}, and suggest that this is strong evidence for a hadronic origin to the TeV emission observed in Cas A \\citep{aharonian01,albert07}. Based on the location of the synchrotron knots, Uchiyama \\& Aharionian suggest that the emission is located primarily at the reverse shock, and \\citet{helder08} reach a similar conclusion.  Here we present forward shock velocity measurements using new {\\it Chandra} ACIS observations of Cas A taken in December 2007 and compare these results to models for SNR evolution with and without efficient shock acceleration.  The new observations show that many nonthermal emission filaments and features have undergone substantial brightness variations over the last four years.  Model fits to the nonthermal emission coming from both the forward shock and the interior filaments indicate that they are quantitatively similar.  We also present evidence for fast variability in forward shock front filaments which argues against the conclusion that rapid variability is a property restricted to emission at the reverse shock. ",
        "conclusions": "We have presented new {\\it Chandra} ACIS observations of Cas A which were taken in late 2007. These new observations, when combined with previous {\\it Chandra} data, allow us to constrain the velocity of the forward shock to be about 4900 km s$^{-1}$.  Combined with results from previous analyses of Cas A's X-ray emission \\citep{laming03,gotthelf01}, we present several models for the evolution of Cas A and find that it's expansion can be well modeled by an $n= 7 - 9$ ejecta profile running into a circumstellar wind.  We also find that the position of the reverse shock in this model is consistent with that measured by \\citet{gotthelf01}. However, in order to match the radius of the forward shock, we found that we must assume that the forward shock is efficiently accelerating cosmic rays.  Rapid changes in Cas A's synchrotron emission are seen for interior and exterior projected filaments, with both showing similar nonthermal spectra as well as inferred magnetic field strengths.  Based on this and the simulations presented by \\citet{bykov08}, it is currently not clear whether the interior filaments are in fact located at the reverse shock as recently argued by \\citet{uchiyama08} and \\citet{helder08}.  Instead, we propose that the interior filaments might be forward shocks seen in projection \\citep{delaney04b}.  In that case, the observed brightness variations might arise from wrinkles in front-facing, forward shock as it moves through an inhomogeneous, local circumstellar medium.  Although we cannot rule out the possibility that interior nonthermal filaments are associated with the reverse shock, the combination of similar spectra, flaring timescale, and our fits to the remnant's dynamics are suggestive that the observed synchrotron flaring for interior filaments arises from forward shock filaments seen in projection toward Cas A's interior rather than at the reverse shock as recently suggested.  At the least, our new X-ray data of Cas A shows that rapid brightness variations like those seen for interior nonthermal filaments can also be exhibited by some outer, nonthermal forward shock filaments."
    },
    "0808/0808.2973_arXiv.txt": {
        "abstract": "We present high-angular resolution sub-millimeter continuum images and molecular line spectra obtained with the Submillimeter Array toward two massive cores that lie within Infrared Dark Clouds; one actively star-forming (G034.43+00.24~MM1) and the other more quiescent (G028.53$-$00.25~MM1). The high-angular resolution sub-millimeter continuum image of G034.43+00.24~MM1 reveals a compact ($\\sim$ 0.03 pc) and massive ($\\sim$ 29\\,\\Msun) structure while the molecular line spectrum shows emission from numerous complex molecules. Such a rich molecular line spectrum from a compact region indicates that G034.43+00.24~MM1 contains a hot molecular core, an early stage in the formation of a high-mass protostar. Moreover, the velocity structure of its $^{13}${\\rmfamily CO}{\\,(3--2)} emission indicates that this B0 protostar may be surrounded by a rotating circumstellar envelope. In contrast, the sub-millimeter continuum image of G028.53$-$00.25~MM1 reveals three compact ($\\lesssim$ 0.06 pc), massive (9--21\\,\\Msun) condensations but with no lines detected in its spectrum. We suggest that the core G028.53$-$00.25~MM1 is in a very early stage in the high-mass star-formation process; its size and mass are sufficient to form at least one high-mass star, yet it shows no signs of localized heating. Because the combination of high velocity line wings with a large IR--mm bolometric luminosity ($\\sim$~10$^{2}$\\,\\Lsun) indicates that this core has already begun to form accreting protostars, we speculate that the condensations may be in the early phase of accretion and may eventually become high-mass protostars.  We, therefore, have found the possible existence of two high-mass star-forming cores; one in a very early phase of star-formation and one in the later hot core phase.  Together the properties of these two cores support the idea that the earliest stages of high-mass star-formation occur within IRDCs. ",
        "introduction": "Emerging evidence suggests that the very earliest stages of high-mass star and cluster formation occur within cold, dense molecular clumps called infrared dark clouds (IRDCs; \\citealp{Simon-msxgrs,Rathborne06,Rathborne07}). Because these molecular clumps have very high column densities ($\\sim 10^{23}$--$10^{25}$ cm$^{-2}$) and low temperatures ($<25$ K), they have predominantly been identified via their absorption of the background Galactic emission at mid-IR wavelengths (e.g., \\citealp{Egan98,Carey98,Carey00,Simon-catalog}).  Because IRDCs are the densest clumps in molecular clouds \\citep{Simon-msxgrs} which are now undergoing the process of fragmentation and condensation, IRDCs are important laboratories to study the pristine, undisturbed physical conditions of cluster-forming clouds before they are shredded apart by stellar winds and radiation.  Millimeter/sub-millimeter continuum studies of IRDCs \\citep{Lis94,Carey00,Garay04,Ormel05,Beuther05,Rathborne05,Rathborne06} reveal that IRDCs contain many compact cores. These cores have typical sizes of $<$~0.5~pc and masses of $\\sim$~120\\,\\Msun\\, \\citep{Rathborne06}.  While most of these cores contain little evidence for active star formation, some are associated with bright 24\\,\\um\\, emission, broad molecular line emission, shocked gas, and maser emission, indicating that they are actively forming stars (e.g., \\citealp{Rathborne05,Wang06,Chambers08}). Indeed, a number of low- and intermediate-mass protostars \\citep{Carey00,Redman03} in addition to high-mass protostars \\citep{Beuther05,Rathborne05,Pillai06,Wang06} have been identified within IRDCs.  Hot Molecular Cores (HMCs) are associated with the early stages of high-mass star formation. They correspond to the stage immediately after a cold dense starless core has formed. During the later stages in their evolution, HMCs are often also associated with methanol masers and ultra-compact \\hii\\, regions. Numerous examples of HMCs have been found throughout the Galaxy (e.g., \\citealp{Garay99,Kurtz00,Churchwell02}), including one within an IRDC \\citep{Rathborne07}. HMCs are compact ($<$0.1 pc), dense ($\\sim$10$^{5}-10^{8}$\\,\\cmc), and massive ($\\sim$10$^{2}$\\,\\Msun;\\citealp{Garay99,Kurtz00,Churchwell02}). Due to the internal heating from the central high-mass protostar (T$\\sim$100\\,K), they are typically very luminous ($>$10$^3$\\,\\Lsun) and show strong emission from complex molecules (e.g., \\citealp{Kurtz00}). Accretion disks are also inferred in the later stage of the HMC phase directly \\citep{Zhang05a,Cesaroni07} or indirectly \\citep{Kurtz00,Beuther02,Zhang01,Zhang05b} by the presence of molecular outflows and maser emission.  The existence of HMCs in IRDCs establishes a possible link between IRDCs and the early stages of high-mass star formation \\citep{Rathborne07}.  To understand the earliest stages in high-mass star formation one needs to identify and study the so-called `high-mass starless cores,' the immediate precursors to high-mass protostars, HMCs, and ultra-compact \\hii\\, regions. High-mass star formation is rare and occurs rapidly; thus, the identification of high-mass starless cores is difficult. Because IRDCs have sizes, masses, and densities similar to cluster-forming molecular clumps but are considerably colder, we suggest that they are the precursors to clusters and their dense, cold cores the precursors to stars. Thus, the high-mass cores within IRDCs that contain no evidence for star formation are good candidates for the elusive `high-mass starless cores.'  Here we present interferometric sub-millimeter continuum and molecular line observations obtained with the Submillimeter Array toward two cores within IRDCs; one showing evidence for active high-mass star formation (\\irdcfortythree), the other showing none (\\irdcthirty). Both cores are at distances d $>$ 3.7\\,kpc and have sizes R $<$0.8 pc, gas masses M $>$400\\,\\Msun, and bolometric luminosities L $>$10$^{2}$\\,\\Lsun\\, (Table~\\ref{core-properties}; \\citealp{Simon-msxgrs,Rathborne06,Rathborne08}). We speculate that the `active' core is in a later evolutionary stage than the more quiescent core. The new interferometric data reveal distinct properties for these cores and suggest that \\irdcfortythree\\, is a more evolved high-mass protostar in the HMC phase and that \\irdcthirty\\, may be in a very early `starless' core phase of high-mass star-formation. ",
        "conclusions": "Using the SMA we have obtained high-angular resolution sub-millimeter continuum images and molecular line spectra toward two cores within IRDCs. In the low-angular resolution data these cores have sizes R $<$0.8\\,pc, gas masses M $>$400\\,\\Msun, and bolometric luminosities L~$\\gtrsim$10$^{2}$\\,\\Lsun\\, \\citep{Rathborne06,Rathborne08}.  Despite their similar sizes, masses, and bolometric luminosities the two cores show remarkable differences in the high-angular resolution data. The sub-millimeter continuum images reveal that \\irdcfortythree\\,remains unresolved and has a size of $\\sim$ 0.03 pc and mass of 29\\,\\Msun. In contrast, \\irdcthirty\\, contains at least three compact ($\\sim$ 0.06 pc), massive (9--21\\,\\Msun) condensations.  Moreover, the molecular line spectra reveal that \\irdcfortythree\\, emits strong lines from many complex molecules, but the core \\irdcthirty\\, does not. This suggests that \\irdcfortythree\\, is hot, compact, and dense, while \\irdcthirty\\, is colder and possibly more extended. We speculate, therefore, that \\irdcfortythree\\, is a HMC in an early stage in the formation of a high-mass protostar, while \\irdcthirty\\, may represent an even earlier phase in the high-mass star-formation process.  Because \\irdcthirty\\, shows some evidence for active star formation in the low-angular resolution data (e.g., shocked gas, high velocity line profiles, and a high bolometric luminosity), we cannot rule out the possibility that the condensations within this core may have already begun to accrete. However, the low temperatures and absence of bright 24\\,\\um\\, point sources or strong molecular line emission from high excitation transitions suggests that there is little localized heating. Because of its size and mass, \\irdcthirty\\, is likely to give rise to at least one high-mass star. We speculate that the detected condensations may continue to accrete material from their surroundings to eventually form high-mass stars. Since the condensations account for only $\\sim$4\\% of the total mass, there is an ample amount of supply of material to accrete.  The properties of these cores support the idea that high-mass star and cluster formation likely occurs in IRDCs. High-angular sub-millimeter continuum images and molecular line spectra toward more cores within IRDCs are needed to unambiguously establish the evolutionary path from a high-mass starless core to a high-mass protostar."
    },
    "0808/0808.2718_arXiv.txt": {
        "abstract": "We study the effects of Supernova (SN) feedback on the formation of galaxies using hydrodynamical simulations in a $\\Lambda$CDM cosmology. We use an extended version of the code GADGET-2 which includes chemical enrichment and energy feedback by Type II and Type Ia SN, metal-dependent cooling and a multiphase model for the gas component. We focus on the effects of SN feedback on the star formation process, galaxy morphology, evolution of the specific angular momentum and chemical properties. We find that SN feedback plays a fundamental role in galaxy evolution, producing a self-regulated cycle for star formation, preventing the early consumption of gas and allowing disks to form at late times. The SN feedback model is able to reproduce the expected dependence on virial mass, with less massive systems being more strongly affected. ",
        "introduction": "Supernova  explosions play a fundamental role in galaxy formation and evolution. On one side, they are the main source of heavy elements in the Universe and the presence of such elements substantially enhances the cooling of gas (White \\& Frenk 1991). On the other hand, SNe eject a significant amount of energy into the interstellar medium. It is believed that SN explosions are responsible of generating a self-regulated cycle for star formation through the heating and disruption of cold gas clouds, as well as of triggering important galactic winds such as those observed (e.g. Martin 2004). Smaller systems are more strongly affected by SN feedback, because  their shallower potential wells  are less efficient in retaining baryons (e.g. White \\& Frenk 1991).   Numerical simulations have become an important tool to study galaxy formation, since  they can track the joint evolution of dark matter and baryons in the context of a cosmological model. However, this has shown to be an extremely complex task, because of the need to cover a large dynamical range and describe, at the same time, large-scale processes such as tidal interactions and mergers and small-scale processes related to stellar evolution.    One of the main problems that  galaxy formation simulations have repeteadly found is the inability to reproduce the  morphologies of disk galaxies observed in the Universe. This is generally refered to as the angular momentum problem that arises when baryons transfer most of their angular momentum to the dark matter components during interactions and mergers (Navarro \\& Benz 1991; Navarro \\& White 1994). As a result, disks are too small and concentrated with respect to real spirals. More recent simulations which include prescriptions for SN feedback have been able to produce more realistic disks (e.g. Abadi et al. 2003; Robertson et al. 2004; Governato et al. 2007). These works have pointed out the importance of SN feedback as a key process to prevent the loss of angular momentum, regulate the star formation activity and produce extended, young disk-like components.    In this work, we investigate the effects of SN feedback on the formation of galaxies, focusing on the formation of disks. For this purpose, we have run simulations of a Milky-Way type galaxy using an extended version of the code {\\small GADGET-2} which includes chemical enrichment and energy feedback by SN. A summary of the simulation code and the initial conditions is given in Section~\\ref{simus}. In Section~\\ref{results} we investigate the effects of SN feedback on galaxy morphology, star formation rates,  evolution of specific angular momentum and chemical properties. We also investigate the dependence of the results on virial mass. Finally, in Section~\\ref{conclusions} we give our conclusions. ",
        "conclusions": "We have run simulations of a Milky Way-type galaxy in its cosmological setting in order to investigate the effects of SN feedback on the formation of galaxy disks. We compare two simulations with the only difference being the inclusion of the SN energy feedback model of Scannapieco et al. (2005, 2006). Our main results can be summarized as follows:  \\begin{itemize} \\item{ SN feedback helps to settle a self-regulated cycle for star formation  in galaxies, through the heating and disruption of cold gas and the generation of galactic winds. The regulation of star formation  allows gas to be mantained in a hot halo which can condensate at late times, becoming a reservoir for recent star formation. This contributes significantly to the formation of disk components. } \\item{When SN feedback is included, the specific angular momentum of the baryons is  conserved and disks with the correct scale-lengths are obtained. This results from the late collapse of gas with high angular momentum, which becomes available to form stars at later times, when the system does not suffer from strong interactions. } \\item{ The injection of SN energy into the interstellar medium generates a redistribution of chemical elements in galaxies. If energy feedback is not considered, only the very central regions were stars are formed are contaminated. On the contrary, the inclusion of feedback triggers a redistribution of metals since gas is heated and expands, contaminating the outer regions of galaxies. In this case, metallicity profiles in agreement with observations are produced. } \\item{ Our model is able to reproduce the expected dependence of SN feedback on virial mass: as we go to less massive systems, SN feedback has stronger effects: the star formation rates (normalized to mass) are lower, and more violent winds develop. This proves that our model is well suited for studying the cosmological growth of structure where large systems are assembled through mergers of smaller substructures and systems form simultaneously over a wide range of scales. } \\end{itemize}"
    },
    "0808/0808.0001_arXiv.txt": {
        "abstract": "We present photometry for globular and open cluster stars observed with the Sloan Digital Sky Survey (SDSS).  In order to exploit over 100 million stellar objects with $r < 22.5$~mag observed by SDSS, we need to understand the characteristics of stars in the SDSS $ugriz$ filters. While star clusters provide important calibration samples for stellar colors, the regions close to globular clusters, where the fraction of field stars is smallest, are too crowded for the standard SDSS photometric pipeline to process.  To complement the SDSS imaging survey, we reduce the SDSS imaging data for crowded cluster fields using the DAOPHOT/ALLFRAME suite of programs and present photometry for 17 globular clusters and 3 open clusters in a SDSS value-added catalog. Our photometry and cluster fiducial sequences are on the native SDSS 2.5-meter $ugriz$ photometric system, and the fiducial sequences can be directly applied to the SDSS photometry without relying upon any transformations. Model photometry for red giant branch and main-sequence stars obtained by Girardi et~al.\\ cannot be matched simultaneously to fiducial sequences; their colors differ by $\\sim0.02$--$0.05$~mag.  Good agreement ($\\la0.02$~mag in colors) is found with Clem et~al.\\ empirical fiducial sequences in $u'g'r'i'z'$ when using the transformation equations in Tucker et~al. ",
        "introduction": "As single-age and (in most cases) single-metallicity populations, Galactic star clusters provide important calibration samples for exploring the relationships between stellar colors and absolute magnitudes as functions of stellar age and heavy-element content. These two observable properties of a star are related to fundamental physical parameters, such as the effective temperature ($T_{\\rm eff}$) and surface gravity ($\\log{g}$), as well as the metallicity.  The color and magnitude relations can be used to test stellar evolutionary theories, to interpret the observed distribution of stars in color-color and color-magnitude diagrams (CMDs), and to derive distances to stars and star clusters via photometric parallax or main-sequence (MS) fitting techniques \\citep[e.g.,][]{johnson:57}.  Because the relationships between magnitude, color, and fundamental stellar properties depend on the filters used, it is necessary to characterize these relations for each filter system.  Galactic globular and open clusters provide an ideal opportunity to achieve this goal because the same distance can be assumed for cluster members with a wide range of stellar masses.  Furthermore, observations of a large number of Galactic clusters can cover a wide range of the heavy-element content, providing an opportunity to explore the effects of metallicity on magnitudes and colors for each set of filter bandpasses.  Among previous and ongoing optical surveys, the Sloan Digital Sky Survey \\citep[SDSS;][]{york:00,edr,dr1,dr2,dr3,dr4,dr5,dr6} is the largest and most homogeneous database of stellar brightnesses currently available. The original goal of the SDSS was to survey large numbers of galaxies and quasars.  However, in the first five years of operation, SDSS-I has made remarkable contributions to our understanding of the Milky Way and its stellar populations \\citep[e.g.,][]{newberg:02,allendeprieto:06,belokurov:06,dejong:08,juric:08}. These successes have initiated the Galactic structure program SEGUE (Sloan Extension for Galactic Understanding and Exploration; B.\\ Yanny et al.\\ 2008, in preparation), one of the surveys being conducted in the ongoing three year extension of the survey (SDSS-II). When SDSS-II finishes, it will provide imaging data for approximately 10,000 square degrees of the northern sky.  SDSS measures the brightnesses of stars using a dedicated 2.5-m telescope \\citep{gunn:06} in five broadband filters $u$, $g$, $r$, $i$, and $z$, with average wavelengths of 3551\\AA, 4686\\AA, 6165\\AA, 7481\\AA, and 8931\\AA, respectively \\citep{fukugita:96,edr}.  The 95\\% detection repeatability limits are 22.0~mag, 22.2~mag, 22.2~mag, 21.3~mag, and 20.5~mag for point sources in $u$, $g$, $r$, $i$, and $z$, respectively.  The rms photometric precision is $0.02$~mag for sources not limited by photon statistics \\citep{ivezic:03}, and the photometric calibration is accurate to $\\sim2\\%$ in the $g$, $r$, $i$ bands, and $\\sim3\\%$ in $u$ and $z$ \\citep{ivezic:04}. The SDSS filters represent a new filter set for stellar observations, and therefore it is important to understand the properties of stars in this system.  Furthermore, future imaging surveys such as the Panoramic Survey Telescope \\& Rapid Response System \\citep[Pan-STARRS;][]{kaiser:02} and the Large Synoptic Survey Telescope \\citep[LSST;][]{stubbs:04} will use similar photometric bandpasses, providing even deeper data in $ugriz$ than SDSS over a larger fraction of the sky.  During the course of SDSS-I, about 15 globular clusters and several open clusters were observed.  Several more clusters were imaged in SDSS-II including M71.  These clusters together provide accurate calibration samples for stellar colors and magnitudes in the SDSS filters.  The SDSS images are processed using the standard SDSS photometric pipelines \\citep[{\\it Photo};][]{lupton:02}. {\\it Photo} pre-processes the raw images, determines the point spread function (PSF), detects objects, and measures their properties.  Photometric calibration is then carried out using observations of stars in the secondary patch transfer fields \\citep{tucker:06,davenport:07}.  In this paper, we simply refer to these calibrated magnitudes as {\\it Photo} magnitudes.  \\begin{figure*} \\epsscale{0.9} \\plotone{f1.ps} \\caption{CMDs of M3 from the SDSS photometric pipeline ({\\it Photo}; {\\it left}) and DAOPHOT reduction in this paper ({\\it right}). Stars within a $30\\arcmin$ radius from the cluster center are shown in the left panel, but the {\\it Photo} photometry is only available on the outskirts of the cluster.  In the right panel, RR Lyraes are scattered off the cluster horizontal branch. \\label{fig:photo}} \\end{figure*}  {\\it Photo} was originally designed to handle high Galactic latitude fields with relatively low densities of Galactic field stars (owing to the primarily extragalactic mission of SDSS-I); however, there are some concerns about its photometry derived in crowded fields \\citep{dr6}. In particular, stellar clusters present a challenge to {\\it Photo}. Firstly, {\\it Photo} slows down dramatically in the high density cluster cores, which are too crowded for {\\it Photo} to process, so it does not provide photometry for the most crowded regions of these scans. Figure~\\ref{fig:photo} compares a CMD for the globular cluster M3 obtained from {\\it Photo} photometry to that obtained from a DAOPHOT \\citep{stetson:87} reduction in this paper, which is specifically designed for crowded field photometry.  The {\\it Photo} photometry is only available on the outskirts of the cluster, and it provides a considerably less well-defined subgiant branch (SGB), red giant branch (RGB), and horizontal branch (HB).  Secondly, there is a concern that the photometry in the area surrounding clusters, and in low Galactic latitude fields, may also be affected by inaccurate modeling of the PSF if stars in crowded regions were selected as PSF stars by {\\it Photo}.  Photometric information in crowded fields can be extracted from the original SDSS imaging data.  For example, \\citet{smolcic:07} used the DoPHOT \\citep{schechter:93} photometry package to explore the structure of the Leo~I dwarf spheroidal galaxy (dSph).  Similarly, \\citet{coleman:07} used the DAOPHOT package to study the stellar distribution of the dSph Leo~II. In this paper, we employ the DAOPHOT/ALLFRAME \\citep{stetson:87,stetson:94} suite of programs to derive photometry for 17 globular clusters and 3 open clusters that have been observed with SDSS.  We derive photometry by running DAOPHOT for SDSS imaging frames where {\\it Photo} did not run. In addition, we reduce imaging data for fields farther away from the clusters, where the {\\it Photo} results are expected to be reliable, in order to set up photometric zero points for the DAOPHOT photometry.  We also compare DAOPHOT and {\\it Photo} results for the open cluster fields to verify the accuracy of the {\\it Photo} magnitudes in these semi-crowded fields.  An overview of the SDSS imaging survey and our sample clusters are presented in \\S~2.  In \\S~3 we describe the preparation of imaging data from the SDSS database.  In \\S~4 we describe the method of crowded field photometry using DAOPHOT/ALLFRAME, and evaluate the photometric accuracy.  In \\S~5 we present cluster photometry and fiducial sequences, and compare them with theoretical stellar isochrones and fiducial sequences in $u'g'r'i'z'$. ",
        "conclusions": ""
    },
    "0808/0808.3366_arXiv.txt": {
        "abstract": "The relevance of storage-ring electron-ion recombination experiments for astro\\-physics is outlined. In particular, the role of low-energy dielectronic-recombination resonances is discussed. A bibliographic compilation of electron-ion recombination measurements with cosmically abundant ions is provided. ",
        "introduction": "Heavy-ion storage rings equipped with electron coolers are an excellent experimental environment for electron-ion collision studies. Some recent studies of dielectronic recombination (DR) focussed on high-resolution spectroscopy of highly-charged ions. Highlights of this research are the measurement of the hyperfine induced decay rate of the $1s^2\\,2s\\,2p\\;^3P_0$ state in berylliumlike Ti$^{18+}$ \\cite{Schippers2007a} utilizing DR at the storage ring TSR of the Heidelberg Max-Planck-Institute for Nuclear Physics, the observation of the isotope shift in DR of three-electron Nd$^{57+}$ \\cite{Brandau2008a} using different isotopes of this ion at GSI's storage ring ESR and the observation of the hyperfine splitting of Sc$^{18+}$ low-energy  DR resonances at the TSR high-resolution electron target \\cite{Wolf2006c}. The latter experiment resulted in the determination of the Sc$^{18+}$($2s_{1/2} - 2p_{3/2}$) energy splitting with an uncertainty of only 4.6 ppm which is less than 1\\% of the few-body effects on radiative corrections \\cite{Lestinsky2008a}. Since these exciting developments have already been reviewed recently \\cite{Schippers2008a}, the present review focusses on the relevance of storage-ring electron-ion experiments for astrophysics.  Storage-ring experiments provide particularly valuable information on DR in low-temperature plasmas such as photoionized plasmas that occur, e.g., in active galactic nuclei (AGN) in the vicinity of super-massive black holes \\cite{Savin2007d}. In such plasmas highly charged ions exist at relatively low temperatures. For many ions, the DR rate coefficients, that determine the charge balance in these plasmas, depend sensitively on the low-energy DR resonance structure at relative electron-ion energies below $\\lesssim$ 3 eV. In the following, the influence of low-energy DR resonances on electron-ion recombination rate coefficients in low-density plasmas and recent efforts of building a recombination data base for astrophysical modeling of photoionized plasmas are briefly discussed. Finally, a compilation of experimental results for the astrophysically most abundant ions is presented. ",
        "conclusions": ""
    },
    "0808/0808.3150_arXiv.txt": {
        "abstract": "We show that the mass-segregation solution for the steady state distribution of stars around a massive black hole (MBH) has two branches: the known weak segregation solution \\citep{bah+77}, and a newly discovered strong segregation solution, presented here. The nature of the solution depends on the heavy-to-light stellar mass ratio $M_{H}/M_{L}$ and on the unbound population number ratio $N_{H}/N_{L}$, through the relaxational coupling parameter $\\Delta\\!=\\!4N_{H}M_{H}^{2}\\left/\\left[N_{L}M_{L}^{2}(3\\!+\\! M_{H}/M_{L})\\right]\\right.$. When the heavy stars are relatively common ($\\Delta\\!\\gg\\!1$), they scatter frequently on each other. This efficient self-coupling leads to weak mass segregation, where the stars form $n\\!\\propto\\! r^{-\\alpha_{M}}$ mass-dependent cusps near the MBH, with indices $\\alpha_{H}\\!=\\!7/4$ for the heavy stars and $3/2\\!<\\!\\alpha_{L}\\!<\\!7/4$ for the light stars (i.e. $\\max(\\alpha_{H}\\!-\\!\\alpha_{L})\\!\\simeq\\!1/4$). However, when the heavy stars are relatively rare ($\\Delta\\!\\ll\\!1$), they scatter mostly on light stars, sink to the center by dynamical friction and settle into a much steeper cusp with $2\\!\\lesssim\\!\\alpha_{H}\\!<\\!11/4$, while the light stars form a $3/2\\!<\\!\\alpha_{L}\\!<\\!7/4$ cusp, resulting in strong segregation (i.e. $\\max(\\alpha_{H}\\!-\\!\\alpha_{L})\\!\\simeq\\!1$). We show that the present-day mass function of evolved stellar populations (coeval or continuously star forming) with a universal initial mass function, separate into two distinct mass scales, $\\sim\\!1\\,\\Mo$ of main sequence and compact dwarfs, and $\\sim\\!10\\,\\Mo$ of stellar black holes (SBHs), and have $\\Delta\\!<\\!0.1$. We conclude that it is likely that many relaxed galactic nuclei are strongly segregated. We review indications of strong segregation in observations of the Galactic Center and in results of numeric simulations, and briefly list some possible implications of a very high central concentration of SBHs around a MBH. ",
        "introduction": "\\label{s:intro}  Early theoretical studies of the dynamics and distribution of stars around a MBH \\citep{pee72,bah+76,bah+77,you80} were triggered by the discovery of quasars \\citep[e.g.][]{mat+63,sch63} and the realization that many galactic nuclei may contain a central massive collapsed object \\citep{lyn69,wol+70}, %  {} as well as by the discovery of X-ray sources in globular clusters \\citep{gia+72}, which were then thought to be accreting MBHs (e.g. \\citealt{wyl70,bah+76}; see also \\citealt{mil+02b} for a recent reevaluation of this possibility). The main motivations for these studies were the prospect of detecting MBHs by the observed stellar density profile and by tidal disruption flares, and the possible role of tidal disruptions of stars in the growth of MBHs \\citep{lig+77,coh+78,sha+78,ree88}  The renewed interest in this problem is driven by observations of stars closely orbiting the Galactic MBH \\citep{eis+05,ghe+05} and the accumulating data on their distribution and dynamics \\citep{ale99,gen+00,gen+03a,sch+03,sch+07}, as well as by the prospects of detecting gravitational waves (GW) from extreme mass ratio inspiral sources by future GW detectors (EMRIs: compact remnants inspiraling into MBHs, see review by \\citealt{ama+07}; \\citealt{hop06}). EMRI rates and properties strongly depend on the stellar density and the stellar dynamical processes within $O(0.01\\,\\mathrm{pc})$ of the MBH, where inspiraling sources originate \\citep[e.g.][]{hop+05,hop+06a}.  Mass segregation occurs in dynamically relaxed systems. MBHs are naturally expected to lie in relaxed cores in scenarios where the MBH is formed by run-away mergers in the extreme central density following core collapse \\citep{ree84}, which occurs on timescales much longer then the relaxation time, $T_{R}$\\texttt{\\textbf{ }}\\citep[e.g. ][]{spi87,qui96b,fre+06b,fre+06c}. Likewise, the extreme mass ratio targets of the planned Laser Interferometer Space Antenna% \\footnote{See LISA mission website\\texttt{ http://lisa.nasa.gov}% } GW detector (LISA) are expected to originate in relaxed nuclei, since LISA design is sensitive to GW from MBHs with mass $\\lesssim\\!10^{7}\\, M_{\\odot}$. The observed correlation between the MBH mass $\\Mbh$ and the typical velocity dispersion of the spheroid of the host galaxy, $\\Mbh\\!\\propto\\!\\sigma^{\\beta}$, $4\\!\\lesssim\\!\\beta\\!\\lesssim5$ \\citep[the $\\Mbh/\\sigma$ relation, ][]{fer+00,geb+00} then implies that such nuclei are dynamically relaxed and very dense \\citep{ale07,mer+07}. This can be seen by assuming for simplicity $\\beta\\!=\\!4$ (a higher value only reinforce these conclusions), and noting that the MBH radius of influence $r_{h}\\!\\sim\\! G\\Mbh/\\sigma^{2}\\!\\propto\\!\\!\\Mbh^{1/2}$ encompasses a stellar mass of order $\\Mbh$, so that the number of stars there is $N_{h}\\!\\sim\\! M_{\\bullet}/\\Ms$, where $\\Ms$ is the typical stellar mass, and the mean stellar density is $\\bar{n}_{h}\\!\\sim\\! N_{h}/r_{h}^{3}\\!\\propto\\! M_{\\bullet}^{-1/2}$. The {}``$nv\\Sigma$'' rate estimate of strong gravitational collisions then implies that $T_{R}^{-1}(r_{h})\\!\\sim\\!\\bar{n}_{h}\\sigma(G\\Ms/\\sigma^{2})^{2}\\!\\propto\\!\\Mbh^{-5/4}$. A more rigorous estimate shows that for the Galactic MBH ($\\Mbh\\!\\simeq\\!4\\times10^{6}\\,\\Mo$, \\citealt{eis+05,ghe+05}), an archetype of LISA targets, $T_{R}\\!\\sim\\! O(1\\,\\mathrm{Gyr})\\!<t_{H}$ (the Hubble time) and $\\bar{n}_{h}\\!\\sim\\! O(10^{5}\\,\\mathrm{pc^{-3}})$. The density in the stellar cusp near the MBH is orders of magnitude higher still (see below). Since $T_{R}\\!\\propto\\!\\Mbh^{-5/4}$, isolated nuclei with $\\Mbh\\!\\lesssim\\!10^{7}\\,\\Mo$ are predicted to be relaxed.  A single mass stellar system around a MBH is expected to relax to a $r^{-\\alpha}$ cusp with $\\alpha\\!=\\!7/4$ \\citep{bah+76}. This results from the fact that the gravitational orbital energy gained by the system when stars are destroyed near the MBH is conserved as it is shared and carried outward by the remaining stars, $\\dot{E}(r)\\!\\sim\\! E(r)N(<r)/T_{R}\\propto r^{-1}r^{3-\\alpha}/r^{\\alpha-3/2}\\!=\\! r^{7/2-2\\alpha}\\!=\\mathrm{const}$ \\citep{bin+87}. When the system includes a spectrum of masses, $M_{L}\\le\\! M\\!\\le\\! M_{H}$, the approach toward equipartition by 2-body interactions decreases the specific kinetic energy of the high-mass stars, while that of the low-mass stars increases. As a result, the high-mass stars sink and concentrate in the center on the dynamical friction timescale $T_{\\mathrm{df}}\\T{\\sim}T_{R}\\left\\langle M\\right\\rangle /M_{H}$, while the low-mass stars float out \\citep{spi87}.  \\citet[hereafter BW77]{bah+77} approximated the mass segregation problem in the Fokker-Planck formalism, and solved for the steady state, angular momentum averaged stellar distribution functions (DFs) $f_{M}(E)$. They found that near the MBH, the DFs can be approximated by power-laws, $f_{M}(E)\\!\\propto\\! E^{p_{M}}$, and that the stellar current into the MBH, $Q_{M}(E)$, is very small, which leads to a specific relation between the stellar mass and the logarithmic slope of the DF, $p_{M}\\!=\\! M/4M_{H}$. In the Keplerian limit near the MBH, these DFs correspond to power-law density cusps, with $\\alpha_{M}\\!=\\!3/2+p_{M}$. The BW77 solution thus predicts a relatively small range of central concentrations, from $\\alpha_{H}\\!=\\!7/4$ for the heaviest stars in the populations, to $\\alpha_{L}\\!\\rightarrow\\!3/2$ for the lightest stars. We show below that the {}``zero-flow'' assumption breaks down when the massive stars are too rare to scatter efficiently against each other, and instead sink to the center by dynamical friction against the light stars. As a result, the BW77 relation between $M$ and $p_{M}$ no longer holds, and the range of central concentrations far exceeds that of the BW77 solution.  The Galactic Center (GC) provides to date the few available observations that directly bear on the question of mass segregation around a MBH. The over-abundance of X-ray transients in the central $\\sim\\!1$ pc of the GC was interpreted as evidence of a high central concentration of neutron stars and stellar black holes (SBHs) \\citep{mun+04}; The central decrease in the surface density of the low-mass horizontal branch {}``red clump'' giants was interpreted as evidence of the evacuation of long-lived light objects by mass segregation \\citep{lev06b,sch+07}. The dynamical upper limit on the distributed dark mass in the inner few mpc around the galactic MBH, $M_{\\mathrm{dm}}/\\Mbh\\!\\sim\\!\\mathrm{few\\times}0.01$ \\citep{mou+05,gil+08,ghe+08}, is still 10--100 times larger than predicted by approximate theoretical estimates \\citep{mor93,mir+00}, Fokker-Planck calculations \\citep{hop+06b}, or by the conservative drain limit \\citep{ale+04}, which places an upper bound on the maximal number of compact remnants that can avoid being thrown into the MBH by mutual 2-body scattering. These theoretical estimates are consistent with the upper bound derived from the observed limits on the diffuse X-ray in the central 0.7 pc ($\\sim2\\times10^{4}$) \\citep{dee+07}.  Detailed analysis of observations of the GC \\citep{lev06b,sch+07} suggests that the degree of segregation between the light and heavy stars is stronger than expected in the BW77 solution, although this interpretation of the data is not unique. Strong segregation in the GC is further supported by analytic \\citep{hop+06b} and Monte-Carlo \\citep[M. Freitag, priv. comm.]{fre+06} mass segregation results (Fig. \\ref{f:FPGCmodel}). These calculations do not take into account star formation and evolution, but instead assume for simplicity a non-evolving mass function that is based on a model of the GC's present day mass mass function (PMF) as an old, continuously star forming population with the {}``universal'' initial mass function (IMF) \\citep[table 2.1]{ale+99a,ale05}. In these theoretical models $\\alpha_{L}\\!\\simeq\\!1.5$, while $\\alpha_{H}\\!\\gtrsim\\!2$ at $r\\!\\sim\\!0.1$ pc. $N$-body simulations of stellar clusters with evolving stellar populations and a central intermediate mass black hole (IMBH) also provide hints of strong mass segregation \\citep[J. Makino, priv. comm.]{bau+04}. As we show below (Fig. \\ref{f:pLH_VR}), these results reflect the fact that the relative fraction of SBHs in old populations is below a critical threshold needed for them to scatter efficiently against each other and maintain the weak segregation solution. Instead, they sink to the center by dynamical friction and settle into the strong segregation solution.  This paper is organized as follows. The Fokker-Planck formulation of the mass segregation problem, the choices of boundary conditions and the various simplifying approximations are described in \\S \\ref{s:FPeq}. The physical meaning of the relaxational coupling parameter, $\\Delta$, is discussed in \\S \\ref{s:rlxparm}. The $\\Delta$ parameters of different stellar populations are explored in \\S \\ref{s:PMF}. A large grid of Fokker-Planck mass segregation calculations is presented and compared to simple analytic estimates in \\S \\ref{s:results}. The results and their implications are discussed and summarized in \\S \\ref{s:discuss}. ",
        "conclusions": "\\label{s:discuss}   \\subsection{Strong segregation and other relaxation processes}  Strong mass segregation, or mass segregation instability, shares some common features with the Spitzer, or equipartition, instability in a cluster \\citep{spi69}, where the heavy stars decouple from the light ones and evolve away from equipartition by forming a dense sub-system in the center with a much higher velocity dispersion, $\\sigma_{H}^{2}/\\sigma_{L}^{2}\\T{>}M_{L}/M_{H}$. However, these two instabilities are distinct effects. The Spitzer instability occurs only when the heavy stars are relatively common in the population, $N_{H}M_{H}^{5/2}/N_{L}M_{L}^{5/2}>\\beta$, where $\\beta\\!\\simeq0.16$ (for $M_{H}\\!\\gg\\! M_{L}$ and $N_{H}M_{H}\\!\\ll\\! N_{L}M_{L}$). In contrast, in the Keplerian potential near a MBH, the Jeans equation dictates that $\\sigma_{H}^{2}/\\sigma_{L}^{2}\\T{=}(5/2+p_{L})/(5/2+p_{H})\\T{\\sim}1$ \\citep{ale+01a}, and so equipartition is never achieved, irrespective of the heavy-to-light mass or number ratios. Strong mass segregation is an instability in the spatial distribution of the heavy stars, which occurs in the opposite limit to the Spitzer instability, when the heavy stars are relatively rare in the population, $\\Delta\\T{\\sim}N_{H}M_{H}/N_{L}M_{H}\\T{<}1$ (for $M_{H}\\T{\\gg}M_{L}$, Eq. \\ref{e:delta}).  \\citet{per+07} focused on the relaxation of light objects, stars, by heavy objects, massive perturbers (e.g. giant molecular clouds or clusters, with $M_{H}\\T{\\gg}M_{L}$), and expressed the efficiency of massive perturber-induced relaxation relative to star-star relaxation by the parameter $\\mu_{2}\\T{=}N_{H}M_{H}^{2}/N_{L}M_{L}^{2}$ ($\\sim\\!(D_{HLs}+D_{HLf})/(D_{LLs}+D_{LLf})\\T{\\simeq}3N_{H}M_{H}^{2}/4N_{L}M_{L}^{2}$). When $\\mu_{2}\\T{\\gg}1$, massive perturbers dominate stellar relaxation. Here we focus on the dynamics of the heavy objects, the SBHs, and so the relaxational coupling parameter $\\Delta$ is defined to also take into account the interactions between the heavy masses. The parameter $\\mu_{2}$ addresses the question ``Which mass component dominates the relaxation of the light stars?'', while $\\Delta$ addresses the question ``Which mass component dominates the dynamics of the heavy stars?''. The two are related by $\\Delta\\T{=}\\mu_{2}4/(3+M_{H}/M_{L})$.  Our approximate treatment of the mass segregation process neglects the relaxation of angular momentum to near radial ({}``loss-cone'') orbits. This, rather than diffusion in energy, is the primary channel for stellar destruction by the MBH. A full treatment of the mass segregation problem in ($E,J$) phase space \\citep[e.g.][]{coh+78} is beyond the scope of this work. Here we follow BW77, who treated the problem approximately in $E$ only, and who further showed that the neglect of an effective loss-cone term (Eq. \\ref{e:FPeq}) did not much change the shape of the DFs. We confirmed that this conclusion also holds for our mass segregation models (Fig. \\ref{f:Delta_flow}).   \\subsection{Possible implications of strong mass segregation}  Strong mass segregation occurs in stellar systems with a relatively\\emph{ lower} fraction of SBHs, that reach a \\emph{higher} central concentration of SBHs very close to the MBH, compared to systems with a higher fraction of SBHs that undergo weak segregation. To compare two such systems, and to determine which has more SBHs enclosed inside some given volume around the MBH, it is necessary to specify the comparison procedure (e.g. assuming an equal total stellar mass; or the same MBH mass and the $\\Mbh/\\sigma$ relation; or an equal number of SBHs within $r_{h}$). The choice depends on the question of interest. Here we do not address such quantitative issues, but limit ourselves to briefly listing some processes that are expected to be affected by the degree of segregation.  \\emph{Accelerated relaxation. }The degree of mass segregation affects both the non-coherent 2-body relaxation timescale, $ $which scales as $1\\left/\\int M^{2}(\\mathrm{d}N/\\mathrm{d}M)\\mathrm{d}M\\right.$ (see \\S \\ref{s:intro}), and the resonant relaxation timescale, which scales as as $\\int M(\\mathrm{d}N/\\mathrm{d}M)\\mathrm{d}M\\left/\\int M^{2}(\\mathrm{d}N/\\mathrm{d}M)\\mathrm{d}M\\right.$ \\citep{rau+96}. In particular, the stronger the mass segregation, the shorter is the resonant relaxation timescale, which does not depend on the number of stars, but only on their typical mass. Efficient resonant relaxation near the MBH may affect stellar orbits and accretion disk dynamics there (see below).  \\emph{GW event rates.} The GW EMRI rate is determined by the number of potential GW sources within the critical radius $r_{\\mathrm{crit}}$, which demarcates the boundary between compact object that inspiral into the MBH those that plunge (infall) into it. The critical radius is a function of the relaxation time, and to good approximation the EMRI rate is $\\Gamma\\T{\\sim}N[r_{crit}(T_{R})]/T_{R}$ \\citep{hop+05}. Strong segregation will affect the EMRI rate both by modifying the 2-body relaxation time and by affecting the number of stars enclosed inside $r_{\\mathrm{crit}}$, as well as by decreasing the resonant relaxation timescale. Similarly, the rates of detectable GW bursts from fly-bys near the Galactic MBH strongly depend on the number of SBHs near it \\citep{hop+07}. Strong segregation may also affect GW emission from close SBH--SBH interactions in a very dense cusp \\citep{ole+08b}.  \\emph{SBH--star interactions.} A higher central concentration of SBHs affects the probability of SBH--star interactions, which can lead to the randomization of stellar orbits, the heating of a stellar disk \\citep{per+08}, the 3-body exchange capture of massive young stars near the MBH \\citep{ale+04}, or the ejection of hyper-velocity stars \\citep{ole+08}.  \\emph{SBH--accretion disk interactions.} A higher central concentration of SBHs within $10^{1-3}$ gravitational radii of the MBH could exert coherent torques on the accretion disk, warp it and possibly affect its hydrodynamics (Bregman \\& Alexander, 2008, in prep.). The SBHs may shock the disk as they cross it, experience drag by it, and be carried by it to the MBH \\citep[e.g.][]{art+93,nay+04,sub+05}.  \\emph{Enhanced gravitational lensing. }SBHs projected near the Einstein angle of the MBH can strongly modify the gravitational lensing properties of the MBH,\\emph{ }in a way similar to the effect of a planet orbiting a Galactic star that is lensing a background source \\citep{ale+01c,cha+01a}.   \\subsection{Summary}  We show that the steady state solution of a relaxed multi-mass stellar system around a MBH has two branches: the known weak (Bahcall-Wolf) mass segregation solution, where the difference in the degree of central concentration of the light and heavy stars is relatively small, and a newly discovered strong segregation solution, where the difference is much larger. The nature of the solution is determined by the global properties of the system (the mass ratio between the heavy and light stars, $M_{H}/M_{L}$, and their number ratio far from the MBH, $N_{H}/N_{L}$) through the relaxational coupling parameter, $\\Delta\\!=\\!4N_{H}M_{H}^{2}\\left/\\left[N_{L}M_{L}^{2}(3\\!+\\! M_{H}/M_{L})\\right]\\right.$. Strong mass segregation occurs when the heavy stars are relatively rare in the population $(\\Delta\\T{\\ll}1)$, and sink to the center by dynamical friction. Weak mass segregation occurs when the heavy stars are common in the population ($\\Delta\\T{\\gg}1$) and settle to the single mass stellar cusp solution. We show that relaxed old coeval or continuously star-forming populations with a universal IMF typically have $\\Delta\\T{<}0.1$, and thus settle to the strong mass segregation solution around a MBH."
    },
    "0808/0808.1539_arXiv.txt": {
        "abstract": "{Little is known about the structure of the interstellar medium and the nature of individual clouds in galaxies at intermediate redshifts. The gravitational lens toward PKS\\,1830--211 offers the unique possibility to study this interstellar gas with high sensitivity and angular resolution in a molecular cloud that existed half a Hubble time ago.} {This multi-line study aims at a better definition of the physical properties of a significantly redshifted cloud.} {Using the Green Bank Telescope (GBT), we searched for ammonia (NH$_3$) toward PKS\\,1830--211.} {We have detected all ten observed metastable $\\lambda$$\\sim$2\\,cm ammonia lines. The ($J$,$K$) = (1,1) to (10,10) transitions, up to $\\sim$1030\\,K above the ground state, were measured in absorption against the radio continuum of the lensed background source. The ammonia absorption appears to be optically thin, with absolute peak flux densities up to 2.5\\% of the total continuum flux density. Measured intensities are consistent with a kinetic temperature of $\\sim$80\\,K for 80--90\\% of the ammonia column. The remaining 10--20\\% are warmer, with at least some of this gas reaching kinetic temperatures of $\\ga$600\\,K. Toward the south-western continuum source, the column density is $N$(NH$_3$) $\\sim$ (5--10)$\\times$10$^{14}$\\,cm$^{-2}$, which implies a fractional abundance of $X$(NH$_3$) $\\sim$ (1.5--3.0)$\\times$10$^{-8}$. Similarities with a hot NH$_3$ absorption component toward the Sgr~B2 region close to our Galactic center, observed up to the (18,18) line, suggest that the Sgr~B2 component also consists of warm diffuse low-density gas. The warm absorption features from PKS\\,1830--211 are unique in the sense that they originate in a spiral arm.} {} ",
        "introduction": "Studies of interstellar emission and absorption lines are complementary. While emission commonly traces an extended gas reservoir, which provides information on global properties such as mass and kinematics, absorption tends to trace a spatially far smaller region confined by the angular extent of a continuum background source. Absorption lines allow us to derive optical depths directly from line to continuum flux density ratios; the strength of the absorption depends on the flux of the background source and is decoupled from the distance to the absorber. Thus, for objects at significant redshifts, studies of interstellar absorption lines have the potential to provide unique information on otherwise inaccessible regions, combining extremely high sensitivity with outstanding spatial resolution (e.g., Chengalur et al. \\cite{chengalur99}; Wiklind \\& Combes \\cite{wc99}).  In spite of extended surveys, multi-line molecular absorption systems have so far only been discovered toward five distant targets (Wiklind \\& Combes \\cite{wc94}, \\cite{wc95}, \\cite{wc96a}; \\cite{wc96b}; Kanekar et al. \\cite{kanekar05}), all of which are located at intermediate redshifts (0.25$\\leq$$z$$\\leq$0.89). The two most notable absorbers, those toward B0218+357 and PKS\\,1830--211, share a number of common properties. They have similar redshifts ($z$=0.68 and 0.89, respectively), show the highest line-of-sight column densities, and originate in gravitational lenses. Toward both systems the respective background sources, probably a BL Lac object and a blazar, respectively (e.g. Kemball et al. \\cite{kem01}; De Rosa et al. \\cite{ros05}), are lensed into three dominant features, a north-eastern and a south-western hotspot and an Einstein ring. Optical or near infrared and radio separations of the hotspots differ due to optical obscuration (Courbin et al. \\cite{courbin98}; York et al. \\cite{york05}). The lensing galaxies themselves are spirals viewed almost face-on (Courbin et al. \\cite{courbin02}; Winn et al. \\cite{winn02}; York et al. \\cite{york05}). The bulk of the absorption arises in both cases from in front of the south-western of the two main radio continuum images, at radial distances of $\\sim$2 and $\\sim$4\\,kpc from the center of the respective lens (e.g., Frye et al. \\cite{frye97}; Swift et al. \\cite{swift01}; Meylan et al. \\cite{mey05}; York et al. \\cite{york05}). In both cases, this absorption originates in the south-western spiral arm of the parent galaxy.  There are also significant differences between B0218+357 and PKS\\,1830--211: one is image separation (334 versus 970\\,mas; e.g., Jin et al. \\cite{jin03}; Wucknitz et al. \\cite{wuck04}); another is time delay (10 versus 25 days; e.g., Lovell et al. \\cite{lov98}; Biggs et al. \\cite{biggs01}). With respect to the detection of molecular absorption lines, the most striking differences are the continuum levels and the complexity of the absorption systems. B0218+357 hosts two molecular components separated by $\\sim$13\\,km\\,s$^{-1}$ observed toward the south-western continuum image (Menten \\& Reid \\cite{menten96}; Jethava et al. \\cite{jethava07}; Muller et al. \\cite{muller07}). Toward PKS\\,1830--211, weak molecular absorption is also measured toward the north-eastern radio continuum hotspot, displaced by \\hbox{--147}\\,km\\,s$^{-1}$ from the dominant south-western absorption component (Wiklind \\& Combes \\cite{wc98}). Furthermore, there is H{\\sc i} absorption at $z$=0.19 (Lovell et al. \\cite{lovell96}).  PKS\\,1830--211 is one of the most prominent compact radio sources in the sky and therefore an ideal target for absorption line studies at radio wavelengths. Molecular species that have been detected include CS, HCN, HCO$^+$, HNC and N$_2$H$^+$ (Wiklind \\& Combes \\cite{wc96b}), CO (Gerin et al. \\cite{gerin97}), OH (Chengalur et al. \\cite{chengalur99}) and C$_2$H, H$_2$CO, C$_3$H$_2$, and HC$_3$N (Menten et al. \\cite{menten99}). The $\\lambda$21\\,cm line of H{\\sc i} was also observed at $z$=0.89 (Chengalur et al. \\cite{chengalur99}). Muller et al. (\\cite{muller06}) determined CNO and sulfur isotope ratios.  PKS\\,1830--211 provides a unique view of a molecular cloud at more than half a Hubble time in the past ($\\Lambda$-cosmology with $H_{0}$=73\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm m}$ = 0.28 and $\\Omega_{\\Lambda}$ = 0.72; Spergel et al. \\cite{sper07}; 1\\,mas corresponds to $\\sim$7.5\\,pc). With a luminosity distance of $\\sim$5.5\\,Gpc it is the farthest molecular cloud known to date that can be studied in detail. The nature of this object is, however, poorly understood. In an attempt to constrain the physical properties of the target and motivated by the recent detection of NH$_3$ toward B0218+357 (Henkel et al. \\cite{henkel05}), we searched for NH$_3$ absorption to evaluate the still poorly constrained physical parameters of the absorbing molecular gas. ",
        "conclusions": "Our ammonia observations of the gravitational lens system PKS\\,1830--211 at a redshift of $z$$\\sim$0.9 reveal the following main results:  \\begin{itemize}  \\item Ammonia (NH$_3$) is detected in absorption in its ten lowest metatstable inversion lines. The ($J$,$K$) = (10,10) transition lies $\\sim$1030\\,K above the ground state.  \\item The NH$_3$ absorption lines are optically thin and reach, in the strongest line, only 2.5\\% of the continuum level. No ammonia absorption is detected toward the system displaced by \\hbox{--147}\\,km\\,s$^{-1}$ relative to the main component.  \\item The absorbing gas is warm: 80-90\\% of the absorption arises from an environment with $T_{\\rm kin}$ $\\sim$ 80\\,K. The remaining gas is with $>$100\\,K considerably warmer; for some gas, $T_{\\rm kin}$$\\ga$600\\,K. The total column density is approximately $N$(NH$_3$) = (5--10)$\\times$10$^{14}$\\,cm$^{-2}$.  \\item Compared with H$_2$ column densities derived from CO at mm-wavelengths, the fractional abundance is $X$(NH$_3$) $\\sim$ 1.5--3.0$\\times$10$^{-8}$. This is a low value when compared with abundances in Galactic cloud cores. The low value may be caused by a low metallicity implying small abundances of nitrogen and dust (NH$_3$ is easily destroyed by UV-photons), by diffuse gas that also shows low NH$_3$ abundances along Galactic lines-of-sight or it may reflect differences between the continuum background morphology at cm- and mm-wavelengths.  \\item If the H{\\sc i} spin temperature is equal to the kinetic temperature of the bulk of the NH$_3$ absorbing gas (80\\,K), the total column must be mostly molecular. If, however, most of the H{\\sc i} is associated with the hot ammonia component, the H{\\sc i} and molecular column densities should be similar.  \\item In the nearby Universe, only the envelope of Sgr~B2 has a cloud component with ammonia properties similar to those seen toward PKS\\,1830--211. This cloud contains gas from either hot cores or from low-density gas forming an envelope; the analogy with PKS1830--211 suggests the latter. The column toward PKS\\,1830--211 is unique in the sense that it arises from a spiral arm and not from the central region of a gas-rich galaxy.  \\end{itemize}  Following Flambaum \\& Kozlov (\\cite{flambaum07}), a detailed comparison of NH$_3$ inversion lines with rotational spectra from linear molecules can provide limits on the space-time variation of the ratio of the proton to electron mass. Making use of the ammonia profiles presented here, such a comparison will be carried out in a forthcoming paper."
    },
    "0808/0808.0223_arXiv.txt": {
        "abstract": "s{Magnetowave induced plasma wakefield acceleration (MPWA) in a relativistic astrophysical outflow has been proposed as a viable mechanism for the acceleration of cosmic particles to ultra high energies. Here we present simulation results that demonstrate the viability of this mechanism. We invoke the high frequency and high speed whistler mode for the driving pulse. The plasma wakefield so induced validates precisely the theoretical prediction. This mechanism is shown capable of accelerating charged particles to ZeV energies in Active Galactic Nuclei (AGN). } ",
        "introduction": "The origin of ultra high energy cosmic rays (UHECR) is an intriguing question in astrophysics. Theories categorized as either ``top-down\" or ``bottom-up\" scenarios are proposed to answer this question. Each scenario faces its own theoretical and observational challenges \\cite{Olinto}. Since the observations from HiRes \\cite{cosmic:HiRes} and Auger \\cite{cosmic:Auger} confirm the Greisen-Zatepin-Kuz'min (GZK) suppression of the cosmic ray flux \\cite{gzk}, the need for top-down exotic models is reduced. Hence the challenge to find a viable ``bottom up\" mechanism for accelerating UHECR becomes more acute.  Shocks, unipolar inductors and magnetic flares are the three most potent, observed, ``conventional\" accelerators that can be extended to account for $\\sim{\\rm ZeV} (=10^{21}{\\rm eV})$ energy cosmic rays \\cite{Blandford:1999}. Radio jet termination shocks and gamma ray bursts (GRB) have been invoked as sites for the shock acceleration, while dormant galactic center black holes and magnetars have been proposed as sites for the unipolar inductor acceleration and the flare acceleration, respectively. Each of these models, however, presents problems \\cite{Blandford:1999}. Evidently, novel acceleration mechanisms that can avoid the difficulties faced by these conventional models should not be overlooked.  Plasma wakefield accelerators \\cite{LWFA:Tajima,PWFA:Chen85} are known to possess two salient features: (1) The energy gain per unit distance does not depend (inversely) on the particle's instantaneous energy; (2) The acceleration is linear. These features are essential for a good acceleration efficiency. Although high-intensity, ultra-short photon or particle beam pulses that excite the laboratory plasma wakefields are not available in the astrophysical setting, large amplitude plasma wakefields can instead be excited by the astrophysically abundant plasma ``magnetowaves\" \\cite{plasma:Chen02}. Protons can be accelerated beyond ZeV energy by riding on such wakefields. This attractive concept has never been validated through self-consistent computer simulations. In this presentation, we report our simulation that confirm this concept \\cite{Chang:2007um}. We also discuss this acceleration mechanism in AGN. ",
        "conclusions": "Through PIC simulations, we have confirmed the concept of plasma wakefield excited by a magnetowave in the magnetized plasma. We have demonstrated how such a wakefield may accelerate particles to ZeV energies in AGN. As a first step, we investigated MPWA in the parallel-field configuration. Since both poloidal and toroidal field components are inevitable in astro-jets, we will further investigate plasma wakefield excitation and acceleration under the cross-field configuration. Besides, we have limited our discussions in the linear regime $a_0\\ll 1$, which is applicable to AGN. However, in other astrophysical settings, such as GRB, the magnetowave could be much stronger such that $a_0\\gg 1$. We will investigate plasma wakefield generations in this regime and therefore explore the MWPA production of UHECR in other powerful astrophysical sites."
    },
    "0808/0808.0153_arXiv.txt": {
        "abstract": "We present deep HST ACS/WFC photometry of the dwarf irregular galaxy NGC 1569, one of the closest and strongest nearby starburst galaxies. These data allow us, for the first time, to unequivocally detect the tip of the red giant branch and thereby determine the distance to NGC 1569.  We find that this galaxy is 3.36$\\pm$0.20 Mpc away, considerably farther away than the typically assumed distance of 2.2$\\pm$0.6 Mpc. Previously thought to be an isolated galaxy due to its shorter distance, our new distance firmly establishes NGC 1569 as a member of the IC 342 group of galaxies. The higher density environment may help explain the starburst nature of NGC 1569, since starbursts are often triggered by galaxy interactions. On the other hand, the longer distance implies that NGC 1569 is an even more extreme starburst galaxy than previously believed. Previous estimates of the rate of star formation for stars younger than $\\la$ 1 Gyr become stronger by more than a factor of 2.  Stars older than this were not constrained by previous studies.  The dynamical masses of NGC 1569's three super star clusters, which are already known as some of the most massive ever discovered, increase by $\\sim$53\\% to 6-7$\\times10^5$\\msun. ",
        "introduction": "Massive starbursts drive the evolution of galaxies at high redshift; they provide chemical enrichment and thermal and mechanical heating of both the interstellar and intergalactic medium.  In the past decade, a large population of star-forming galaxies has been discovered at high redshift, highlighting the importance of these galaxies on a cosmological scale.  However, starburst galaxies can only be studied in detail in the nearby Universe where they are much rarer. The dwarf irregular galaxy NGC 1569 is one of the closest examples of a true starburst.  Star formation in NGC 1569 has been studied extensively with the Hubble Space Telescope (HST; e.g., Greggio et al.~1998, hereafter G98, Aloisi et al.~2001, Angeretti et al.~2005, hereafter A05), with results showing that its star formation rates (SFRs) are 2-3 times higher than in other strong starbursts (e.g.~NGC 1705) and 2-3 orders of magnitude higher than in Local Group irregulars and the solar neighborhood, if the SFR per unit area is considered. NGC 1569 is also home to three of the most massive super star clusters (SSCs) ever discovered.\\looseness=-2  Although it lies on the sky in the same direction as the IC 342 group of galaxies (see Fig.~2 in Karachentsev et al.~2003), the typically assumed distance of 2.2$\\pm$0.6 Mpc (Israel 1988) is based on the luminosity of the brightest resolved stars (Ables 1971) and places NGC 1569 well in front of IC 342 ($D = 3.28 \\pm 0.27$ Mpc; Saha, Claver, \\& Hoessel 2002). NGC 1569 has therefore generally been viewed as an isolated starburst galaxy just beyond the outskirts of the Local Group. The isolated environment makes it more complicated to understand the trigger of the starburst, since starbursts are often associated with galaxy interactions and are not generally believed to be internally driven. However, NGC 1569's distance is relatively uncertain, and a large range of possible distances exists in the literature.  For example, using HST/WFPC2 photometry, Makarova \\& Karachentsev (2003) attempted to measure the brightness of stars at the tip of the red giant branch (TRGB) in NGC 1569 and found, due to the limited number of stars available at their detection limit, two possible solutions; a short distance of $1.95\\pm0.2$ Mpc or a long distance of $2.8\\pm0.2$ Mpc, with no preference for one over the other.  We note also that distance determinations based on the brightest resolved stars in dwarf galaxies are prone to considerable uncertainty due to stochastic effects (e.g.~Greggio 1986), and that, at the distance of NGC 1569, ground-based observations suffer from the possibility of misinterpreting star clusters or pairs of stars as individual stars.  In an effort to better determine its star formation history (SFH), we have used the HST ACS/WFC to obtain deep $V$- and $I$-band photometry of NGC 1569. These data also provide for the first time an accurate distance based on unequivocal identification of the TRGB, which we present and discuss in this Letter. ",
        "conclusions": ""
    },
    "0808/0808.2156_arXiv.txt": {
        "abstract": "A new method for determining the stellar rotation period is proposed here, based on the detection of starspots during transits of an extra-solar planet orbiting its host star. As the planet eclipses the star, it may pass in front of a starspot which will then make itself known through small flux  variations in the transit light curve. If we are lucky enough to catch the same spot on two consecutive transits, it is possible to estimate the stellar rotational period. This method is successfully tested on transit simulations on the Sun yielding the correct value for the solar period. By detecting two starspots on more than one transit of HD 209458 observed by the Hubble Space Telescope, it was possible to estimate a period of either 9.9 or 11.4 days for the star, depending on which spot is responsible for the signature in the light curve a few transits later. Comparison with period estimates of HD209458 reported in the literature indicates that 11.4 days is the most likely stellar rotation period. ",
        "introduction": "Four centuries have elapsed since the first determination of the rotational period of a star, namely the Sun, by observation of the apparent movement of starspots on its surface. For star other than the Sun, the determination of the rotational period is made basically by two methods: either from the rotational broadening of spectral lines or by the periodic modulation of the stellar flux due to the dark and bright features on the stellar surface which rotate with it. The latter method has the advantage of determining directly the period of a star, without the sin$i$ spectroscopic uncertainty. Moreover, it can also be applied to stars with long rotational period, which cannot be determined from Doppler broadening of spectral lines.  Here I would like to propose a new way of estimating the stellar rotation by using planetary transits. If by any chance, during a transit, the planet passes in front of a starspot, as in Silva (2003), and in a consecutive transit the configuration is such that this same spot is again occulted, then it is possible to estimate the stellar rotation as it was done for the Sun four centuries ago.  Previous authors have used a similar method to estimate, among other things, the rotational period of the components of eclipsing binaries. Eaton and Hall (1979) explained successfully the light curve variations of RS CVn type stars as due to starspots. Besides the rotation period and its variation, durations of activity cycles and surface area coverage of spots have also been obtained for the prototypes of active binaries RS CVn (Rodono et al., 1995) and RT Lac (Lanza et al. 2002). These are very active stars, nevertheless it will be possible to study moderately active stars such as our Sun with space missions such as CoRoT, MOST, and Kepler. In anticipation to these satellites, the rotational modulation of the Sun as a star has been modeled using the VIRGO/SoHO data (Lanza et al. 2003, 2004).  The method described here, however, has a much better spatial resolution than that of eclipsing binaries, because of the smaller planet diameter in comparison to that of the stars. Moreover, planetary transits such as the ones used here will be quite commonly detected by the CoRoT mission already in operation. Section \\ref{transit} presents the method, which is tested on the Sun with a fiducially transiting planet (Section \\ref{solar}) and then applied to HD209458 (Section \\ref{hd}). The last section discusses the results of this method and presents the conclusions. ",
        "conclusions": "This work proposes to estimate the rotational period of a star by following the apparent shift in longitude position of its surface spots, similar to what was done by Galileo and his contemporaries four centuries ago for the Sun. Tracking of the spots is done by identifying ``bumps\" in the light curves of successive planetary transits.  This method was successfully tested for the Sun yielding the correct value for the solar period from simulated transits taken only three days apart even though the solar period is about 27 days. Moreover, by modeling of different sunspot groups it was possible to verify the solar differential rotation. Supposing that the small variations detected in the light curve during planetary transits is due to occultation of starspots, the model was also applied to HD 209458 using HST observations obtained by Brown et al. (2001) in April and May of 2000. The star was modeled as having two spots during the four transits. The ``bumps\" detected on two transits separated by three orbital periods yield a stellar rotation period of 9.9 or 11.4 days depending on which of the spots detected on April 25th is considered to cause the intensity variation on the May 5th transit light curve.  Several observations of $v \\sin i$ of HD 209458 are found in the literature. Assuming that the inclination angle is 86.68$^\\circ$ and the stellar radius 1.148 R$_\\odot$, these observations yield stellar periods of 14.4 $\\pm$ 2.1 days (Mazeh et al. 2000) and 15 $\\pm$ 6 days (Queloz et al. 2000). More recently, shorter periods have been found, 12.3 $\\pm$ 0.5 days (Winn et al. 2005) and 12 days (Fisher \\& Valenti 2005).  The periods obtained here, 9.9 and 11.4 days, are a little shorter than those found in the literature. It seems that the 11.4 days period is the true one for HD 209458, even though it is still a little shorter than the periods listed in the literature. However, HD 209458 probably presents differential rotation, similar to the Sun. In this case, the period obtained by other authors, which was based on line broadening observations, is actually an average of the periods of the whole stellar disk, whereas the period determined here represents the rotational velocity at that specific latitude which is close to the equator. As the planet size spans about 20$^o$ in stellar latitude, we are probing latitudes from -22$^o$ to -38$^o$. Since the work of Winn et al. (2005) was based on the observation of the Rossiter-McLaughlin effect during transits, therefore obtained at the same latitudes we sample, this is the result which should better agree with the one presented here.  Thus far, over 50 transiting planets have been detected (The Extrasolar Planets Encyclopaedia - http://exoplanet.eu) and many more are expected especially in the next months thanks to observations by the CoRoT satellite. The method proposed here can easily be applied to the planetary transits of newly discovered planets and checked against periodic modulation of the stellar flux outside transits themselves."
    },
    "0808/0808.3293_arXiv.txt": {
        "abstract": "Tidal tails of star clusters are not homogeneous but show well defined clumps in observations as well as in numerical simulations. Recently an epicyclic theory for the formation of these clumps was presented. A quantitative analysis was still missing. We present a quantitative derivation of the angular momentum and energy distribution of escaping stars from a star cluster in the tidal field of the Milky Way and derive the connection to the position and width of the clumps. For the numerical realization we use star-by-star $N$-body simulations. We find a very good agreement of theory and models. We show that the radial offset of the tidal arms scales with the tidal radius, which is a function of cluster mass and the rotation curve at the cluster orbit. The mean radial offset is 2.77 times the tidal radius in the outer disc. Near the Galactic centre the circumstances are more complicated, but to lowest order the theory still applies. We have also measured the Jacobi energy distribution of bound stars and showed that there is a large fraction of stars (about 35\\%) above the critical Jacobi energy at all times, which can potentially leave the cluster. This is a hint that the mass loss is dominated by a self-regulating process of increasing Jacobi energy due to the weakening of the potential well of the star cluster, which is induced by the mass loss itself. ",
        "introduction": "\\label{sec-intro}  Recently well-defined clumps were observed in the tidal tails of globular clusters by \\citet{Le00} for NGC 6254 and Pal 12 and by \\citet{Od01,Od03} for Pal 5. External perturbations like crossing of the galactic disc, peri-centre passage or near-encounters with other globular clusters were discussed as the source of the clumps \\citep{Ca05}. In \\citet{Ca05} the formation of these clumps in tidal tails of star clusters on eccentric orbits were confirmed by numerical studies. But the clumps occur also in the tidal tails of star clusters on circular orbits with no external push. Recently, \\citet{Ku08} presented a theoretical explanation for the clump formation in a constant tidal field. It is essentially due to the epicyclic motion of the stars lost by the star cluster.  We present a quantitative analysis of the tidal tail structure for star clusters moving on a circular orbit in the galactic disc. The analysis is based on numerical simulations with realistic particle numbers including an initial mass function (IMF) and stellar evolution. We compare the results for star clusters at the solar circle and near the Galactic centre. Since the epicycle theory is a perturbation theory with respect to a circular orbit with constant tidal field, we cannot apply it for predictions of clump distances to the eccentric orbits of the observed globular clusters. On the other hand the observation of tidal tail clumps of open clusters in the galactic disc are hampered by the overwhelming number of field stars with similar properties. For an identification the contrast in density and velocity with respect to the field stars may be too small. Additionally the tidal tails may be destroyed quickly by the same gravitational scattering process, which is also responsible for the dynamical heating of the stellar disc. We discuss the observability further in Sect.~\\ref{sec-sum}.  In Section \\ref{sec-dyn} we present the epicyclic theory for the stars in the tidal tails and the connection to the mass loss and the orbit of the star cluster. In Section \\ref{sec-num} we present the numerical codes used and the properties of the star clusters. Section \\ref{sec-res} contains the quantitative comparison of the numerical results with the theoretical predictions. In Section \\ref{sec-sum} we summarize our results. ",
        "conclusions": "\\label{sec-sum}  We presented a quantitative derivation of the angular momentum and energy distribution of escaping stars from a star cluster in the tidal field of the Milky Way. Despite the motion on a circular orbit, the tidal tails are clumpy due to the epicyclic motion of the stars. We compared the derived distances and widths of the clumps with numerical simulations using star-by-star simulations. For star clusters at the solar circle we included an IMF and mass loss due to stellar evolution in the calculations. The same equations were applied to a star cluster very close to the Galactic centre, where the tidal forces are very strong.  We find a very good agreement of theory and models concerning the tidal tail structure. The positions of the clumps are determined by the angular momentum offset of the stars, which lead to a radial offset of the epicenters with respect to the cluster orbit. The investigation of \\citet{Ku08} is a special case of our investigations but for a Kepler potential. We find that the radial offset of the tidal arms is proportional to the tidal radius. However near the Galactic centre the factor of proportionality is considerably smaller.  The tidal arm structure at large galactocentric radii is symmetric, whereas the asymmetry near the galactic centre is considerable. This can be reproduced only partly by taking into account the correction of the epicyclic frequency at the epicentre radii.  We have also measured the Jacobi energy distribution of bound stars and showed that there are 35\\% of stars above the critical Jacobi energy independent of the evolutionary state of the cluster. These stars can potentially leave the cluster. This is a hint, that mass loss is dominated by a self-regulating process of increasing Jacobi energy due to the diminishing gravitational potential of the star cluster induced by the mass loss itself.  Finally we consider the observability of the predicted clump properties in the tidal tails of star clusters. The identification of tidal tail stars of open clusters on a circular orbit is strongly hampered by the large number of nearby field stars with similar properties. But with differential methods using high quality data for distances and velocities it may be possible to identify the most prominent first clumps. We have shown that the maximum density in the first clump does not decrease with time until dissolution of the cluster. The first clumps are formed by escaping stars with a time delay determined by the epicyclic period $T_\\kappa\\approx 150$\\,Myr. The tidal tail structure will probably survive the gravitational scattering process, which is also responsible for the galactic disc heating \\citep{Wie77}. On a timescale of one epicyclic period, we expect only small perturbations of the tidal clump position and velocity but no destruction. The density maximum of the first clump is of the order of a few percent of the field density of the galactic disc. The velocity imprint by the epicyclic motion is of the order of 2\\,km/s. An overdensity with these properties may be difficult to observe for clusters on exact circular orbits, if there is no additional separating property. For young star clusters the high fraction of early type stars can serve for such a discrimination. On the other hand most star clusters are identified by the systematic peculiar motion with respect to the field stars. If the positions and velocities of the tidal clump stars are properly predicted, they may be observable as moving groups.  For an application of the tidal clump theory to the eccentric orbits of globular clusters a perturbation theory with respect to the 'free falling' comoving coordinate system would be necessary. In this case the Jacobi energy is no longer a constant of motion. This will be a matter of future investigations. Some numerical test runs have shown that tidal tail clumps are formed also on highly eccentric orbits without an additional external perturbation, but the geometry and density is modulated along the orbital position."
    },
    "0808/0808.1069_arXiv.txt": {
        "abstract": "{The remarkable pre-main-sequence object \\mbox{V718 Per} (HMW 15, H187) in the young cluster IC 348 periodically undergoes long-lasting eclipses caused by variable amounts of circumstellar dust in the line-of-sight to the star. It has been speculated that the star is a close binary and similar to another unusual eclipsing object, KH 15D.} {We have submitted \\mbox{V718 Per} to a detailed photometric and spectroscopic study to investigate how regular the recurrent eclipses are, to find out more about the properties of the stellar object and the occulting circumstellar material, and to look for signatures of a possible binary component.} {\\mbox{V718 Per} was monitored photometrically from the optical to the near-infrared (NIR). We also obtained high-resolution optical spectra with the Keck telescope at minimum as well as at maximum brightness. We derived the fundamental photospheric parameters of this star by comparing with synthetic spectra.} {Our photometric data show that the eclipses are very symmetric and persistent, and that the extinction law of the foreground occulting dust deviates only little from what is expected for ``normal'' interstellar material. The stellar parameters of \\mbox{V718 Per} indicate a primordial abundance of Li and a surface temperature of $T_{\\rm eff} \\approx 5200$ K. Remarkably, the in-eclipse spectrum shows a significant broadening of the photospheric absorption lines, as well as a slightly lower stellar surface temperature. In addition, weak emission components appear in the absorption lines of H$\\alpha$ and the Ca II IR triplet lines. We did not detect any signs of atomic or molecular features related to the occulting body in the in-eclipse spectrum. We also found no evidence of any radial velocity changes in \\mbox{V718 Per} to within about $\\pm 80$ m s$^{-1}$, which for an edge-on system corresponds to a maximum companion mass of ${\\sim}6~M_{\\rm Jup}$.} {Our observations suggest that \\mbox{V718 Per} is a single star, and thus very different from the related binary system KH 15D. We conclude that \\mbox{V718 Per} is surrounded by an edge-on circumstellar disk with an irregular mass distribution orbiting at a distance of 3.3 AU from the star, presumably at the inner disk edge. To produce the prolonged eclipses, the occulting feature must extend along more than half of the inner disk edge. The change in stellar surface temperature and the emission line activity observed could be related to spot activity. We ascribe the broadening of photospheric absorption lines during the eclipse to forward scattering of stellar light in the circumstellar dust feature.} ",
        "introduction": " ",
        "conclusions": "Our new photometric data of \\mbox{V718 Per} (H187) extends previous measurements and confirms that this object shows long-lasting eclipses with a period of 4.7 years. The eclipses, which are very symmetric,  are caused by occulting dust, and the colour changes suggest an extinction law of the foreground dust that deviates only little from what is expected for ``normal'' interstellar grains.  It has been speculated that \\mbox{V718 Per} may be similar to the unusual close binary KH 15D, also showing periodic eclipses, which presumably are caused by different coverage of the orbiting stars by a circumbinary disk. We obtained two high-dispersion spectra of \\mbox{V718 Per}, the first close to the deepest point of the eclipse and the other at a time outside eclipse with a time difference corresponding to roughly a quarter of the eclipse period. From these spectra we found no evidence of any change in the radial velocity of the star to within $\\pm 80$ m s$^{-1}$. Although we cannot be fully certain on the basis of two spectra, the absence of radial velocity variations in a system that experiences eclipses from its circumstellar material makes it very unlikely that \\mbox{V718 Per} is a close binary. \\mbox{V718 Per} seems therefore very different from KH 15D.  We confirm that the eclipses are caused by a stable extended, dusty structure orbiting at ${\\sim} 3.3$ AU from the star in an edge-on circumstellar disk. We have derived the SED of \\mbox{V718 Per} outside eclipse, including IR fluxes obtained with {\\it Spitzer}. This indicates that the star could be surrounded by a thin, low-mass disk. Because of the extended eclipse duration, the structure that causes the eclipses must extend over half a circle along the disk edge. A low-mass companion can in principle induce this structure, but the mass of any planet cannot exceed $6~M_{\\rm Jup}$ (assuming an edge-on system). The full amplitude of an eclipse amounts to 1.1 magnitudes in the V-band, but there are no signs of any enhanced absorption features from circumstellar gas during eclipse. It appears that the occulting structure is rather void of gas.  \\mbox{V718 Per} has a typical late-type absorption line spectrum without strong emission lines of e.g H$\\alpha$. We have used theoretical synthetic spectra from model atmosphere and derived fundamental photospheric parameters. Our spectroscopic analysis shows that \\mbox{V718 Per} has a primordial abundance of Li and a surface temperature of $T_{\\rm eff} \\approx 5200$\\,K. From the luminosity derived from its SED, and by comparing with theoretical evolution tracks, we find that \\mbox{V718 Per} is in its post-T Tauri phase of evolution.  However, there are remarkable differences between the in-eclipse and out-of-eclipse spectrum. During the eclipse several spectral features show that the surface temperature is slightly lower than in the out-of-eclipse spectrum, corresponding to a change in spectral type from G9 to K0. In addition, narrow emission components appear in the absorption cores of the H$\\alpha$ and the Ca {\\sc ii} IR triplet lines during eclipse, and the photospheric absorption lines become slightly broader. The change in stellar surface temperature and the emission line activity observed is puzzling. {\\it Since \\mbox{V718 Per} shows no short-term (rotational) photometric variability, this cannot be explained as the result of a variable coverage by starspots.} However, the observed spectral changes could be related to long-term changes in the activity near the polar regions. The broadening of the absorption lines we ascribe to forward scattering of stellar light in the circumstellar dust feature when it passes through the line-of-sight."
    },
    "0808/0808.1633_arXiv.txt": {
        "abstract": "By using images taken with WFCAM on UKIRT and SofI on the NTT and combining them with 2MASS we have measured proper motions for 126 L and T dwarfs in the dwarf archive. Two of these L dwarfs appear to have M dwarf common proper motion companions, and 2 also appear to be high velocity dwarfs, indicating possible membership of the thick disc. We have also compared the motion of these 126 objects to that of numerous moving groups, and have identified new members of the Hyades, Ursa Major and Pleiades moving groups. These new objects, as well as those identified in \\citet{jameson08} have allowed us to refine the L dwarf sequence for Ursa Major that was defined by \\citet{jameson08b}. ",
        "introduction": "Brown dwarfs may be thought of as failed stars. These low mass ($\\leq$70 M$_{\\rm Jup}$ \\citealt{burrows01}), cool objects are the lowest mass objects that the star formation process can produce. The majority of the brown dwarfs that have been discovered to date are field objects found using surveys such as the Two Micron All Sky Survey (2MASS; \\citealt{skrutskie06}, see \\citealt{leggett02} for examples), the DEep Near-Infrared Sky survey (DENIS; \\citealt{denis05}, see \\citealt{delfosse99} for examples), the Sloan Digital Sky Survey (SDSS;\\citealt{york00} see \\citealt{hawley02} for examples) and the UKIRT Deep Infrared Sky Survey (UKIDSS; \\citealt{lawrence06}, see \\citealt{kendall07, lodieu07b} for examples). However, to study brown dwarfs in depth, a knowledge of their age is essential, which means we must study brown dwarfs in open star clusters or moving groups.  Once a brown dwarf has been proved to belong to an open star cluster, or a moving group, then its age is known, allowing meaningful comparisons to evolutionary models to be made. The most recent example of this is the study done by \\citet{bannister07} who used existing proper motions and parallax measurements to show that a selection of field dwarfs in fact belong to the Ursa Major and Hyades moving groups. The importance of this study, is that these are the first brown dwarfs to be associated with an older cluster or group (age $>$200 Myr). Older clusters such as the Hyades are expected to contain very few or no brown dwarfs or low mass members, due to the dynamical evolution of the cluster over time \\citep{adams02}. However, these escaped low mass objects may remain members of the much larger moving group that surrounds the cluster.  \\citet{jameson08} followed this work by using the wide field camera (WFCAM, \\citealt{casali07}) on the United Kingdom Infrared Telescope (UKIRT) to image 143 known field L dwarfs. These images provided a second epoch for proper motion measurements, when combined with existing 2MASS images, typically taken 7 years previously. Using the proper motions and a distance calculated using the spectral type of the L dwarf given by \\citet{cruz03}, the moving group method was applied, and all 143 objects were scrutinised to check if their direction and magnitude of motion made them candidates of the many moving groups known. Members of the Hyades, Ursa Major and Pleiades moving groups were found. Radial velocity measurements such as those of \\citet{zapatero07} are required however, before it can be determined if these moving group members  are cluster members that have ``escaped'' as the cluster has dynamically evolved \\citep{adams02}. Is should also be noted that galactic resonances can produce effects similar to moving groups, and so all members may not be coeval \\citep{dehnen98}.  To continue the study started by \\citet{bannister07} and \\citet{jameson08}, we have measured proper motions for the majority of the remaining known field L dwarfs listed in the online L and T dwarf archive (http://spider.ipac.caltech.edu/staff/davy/ARCHIVE/). This has again been accomplished using the WFCAM on UKIRT and for the more southern objects, Son of ISAAC (SofI) on the 3.58m ESO New Technology Telescope (NTT). Using these wide field images and existing catalogue data, we have measured proper motions for an additional 126 L and T dwarfs listed in the dwarf archive.  These proper motion data may be put to a number of uses. Using reduced proper motion diagrams they can be used as an approximate measure of distance. The proper motion measurements can also be used to help identify objects as members of a star cluster or members of a moving group via the moving cluster method.  Taken with measured radial velocities and distances, it can yield all three components of velocity (U,V,W). As brown dwarfs tend to be faint, measuring their radial velocities is very difficult. As a result,  very few L or T dwarfs have measurements, none of which are in this sample. This means we cannot determine whether these moving group members are escaped cluster members or otherwise.  Our proper motion data are discussed and listed in section 2 of this paper. ",
        "conclusions": ""
    },
    "0808/0808.3400_arXiv.txt": {
        "abstract": "We present a framework for analyzing weak gravitational lensing survey data, including lensing and source-density observables, plus spectroscopic redshift calibration data.  All two-point observables are predicted in terms of parameters of a perturbed Robertson-Walker metric, making the framework independent of the models for gravity, dark energy, or galaxy properties.  For Gaussian fluctuations the 2-point model determines the survey likelihood function and allows Fisher-matrix forecasting.  The framework includes nuisance terms for the major systematic errors: shear measurement errors, magnification bias and redshift calibration errors, intrinsic galaxy alignments, and inaccurate theoretical predictions.   We propose flexible parameterizations of the many nuisance parameters related to galaxy bias and intrinsic alignment. For the first time we can integrate many different observables and systematic errors into a single analysis. As a first application of this framework, we demonstrate that: uncertainties in power-spectrum theory cause very minor degradation to cosmological information content; nearly all useful information (excepting baryon oscillations) is extracted with $\\approx 3$ bins per decade of angular scale; and the rate at which galaxy bias varies with redshift substantially influences the strength of cosmological inference.  The framework will permit careful study of the interplay between numerous observables, systematic errors, and spectroscopic calibration data for large weak-lensing surveys. ",
        "introduction": "Weak gravitational lensing of background sources can produce exceptionally strong constraints on cosmological parameters and tests of General Relativity.  Initial analyses considered the 2-point correlation function (or, equivalently, power spectrum) of the shear pattern induced on a single population of background galaxies \\citep{Jordi, Kaiser92, Blandford91}.  A wealth of new statistics, however, have been suggested as more powerful means to extract information from weak lensing (WL): cross-power spectra of multiple source populations with distinct redshift distributions (a.k.a. ``tomography'') \\citep{Hu}; the correlation of shear with foreground galaxy clusters \\citep{JainTaylor}, or more generally the cross-correlation of lensing shear with the galaxy distribution \\citep{BJ04, Zhang}; joint analyses of density-density, density-shear, and shear-shear correlations in an imaging survey \\citep{HJ04}; cross-correlation of magnification as well as shear \\citep{Jainmag}; use of the CMB \\citep{HO02,HS03} or recombination-era 21~cm signals \\citep{MW07,Pen04,ZZ06} as source planes; cross-correlation of source density or shear with a distinct spectroscopic galaxy survey population \\citep{Newman, SKZC}; and the use of 3-point statistics \\citep{TJ04} or statistics such as peak counts \\citep{Hennawi,Wang,Laura} to move beyond 2-point information.  Each of these potential innovations has been individually analyzed and shown to improve cosmological constraints.  The first goal of this paper is to consider the {\\em simultaneous} use of all of these observable statistics: can we forecast the cosmological information that they will yield collectively in future surveys?  Can we start to develop a framework in which all these signals could be analyzed simultaneously in a real experiment?  In parallel with the increasing variety of proposed WL signals, the community has identified a series of potential astrophysical and instrumental non-idealities in WL data which, if ignored, would lead to substantial systematic errors in the inferred cosmology.  These include: finite accuracy in our ability to predict the deflecting mass power spectrum due to nonlinearities \\citep{JS97} and baryonic physics \\citep{Zhan06, Jing}; intrinsic alignments (IA) between galaxy shapes \\citep{CroftMetzler} and between galaxy shapes and the local mass distribution \\citep{Hirata2} that are not induced by lensing; multiplicative ``shear calibration'' errors in the derivation of lensing shear from galaxy images \\citep{Ishak04,HTBJ}; additive ``spurious shear'' due to uncorrected PSF ellipticity or other imaging systematics \\citep{HTBJ,AmaraRefregier}; and errors in the assignment of redshifts to the source populations \\citep{MaHutererHu}.  The impact of these systematic-error sources on cosmological inferences have been analyzed by different means, but a second goal of this paper is to produce a comprehensive forecast that considers the presence of them all simultaneously.  Previous work has shown that these multiple sources of information and systematic in WL surveys can interact in interesting ways.  For example, in the presence of tomographic data, many systematics are readily distinguishable from cosmological signals and can hence be diagnosed and corrected internally to a survey; this approach is called {\\em self-calibration} \\citep{HTBJ}.  It has also been shown that combining galaxy density and lensing correlations can lead to self-calibration of shear calibration errors \\citep{BJ04} and the uncertainties in galaxy biasing \\citep{HJ04, Zhan06}.  Intrinsic alignments of galaxies can be diagnosed and corrected if tomographic information is available \\citep{KingSchneider}, however this places substantially greater demands on the precision and accuracy of redshift assignment than would otherwise be needed \\citep{BridleKing}. These investigations raise important practical questions: will the self-calibration techniques continue to succeed when we attempt to simultaneously self-calibrate several different systematic errors?  Do cross-correlation techniques reduce uncertainties in redshift distributions to negligible levels, or is it necessary to make a complete spectroscopic redshift survey of some size to measure redshift distributions directly \\citep{MaBernstein}?  This paper will present a formalism through which all these questions can be answered, but we defer to later papers the application of the framework to these issues.  A third goal of this work is to describe the constraints by WL in a language that is not tied to a specific cosmological model.  Most forecasts for WL survey constraints are done within the context of a Universe that has homogeneous dark energy with equation of state $w=w_0 + w_a(1-a)$.  Projecting the WL experiment onto this model gives concrete predictions, but obscures what the WL is really measuring.  So the analysis framework presented here will be {\\em dark-energy agnostic}, meaning that no specific model is assumed. We will be very explicit about the assumptions made in the analysis and try to keep them to a minimum.  In fact a great strength of WL experiments are their ability to test General Relativity itself, so we seek an analysis method that is general enough to incorporate such tests. Similar to the approach of \\citet{KST}, our analysis results in constraints on the distance and growth functions $D(z)$ and $g_\\phi(z)$, without reference to the particular dark-energy or gravity modifications that might cause deviations from $\\Lambda$CDM.  In the following section we describe a ``kitchen-sink'' formalism for WL survey observables that allows the incorporation of all suggested 2-point statistics and very general treatments of nearly all proposed systematic errors.  In \\S\\ref{speclike} we give a likelihood function and Fisher matrix for an unbiased spectroscopic redshift survey of source galaxies.  Then we briefly describe a software implementation of the lensing and spectroscopy likelihood calculations. We describe our model for the evolution of the lensing-potential power spectrum in \\S\\ref{powermodel}, and \\S\\ref{nuisance} we describe generic models used for the nuisance functions required in the lensing-survey analysis. In \\S\\ref{tuneup} we use the implementation of these methods to investigate the proper choices for the bin sizes and grid spacings needed to turn the lensing analysis into a tractable finite-dimensional problem.  Further application of the framework to survey forecasting will be done in future papers.  An earlier version of this WL analysis formalism was used to generate forecasts for the Dark Energy Task Force \\citep{DETF}, and is described in an appendix to that report. ",
        "conclusions": "The core of this paper are the expressions (\\ref{gg1})--(\\ref{kk1}) for the two-point correlation matrix of the lensing and density observable multipoles produced by a typical lensing survey.  This was derived under a very limited set of assumptions: a homogeneous and isotropic 4-dimensional metric Universe with scalar perturbations; plus the weak-lensing limit, the Limber and Born approximations, and an approximation that lensing magnification bias and intrinsic density fluctuations are additive. The last four assumptions could be relaxed at the expense of computational complexity.  We thus hope that data analyses based on this framework could be used to constrain a wide variety of potential explanations for the acceleration phenomenon, including gravity modifications as well as new fields in the Universe.  In the limit of Gaussian fluctuation fields, the two-point information is a complete description of the likelihood and hence can be used to construct Fisher matrices or analyze data. As currently configured, the analysis yields the survey's ability to constrain the distance function $D(z)$ and linear growth function $g_\\phi(z)$, without reference to particular dark-energy models.  It would be straightforward to implement scale-dependent linear-growth functions.  This framework subsumes all of the information (up to 2-point level) that is likely to be obtained from lensing observations: density-density, lensing-density, and lensing-lensing correlations, plus redshift distributions from unbiased spectroscopic surveys (\\S\\ref{speclike}).  Furthermore it allows for the most important expected forms of systematic error: photo-z calibration errors, shear and magnification-bias calibration errors, intrinsic alignments, and inaccuracies in power-spectrum theory.  Systematics that are additive to shear (\\eg uncorrected PSF ellipticity) or to density (\\eg uncorrected foreground extinction) have not been included.  We have not done so since the additive errors could, in principle, exhibit almost any arbitrary signature in the covariance matrix of the observables.  Hence a completely general model for additive errors would be degenerate with almost all other signals. For the additive systematics, it is better to determine the level at which they would bias the cosmological results than to attempt to fit a model.  \\citet{AmaraRefregier} is a good example of this approach.  Since the analysis framework is independent of models for dark energy, gravity, power-spectrum evolution, or galaxy bias, we get a stripped-down look at what parameters are truly constrained by the data, and what nuisance functions must be modeled in order to extract the cosmological information.  There is a substantial suite of biases and correlation functions involved in understanding the full survey data.  In other work these have been ignored, or have been quantified by reference to halo occupation models \\citep{HJ04,CvdBMLMY}. Here we introduce generic functions for bias and calibration nuisance functions that are not based on any particular physical model.  We implement one possible model for the evolution of the lensing-potential power spectrum, based on General Relativity but allowing for failure of the growth equation.  It is straightforward to implement other potential deviations from General Relativity.  In the current implementation, the end result of the Fisher analysis is a forecast of the ability to constrain the functions $D_A(z)$ and $g_\\phi(z)$.  Since the analysis must be discretized in redshift and angular scale in order to be feasible, we investigated the bin sizes or bandwidths of nuisance functions that should be chosen.  We find that $\\approx 3$ bins per decade of angular scale suffice to extract all information (apart from baryon acoustic oscillations), and that nuisance functions should be specified no finer than this.  Nuisance functions for power-spectrum theory errors and for shear and magnification-bias calibration errors can be specified coarsely in redshift space ($\\Delta \\ln a \\approx 1$), but the galaxy biases, correlations, and intrinsic-alignments must be modeled with potentially finer structure in redshift ($\\Delta \\ln a \\approx 0.1$) to immunize against potential astrophysical systematics.  In future papers we will use this framework and its implementation to investigate the requirements for spectroscopic calibration of photo-z's in large lensing surveys, and other practical issues.  As a simple first application of our framework, we have shown here that power-spectrum theory uncertainty does not significantly degrade the cosmological power of a nominal lensing survey at $10<\\ell<10^4$. Non-Gaussian statistics are a much more important factor to consider.  {\\tt C++} Code to implement Fisher forecasting using this framework has been written and runs quickly on desktop computers despite the large number of free parameters in these general models.  Interested parties should contact the author for access to the code."
    },
    "0808/0808.0129_arXiv.txt": {
        "abstract": "Tightest known quadruple systems VW~LMi consists of contact eclipsing binary with $P_{12}$ = 0.477551 days and detached binary with $P_{34}$ = 7.93063 days revolving in rather tight, 355.0-days orbit. This paper presents new photometric and spectroscopic observations yielding 69 times of minima and 36 disentangled radial velocities for the component stars. All available radial velocities and minima times are combined to better characterize the orbits and to derive absolute parameters of components. The total mass of the quadruple system was estimated at 4.56 M$_\\odot$. The detached, non-eclipsing binary with orbital period $P$ = 7.93 days is found to show apsidal motion with $U \\approx 80$ years. Precession period in this binary, caused by the gravitational perturbation of the contact binary, is estimated to be about 120 years. The wide mutual orbit and orbit of the non-eclipsing pair are found to be close to coplanarity, preventing any changes of the inclination angle of the non-eclipsing orbit and excluding occurrence of the second system of eclipses in future. Possibilities of astrometric solution and direct resolving of the wide, mutual orbit are discussed. Nearby star, HD95606, was found to form loose binary with quadruple system VW~LMi. ",
        "introduction": "\\label{intro}  The photometric variability of VW~LMi (HIP~54003, HD~95660, sp. type F3-5V, $V_{max}$=8.0), was found by the Hipparcos mission \\citep{hipp}, where it was correctly classified as a W~UMa-type eclipsing binary with an orbital period of 0.477547 days. The first ground-based photometric observations of the system obtained in 1999 and 2000 (taken in the $B$ and $V$ Johnson filters) were published by \\citet{dumi2000}. Analysis of these light curves \\citep{dumi2003} lead to the photometric mass ratio $q_{ph}$ = 0.395 and inclination $i$ = 72.4\\degr and contact configuration for the system. Later \\citet{gome2003} presented new $BV$ photometry and its preliminary analysis. Assuming convective envelopes for both components and the temperature of the primary as $T_1$ = 6700 K the authors estimated  the mass ratio as $q=0.4$, and inclination around 70\\degr. Fourier analysis of its Hipparcos light curve (hereafter LC) presented by \\citet{sela2004} yielded quite different parameters: $q = 0.25$, $i$ = 72.5\\degr~and fill-out factor $f=0.4$. The discovery of the second (non-eclipsing) binary in VW~LMi by \\citet{ddo11} makes all previous photometric solutions almost useless due to strong light contribution of the second pair of about $(L_3 + L_4)/(L_1 + L_2) = 0.42$ (at the maximum brightness of the contact pair) which was not taken into account.  \\citet{ddo11} presented long-term spectroscopy (209 spectra taken between 1998 and 2005) of the system obtained at David Dunlap Observatory (hereafter DDO) which enabled to disentangle all three orbits in this tight multiple system: the contact binary with the $P_{12}$ = 0.4775 days period is orbiting another binary with $P_{34} = 7.93$ days in a relatively tight, 355-days, mutual orbit. Using preliminary inclination angle of the contact-binary orbit found by photometric analysis, $i_{12}$ = 80.1$\\degr$, the authors determined masses of all components and found that the orbits of the binaries are not coplanar. Light-time effect (hereafter LITE) with peak-to-peak range of $2A$ = 0.0074 days was predicted to be seen in the minima of the contact pair as a result of the mutual revolution of the binaries. The LITE was found in published minima by \\citet{bajo2007}. The corresponding orbital parameters, $A$ = 0.0037(4) days, $P_{1234}$ = 353(2) days, $e$ = 0.5(2) and $\\omega$ = 2.8(4) rad, are rather preliminary due to few available minima. The eccentricity corresponding to their orbital solution is much higher than predicted spectroscopically by \\citet{ddo11}.  The spectral type of VW~LMi was estimated as F5V by \\citet{ddo11}, while observed Tycho-2 $(B-V)$ = 0.21 and 2MASS $J-K$ = 0.34 colors correspond to F2V spectral type. Both determined spectral type and colors refer to the whole quadruple system.  VW~LMi is tightest quadruple system known \\citep{toko2008}. Also it has the shortest period of the outer orbit within multiple systems harboring contact binaries. The ratio of the outer orbital period and orbital period of the non-eclipsing pair is only about $P_{1234}/P_{34}$ = 44.5 hence we can expect secular orbital changes on the timescales as short as decades. The chances of resolving of components (binaries) by either speckle or long-baseline interferometry are rather meager: the maximum angular separation of the components was estimated to be only about 10mas \\citep{ddo11}. Astrometric observations of VW~LMi do not indicate its multiplicity.  Goals of the present paper are as follows: (i) to present and analyze new photometric and spectroscopic observations, (ii) to perform simultaneous solution of LITE using both minima times of the contact binary and radial velocities (hereafter RVs) of the individual components, (iii) to assess possibility of the tidal disturbances of the inner orbits and resulting precession, (iv) determine absolute parameters of all four components. ",
        "conclusions": "Tightest known quadruple system VW~LMi is really unique. It definitely deserves additional observations and analysis. The study of this system could bring light on (i) the evolution and origin of binary stars in multiple systems of stars (ii) tidal interaction of third body and its influence on the Roche geometry (iii) long-term evolution of orbits in tight multiple systems. The system is useful for further analysis since all four components are visible in the spectrum. The determination of the individual component's parameters like $T_{eff}$, $\\log g$, metalicity could benefit from spectra disentangling (see \\citet{hadr1995}). The absolute parameters of all components which had to originate at practically same time could be used to determine the evolutionary status and history of the system.  The study of VW~LMi is, however, complicated by (i) the outer orbital period being close to one year which complicates its phase coverage by Earth-bound observer (ii) very small angular separation of the components making direct resolving of the mutual wide orbit and reliable determination of $i_{1234}$ difficult (which would definitely improve determination of individual masses), (iii) fast rotation of the components of the contact binary making RV measurements unsure and rectification to real continuum impossible (due to the line-blanketing) (iv) eclipses and Roche geometry in the contact pair making usual assumption of the spectra disentangling techniques not valid (line profiles of individual components cannot change with orbital revolution).  It is also interesting to note, that VW~LMi and HD95606 show very similar proper motion, parallax and RVs (see Table~\\ref{tab02}) - the stars definitely form a loosely bound pair and all components very probably evolved from the same protostellar cloud (see \\citet{oswa2007}).  The orbit of the detached pair in VW~LMi with 7.93-days period is almost circular. This is the case of another two quadruples detected by the DDO observations, TZ~Boo and V2610~Oph (DDO series No. XIV, \\citet{ddo14}). That means that either it evolved with such orbit or it was circularized by the gravitational interaction with the contact binary. Detected apsidal motion in VW~LMi requires further observing to reliably determine apsidal period. Investigation of the observed eccentricities of the second binaries and their predicted synchronization timescales in quadruple systems with contact binaries could shed light on the age and evolutionary status of contact binaries \\citep{ruci-priv2008}.  Better characterization of VW~LMi calls for long-term monitoring to cover the whole 355-days orbital cycle. Especially, times of minima should be obtained free of any systematic effects. The understanding of the system would greatly benefit from visual orbit obtained by means of long-baseline interferometry. Multi-color photometry and/or echelle spectroscopy could lead to reliable determination of component's temperatures and luminosities.  \\medskip  The stays of TP at DDO have been supported by a grant to Slavek M. Rucinski from the Natural Sciences and Engineering Council of Canada. This research has been supported in part by the Slovak Academy of Sciences under grants No. 2/7010/7 and 2/7011/7, and grant of the \\v{S}af\\'arik University VVGS 9/07-08. MV's research is supported by a Marie Curie ``Transfer of Knowledge'' Fellowship within the 6th European Community Framework Programme. The observations at Astronomical Observatory at Kolonica Saddle (part of Vihorlat Observatory) were partially supported by APVV grant LPP-0049-06 and APVV bilateral grant SK-UK-01006. DB thanks to Miron Kerul-Kmec for technical assistance with CCD camera at the Roztoky Observatory."
    },
    "0808/0808.0403_arXiv.txt": {
        "abstract": "{The magnetorotational instability (MRI)  of differential rotation  under the simultaneous presence of axial and azimuthal  components of the (current-free) magnetic field is considered. For rotation with uniform specific angular momentum the  MHD equations for axisymmetric  perturbations  are solved  in a local short-wave  approximation. All the solutions  are overstable for $B_z \\cdot B_\\phi\\neq 0$ with eigenfrequencies   approaching the viscous frequency. For more flat rotation laws the results of the local approximation do not comply with the results of a global calculation of the MHD instability of   Taylor-Couette flows between rotating cylinders. -- With  $B_\\phi$  and $B_z$  of the same order the traveling-mode solutions  are  also prefered for flat rotation laws such as the quasi-Kepler rotation.  For magnetic Prandtl number ${\\rm Pm}\\to 0$ they scale  with the   Reynolds number  of  rotation rather than with the magnetic Reynolds number (as for standard MRI) so that they  can easily be realized in MHD laboratory experiments. -- Regarding the  nonaxisymmetric modes one  finds a remarkable influence of the ratio  $B_\\phi/B_z$ only for the extrema. For $B_\\phi\\gg B_z$ and for not too small  Pm the nonaxisymmetric modes  dominate the traveling axisymmetric modes. For standard MRI  with $B_z\\gg B_\\phi$, however,   the critical Reynolds numbers of the nonaxisymmetric modes  exceed the values for the axisymmetric modes by many orders  so that they   are never  prefered. } ",
        "introduction": "It has been shown in previous publications starting with  Hollerbach \\& R\\\"udiger (2005) that the magnetorotational instability (MRI)  under the presence of both current-free axial and azimuthal components of the magnetic field ('Helical' fields, hence HMRI) is always characterized by an eigenoscillation frequency. In combination with the vertical wavenumber the resulting instability pattern is thus an axisymmetric  wave  traveling along the rotation axis. In all our considerations the magnetic Prandtl number Pm plays the basic role. For $\\rm Pm \\to 0$ the HMRI scales with the Reynolds number Re rather than with the magnetic Reynolds number Rm as it does for  the standard magnetorotational instability   for  an axial magnetic field.   The questions arise whether this frequency reflects the geometry of the magnetic field and whether it is observable  with real astrophysical objects such as protoneutron stars and/or accretion disks. Also the relation to the  Azimuthal MagnetoRotational Instability (AMRI, see R\\\"udiger et al. 2007b) for (current-free) toroidal fields which is nonaxisymmetric and which scales with the magnetic Reynolds number  for $\\rm Pm \\to 0$ must be considered.  We find  the AMRI only weakly (if ever) influenced by the addition of an axial field which is not much stronger than the toroidal field.  For $\\rm Pm\\simeq 1$ the difference between Re and Rm disappears so that the main differences of the instabilities also disappear.    In the present paper we start with a  local  approximation  using   analytical methods for the most simple rotation law for constant specific angular momentum, i.e. the Rayleigh limit. Global calculations of the stability of the same rotation law  between two rotating perfect-conducting cylinders and threaded by a helical current-free  magnetic field lead to an almost perfect coincidence of the results of both methods. This is no longer true, however, for more flat rotation laws such as quasikeplerian rotation in a finite gap where the differences of the short-wave approximation  (which only holds for infinitely thin gaps) and global models are so strong that the results completely differ (R\\\"udiger \\& Hollerbach 2007). ",
        "conclusions": ""
    },
    "0808/0808.0918_arXiv.txt": {
        "abstract": "We present stellar velocity dispersion measurements in the host galaxies of 10 luminous quasars (M$_{V}<-$23) using the Ca H\\&K lines in off-nuclear spectra. We combine these data with effective radii and magnitudes from the literature to place the host galaxies on the Fundamental Plane (FP) where their properties are compared to other types of galaxies. We find that the radio-loud (RL) QSO hosts have similar properties to massive elliptical galaxies, while the radio-quiet (RQ) hosts are more similar to intermediate mass galaxies. The RL hosts lie at the upper extreme of the FP due to their large velocity dispersions ($\\langle \\sigma_{*} \\rangle$ = 321~km~s$^{-1}$), low surface brightness ($\\langle \\mu_{e}(r) \\rangle$ = 20.8~mag~arcsec$^{-2}$), and large effective radii ($\\langle R_{e} \\rangle$ = 11.4~kpc), and have $\\langle M_{*} \\rangle$ = 1.5 x 10$^{12}$ M$_{\\sun}$ and $\\langle M/L \\rangle$ = 12.4. In contrast, properties of the RQ hosts are $\\langle \\sigma_{*} \\rangle$ = 241~km~s$^{-1}$, $\\langle M_{*} \\rangle \\sim$ 4.4 x 10$^{11}$ M$_{\\sun}$, and M/L $\\sim$ 5.3. The distinction between these galaxies occurs at $\\sigma_{*}\\sim$~300~km~s$^{-1}$, R$_{e} \\sim$ 6~kpc, and corresponding M$_{*} \\sim$ 5.9 $\\pm$ 3.5 x 10$^{11}$ M$_{\\sun}$. Our data support previous results that PG QSOs are related to gas-rich galaxy mergers that form intermediate-mass galaxies, while RL QSOs reside in massive early-type galaxies, most of which also show signs of recent mergers or interactions. Most previous work has drawn these conclusions by using estimates of the black hole mass and inferring host galaxy properties from that, while here we have relied purely on directly measured host galaxy properties. ",
        "introduction": "A growing understanding of the connection between galaxies and their central black holes has emerged over the last decade. We now know that all galaxies with a bulge contain supermassive black holes \\citep{kormendy04} and that black hole mass is correlated with host galaxy stellar velocity dispersion \\citep{gebhardt00a,ferrarese00,tremaine02}. Furthermore, the inclusion of an amount of energy equal to that expected from AGN feedback to quench star formation above a critical halo mass in semi-analytic galaxy formation models \\citep{cattaneo06,dekel06} reproduces the galaxy demographics and bimodality of properties observed in large surveys (SDSS: \\citet{kauffmann03a,kauffmann03b,hogg03,baldry04,heavens04,cidfernandes05}; GOODS: \\citet{giavalisco04}; COMBO-17: \\citet{bell04}; DEEP/DEEP2: \\citet{koo03,koo05,faber07}; MUNICS: \\citet{drory01}; FIRES: \\citet{labbe03}; K20: \\citet{cimatti02}; GDDS: \\citet{mccarthy04}). These facts suggest that the growth mechanisms of the black hole and galaxy must be connected. However, details of the physical processes that make this connection, such as how AGN energy interacts with and is dissipated by surrounding halo gas, are not yet understood.  One way to investigate these processes is to understand the nature of the host galaxies. Does something in the galaxy trigger AGN activity? Do active quasars exist in galaxies with similar properties? Studies of AGN host galaxies have reached different conclusions. One group of collaborators, \\citep{MKDBOH99,HKDB00,NDKHBJ00} believes these objects to be predominantly normal massive ellipticals, including Nolan et al. (2001) who found that most quasar host galaxies had evolved stellar populations, 10~Gyr old, with only a very small amount of recent star formation.  However, \\citet{mil81}, in the first spectroscopic investigation of a sample of these objects, concluded that they are not normal luminous ellipticals, a result which was later confirmed with deeper spectroscopy from the Keck telescope \\citep{mts96,she01,miller03}. Moreover, \\citet{CS00,CS01} have seen evidence of star formation within the past 100~Myr in quasars hosts with far-infrared excesses using deep spectra from Keck.  There is still debate about whether the different results are due to better quality spectra taken closer to the nucleus using 8-10~m class telescopes, or whether real differences exist in the host galaxy properties of the different quasar samples studied \\citep{lacy06}. This question will no doubt be answered as more data are analyzed from the larger extra-galactic surveys.  In this work we study luminous quasars (M$_{V}<-$23) in which the galaxy is actively feeding the central supermassive black hole. It is here that we should be able to investigate connections between the black hole and its surrounding galaxy from which we can draw conclusions about how the black hole may or may not affect galaxy formation and evolution. We begin in this first paper by analyzing the structural properties of the QSO host galaxies with the use of directly measured stellar velocity dispersions, previously unobtainable for quasars this luminous. We use these data to place the host galaxies on the Fundamental Plane and ascertain their structural properties relative to other types of galaxies. In future papers we will investigate whether these objects follow the M$_{BH}$-$\\sigma$ relation and analyze the host galaxy stellar populations to look for indications of star formation activity relative to quasar activity.  This paper is organized in the following manner. In \\S \\ref{data_section} we describe the sample selection and data analysis, including our removal of scattered quasar light from the observed spectra and the measurement of stellar velocity dispersions. We also present the comparison objects from the literature that are used in our analysis. In \\S \\ref{result_section} we use the fundamental parameters derived in \\S \\ref{data_section} to analyze the Fundamental Plane locations and mass-to-light ratios of these objects. In \\S \\ref{discussion} we discuss the properties of our QSO host galaxies relative to the comparison objects and their implication that two different classes of objects are present in our sample. Finally, \\S \\ref{summary} summarizes our work and presents the main conclusions. Further details about our stellar velocity dispersion measurement limitations and potential biases can be found in Appendix \\ref{bias}. ",
        "conclusions": "}  We have for the first time directly measured host galaxy stellar velocity dispersions for very luminous (M$_{V}<-$23) quasars, including both radio-loud and radio-quiet objects, and analyzed their structural properties. We compare the properties of these host galaxies to those of normal early-type galaxies \\citep{bernardi03a}, giant early-type galaxies \\citep{bernardi06}, a sample of radio-quiet PG QSO hosts \\citep{dasyra07}, and galaxy merger remnants \\citep{rothberg06}. The following summarizes our main conclusions.  \\begin{enumerate} \\item{ The six radio-loud QSO host galaxies lie at the upper extreme the Fundamental Plane of early-type galaxies. They occupy this location due to their large velocity dispersions (with an average of 321~km~s$^{-1}$), large effective radii (average of 11.4 kpc), low surface brightness (average of 20.8~mag~arcsec$^{-2}$), and high M/L (average of 12.4). The properties of these radio-loud host galaxies are similar to those of giant early-type galaxies in the SDSS, although only one has the spectrum of a purely old giant elliptical galaxy.  } \\item{ The four radio-quiet QSO host galaxies reside on the Fundamental Plane among normal early-type galaxies and at the high end of other PG QSO hosts, with a mean velocity dispersion of 241~km~s$^{-1}$. Their surface brightness and effective radii are slightly higher than normal early-type galaxies. The M/L's are slightly below normal early-type galaxies.  } \\item{ The radio-loud hosts in our study are either elliptical galaxies or interacting, while the radio-quiet hosts show a mixture of spiral and elliptical structure. The distinction between the two groups is due to galaxy mass, inferred from measured structural parameters, rather than morphological galaxy type. The separation occurs at galaxy velocity dispersions of $\\sigma_{*}\\sim$~300~km~s$^{-1}$, effective radii of R$_{e} \\sim$ 6~kpc, and corresponding stellar masses of M$_{*} \\sim$ 5.9 $\\pm$ 3.5 x 10$^{11}$~M$_{\\sun}$. } \\item{ We confirm a correlation between radio luminosity and stellar velocity dispersion, and thus implied black hole mass, of the host galaxies that suggests a higher slope (L$_{5 GHz}\\varpropto$~M$_{BH}^{3.8}$ with rms scatter of 1.36 dex) than found by Franceschini et al.~(L$_{5 GHz}\\varpropto$~M$_{BH}^{2.66}$), though it could be consistent with previous work given the large scatter in this relation \\citep{mclure04}, our small sample size, the limited M$_{BH}$ range of our data, and differences in M$_{BH}$ estimated from different techniques. We find a tighter correlation between radio luminosity and host galaxy bulge mass, $L_{radio} \\sim M_{bulge}^{3.56}$ with rms scatter of 1.09 dex }  \\end{enumerate}"
    },
    "0808/0808.2630_arXiv.txt": {
        "abstract": "{} {We seek to determine whether the late-type star 2MASS J12354893$-$3950245 (2M1235$-$39) is a member of the TW Hya Association (TWA), a hypothesis suggested by its association with a bright X-ray source detected serendipitously by ROSAT and XMM-Newton and its ($\\sim3'$) proximity to the well-studied (A+M binary) system HR 4796.} {We used optical spectroscopy to establish the Li and H$\\alpha$ line strengths of 2M1235$-$39, and determined its proper motion via optical imaging. We also considered its X-ray and near-IR fluxes relative to the M star HR 4796B.} {The optical spectrum of 2M1235$-$39 displays strong Li absorption and H$\\alpha$ emission (equivalent widths of 630 m\\AA\\ and $-6.7$ \\AA, respectively). Comparison of the spectrum with that of a nearby field star, along with the DENIS catalog $IJK$ magnitudes, indicates the spectral type of 2M1235$-$39 is M4.5. We measure a proper motion for 2M1235$-$39 that agrees, within the errors, with that of HR 4796.} {The Li absorption and H$\\alpha$ emission line strengths of 2M1235$-$39, its near-IR and X-ray fluxes, and its proper motion all indicate that 2M1235$-$39 is a TWA member. Most likely this star is a wide (13,500 AU) separation, low-mass, tertiary component of the HR 4796 system.} ",
        "introduction": "As of little more than a decade ago, astronomers were almost oblivious to the presence of low-mass, pre-main sequence stars within $\\sim100$ pc of Earth. The intervening years have seen the identification of a few hundred such stars, with ages ranging from 8 to 100 Myr, as part of numerous post-T Tauri associations (Zuckerman \\& Song 2004, hereafter ZS04, and references therein; Torres et al.\\ 2006, 2008). Perhaps the greatest excitement associated with the recognition of the existence of nearby young stars has been the opportunity to study, at close range, the evolution of youthful planetary systems, via direct thermal imaging of warm massive planets (e.g., Chauvin et al.\\ 2004; Song et al.\\ 2006) and via imaging and spectroscopy of debris disks (e.g., Rebull et al.\\ 2008 and references therein). Young, local stellar groups also afford unique insight into the early evolution of low-mass stars and ultracool dwarfs (e.g., Looper et al.\\ 2007; Cruz et al.\\ 2008; and references therein).  The difficulty inherent in identifying young stars and young star groups near Earth reflects the fact that such groups are spread over large areas of the sky (ZS04). Furthermore, while the local young groups are usually ``spearheaded'' by a handful of well-studied, individual systems that feature, e.g., strong H$\\alpha$ emission, enormous IR excesses, and/or easily imaged debris disks (TW Hya and $\\beta$ Pic being cases in point), the vast majority of nearby young stars are otherwise unremarkable late-type (K through M) dwarfs that do not stand out or even turn up in optical emission-line or far-infrared (e.g., IRAS) surveys.  However, all $\\sim10$--100 Myr-old stars of types F through M are at or near the peaks of their lives in terms of their X-ray luminosities relative to bolometric (with ``saturated'' values of $L_X/L_{\\rm bol} \\sim 10^{-3}$, ZS04 [their Fig.~4]; see also Kastner et al.\\ 1997 and Preibisch \\& Feigelson 2005). Hence, X-ray point source catalogs, in tandem with recently released, comprehensive catalogs of distances and proper motions of stars in the solar neighborhood, have served as the main resources with which to isolate stars that are likely nearby and young. Followup optical spectroscopy and/or imaging then readily confirms (or refutes) membership in the ``nearby young star club,'' via determination of surface Li abundances and relative ($UVW$) Galactic space motions.  Here, we demonstrate that serendipitous XMM-Newton and ROSAT X-ray detections of 2MASS J12354893$-$3950245 (hereafter 2M1235$-$39), combined with its optical spectrum and proper motion, establishes this star as a member of the quintessential local young star group, the TW Hya Association (TWA; Kastner et al.\\ 1997; Webb et al.\\ 1999; Zuckerman et al.\\ 2001). Indeed, 2M1235$-$39 is, likely, the tertiary component of the well-studied HR 4796 (A+M star) binary system (Jura et al.\\ 1993; Stauffer et al.\\ 1995), which is designated TWA 11. ",
        "conclusions": "All measured properties of 2M1235$-$39 presented here --- its Li absorption and H$\\alpha$ emission line strengths, its near-IR and X-ray fluxes, and its proper motion --- are compatible with TWA membership. Based on these results and on the similarity of its common proper motion to that of HR 4796A, we conclude that 2M1235$-$39 is most likely a wide (13,500 AU) separation, low-mass, tertiary component of the HR 4796 system."
    },
    "0808/0808.0773_arXiv.txt": {
        "abstract": "We show that the smoothed particle hydrodynamics (SPH) method, used with individual time-steps in the way described in the literature, cannot handle strong explosion problems correctly.  In the individual time-step scheme, particles determine their time-steps essentially from a local Courant condition.  Thus they cannot respond to a strong shock, if the pre-shock timescale is too long compared to the shock timescale.  This problem is not severe in SPH simulations of galaxy formation with a temperature cutoff in the cooling function at $10^4~{\\rm K}$, while it is very dangerous for simulations in which the multiphase nature of the interstellar medium under $10^4~{\\rm K}$ is taken into account.  A solution for this problem is to introduce a time-step limiter which reduces the time-step of a particle if it is too long compared to the time-steps of its neighbor particles.  Thus this kind of time-step constraint is essential for the correct treatment of explosions in high-resolution SPH simulations with individual time-steps. ",
        "introduction": "Hierarchical (individual) time-step method \\citep[e.g.,][]{McMillan1986, HernquistKatz1989, Makino1991IndividualTimeStep} is widely used in simulations of galaxy formation and star formation based on smoothed particle hydrodynamics (SPH) method \\citep{Lucy1977, GingoldMonaghan1977}. This method allows particles to have different time-steps, and can significantly reduce the total calculation cost when there is a large variation in the timescales of particles. Almost all implementations of the individual time-steps method used for particle systems violates Newton's third law (see \\citeauthor{FarrBerschinger2007} \\citeyear{FarrBerschinger2007} for one of exception).  As long as physical quantities are integrated with sufficient accuracy, this violation is not a severe problem.  However, it is not always possible to maintain the accuracy.  To our knowledge, all existing implementations of individual time-step for SPH rely on the determination of the time-step at the end of previous time-steps. Therefore, if something unforeseen occurs during the time-step of one particle, the particle might fail to catch that event, resulting in a large integration error.  A supernova (SN) explosion is an example of such an event.  A SN generates a small amount of very hot gas ($T \\sim 10^8~{\\rm K}$) in a large clump of cold gas ($T \\sim 10~{\\rm K}$).  The difference in the time-steps of the hot gas and surrounding cold gas particles becomes quite large (typically the difference reaches $\\sim 10^3$).  Thus, hot gas particles step forward $\\sim 10^3$ or more time-steps before neighboring cold gas particles respond to the SN event.  This means that the evolution of both hot and cold gas particles is completely wrong, since the surrounding cold gas particles do not react the explosion for a duration much longer than the timescale in which the blast wave would propagate the inter-particle distance.  This problem is not severe for ordinary SPH simulations of galaxy formation because of the temperature cutoff in cooling functions at $10^4~{\\rm K}$, while it becomes very serious for simulations involving the multiphase nature of the interstellar medium under $10^4~{\\rm K}$, because the mach number can be very high.  This problem of sudden change occurs in any dynamical simulation with individual time-steps, as long as the time-steps are determined with the usual explicit method.  In principle, a fully implicit method in which the time-step itself is also determined implicitly \\citep{Makino+2006} or a method which satisfies Newton's third law \\citep{FarrBerschinger2007} can solve this problem, but there are no implementations of such methods for SPH simulations yet {\\footnote {Recently, \\cite{Springel2009} has developed a mesh-based scheme employing individual time-steps where the smaller time-step is adopted for integrations between neighboring meshes. This scheme does satisfy Newton's third law.}}.  In simulation of star clusters, the SN and its kick introduces a sudden change in the orbit (and mass) of the exploded star. Here, a rather simple prescription in which either all stars or at least nearby stars are synchronized to the time of explosion and restart the integration has been used. This prescription, at least the version which synchronizes all particles in the system, is impractical for $N$-body/SPH simulations of galaxies, because the number of SN events and therefore the increase in the calculation cost is too great.  We propose a simple limiter for hydrodynamical time-steps in the individual time-step method which mitigates this problem.  With this limiter the behavior of an explosion integrated by individual time-steps becomes essentially the same as that integrated by global time-steps.  In \\S 2, we describe this limiter, and in \\S 3 we report the result of numerical experiments. ",
        "conclusions": ""
    },
    "0808/0808.1668_arXiv.txt": {
        "abstract": "We carried out a comparison of the signals seen in contemporaneous BiSON and GOLF data sets. Both instruments perform Doppler shift velocity measurements in integrated sunlight, although BiSON perform measurements from the two wings of potassium absorption line and GOLF from one wing of the NaD1 line. Discrepancies between the two datasets have been observed. We show,in fact, that the relative power depends on the wing in which GOLF data observes. During the blue wing period, the relative power is much higher than in BiSON datasets, while a good agreement has been observed during the red period. ",
        "introduction": "P-modes are standing acoustic waves in the solar interior. There are instances when it has been shown that solar flares are correlated with  modes of oscillation in the Sun just as the Earth is set ringing after a major earthquake. These events have been seen in high-degree modes (\\cite{fog},\\cite{kos}). We are interested in the possibility that the global modes are similarly stimulated and we seek to use the extensive BiSON dataset to search for these events. Large power events with strengths many times the mean level are seen and, moreover, the number is in excess of the predictions of stochastic excitation \\cite{Cha1}. Our long-term aim is to categorize these rare events through their temporal features in both power and phase with the hope of  being able to distinguish natural from forced events. As part of this process of categorization of large excitations, we carried out a comparison of the signals seen in contemporaneous BiSON and GOLF data. The intention is to show that the power and the phase returned for the signal is in agreement for the two datasets. Given that, we can go on to use the very extensive BiSON data on its own. We show here that relative power in the two datasets depends on the wing in which the GOLF data are observed. We compare the observed power ratio to the predictions based on the heights in the atmosphere that the respective spectral lines are formed. We show also that the relative phases returned by the analysis are comparable. ",
        "conclusions": "We have used the comparison of contemporaneous GOLF and BiSON data to demonstrate that the power and phase evolution of the solar oscillations data are consistent between the two datasets and are therefore likely to be true characteristics of the oscillations of the Sun and not instrumental in origin. The red-wing GOLF data were found to be most useful for this.  The BiSON data were shown not to be compromised by their lack of perfect fill nor their ground-based observations. This opens up the possibility to use the full BiSON dataset spanning  three solar cycles in duration for further study of the large excitations and their links with solar activity. The next step in this project is to continue with the categorization of events with the eventual hope of being able to identify those that are stimulated by activity on the Sun."
    },
    "0808/0808.3806_arXiv.txt": {
        "abstract": "{ In an attempt of clarifying the connection between the photospheric abundance anomalies and the stellar rotation as well as of exploring the nature of ``normal A'' stars, the abundances of seven elements (C, O, Si, Ca, Ti, Fe, and Ba) and the projected rotational velocity for 46 A-type field stars were determined by applying the spectrum-fitting method to the high-dispersion spectral data obtained with BOES at BOAO. We found that the peculiarities (underabundances of C, O, and Ca; an overabundance of Ba) seen in slow rotators efficiently decrease with an increase of rotation, which almost disappear at $v_{\\rm e}\\sin i \\ga 100$~km~s$^{-1}$. This further suggests that stars with sufficiently large rotational velocity may retain the original composition at the surface without being altered. Considering the subsolar tendency (by several tenths dex below) exhibited by the elemental abundances of such rapidly-rotating (supposedly normal) A stars, we suspect that the gas metallicity may have decreased since our Sun was born, contrary to the common picture of galactic chemical evolution. } ",
        "introduction": "Since unevolved A-type stars on (or near to) the upper main-sequence have masses around $\\sim 2 M_{\\odot}$, their surface abundances may retain information of the composition of the past galactic gas ($\\la 10^{9}$~yr ago) from which they formed; this would provide us with an important opportunity to investigate the ``recent'' chemical evolution history of the Galaxy. However, such a study using A stars as a probe of late-time history of chemical evolution has rarely been done in spite of its potential significance\\footnote{For example, Takeda, Sato, \\& Murata (2008) found in their extensive study of late-G giants (also having masses around $\\sim 2 M_{\\odot}$; evolved counterparts of A dwarfs) that their metallicities show an appreciable diversity as large as $\\sim 1$~dex ($-0.8 \\la$~[Fe/H]~$\\la +0.2$) with a subsolar trend on the average (cf. Fig. 14 therein), from which they argued that some special event (like a mixing of metal-poor primordial gas caused by infall) might have occurred $\\sim 10^{9}$~yr ago. However, before making any speculation, it is important to confirm whether such a trend is also observed in their progenitors in the upper main-sequence.} in contrast to the case of old-time history where a number of researches (using longer-lived F--G--K dwarfs) are available. This is related to the fact that a large fraction of them are rapid rotators (typically $v_{\\rm e} \\sin i \\sim$ 100--200~km~s$^{-1}$ on the average; see, e.g., Royer et al. 2002a, b) whose spectra are technically difficult to analyze because lines are broad and smeared out, while sharp-lined slow rotators easy to handle tend to show abundance anomalies (chemically peculiar stars or CP stars).\\footnote{Although many things are left unresolved concerning the origin and nature of the CP phenomena, it is widely considered (at least in the qualitative sense) that the chemical segregation in the stable atmosphere is responsible for the abundance anomalies at the surface: i.e., an element becomes over- or under-abundant depending on the balance of upward radiation force and the downward gravitational force. In this case, rotation would act against an efficient built-up of such anomalies because it enhances mixing of outer stellar layers via shear instability or meridional circulation.} As a matter of fact, most spectroscopic analyses of A stars have focused on sharp-lined ones with $v_{\\rm e}\\sin i \\la 50$~km~s$^{-1}$. For this reason, nobody could be sure whether the result obtained for star classified as ``normal A'' really reflects the initial composition free from any peculiarities.  Therefore, in order to make a further step forward, it is requisite to challenge abundance determinations for ``unbiased'' sample of A-type stars in general (i.e., without sidestepping rapidly rotating ones), which inevitably requires an application of the spectrum synthesis technique, since reliably measuring the equivalent widths of individual spectral lines is almost hopeless for rapid rotators. Admittedly, while such trials of determining abundances from spectra of A dwarfs including broad-lined ones have recently emerged thanks to the improvement in the method of analysis as well as the data quality, their interests are mainly directed to objects of specific types; e.g., Vega-like stars (Dunkin et al. 1997) or $\\lambda$ Bootis stars (Andrievsky et al. 2002) or open-cluster stars (Takeda \\& Sadakane 1997; Varenne \\& Monier 1999; Gebran et al. 2008; Gebran \\& Monier 2008; Fossati et al. 2008). Namely, a systematic study attempting to clarify the characteristics of normal field A-type stars in general, especially in terms of their abundance--rotation connection, seems to have been rarely attempted. To our knowledge, only one such study is Lemke's (1990, 1993) determinations of C and Ba abundances for some 20 rapidly-rotating A stars with $v_{\\rm e} \\sin i$ up to $\\sim$~200~km~s$^{-1}$, which however appear to be still insufficient and inconclusive as judged from his adopted method of approach as well as the number of elements studied.  Considering this situation, we decided to revisit this problem in our own manner based on the high-dispersion spectral data of $\\sim 50$ A-type stars in a wide range of $v_{\\rm e} \\sin i$ (0--300~km~s$^{-1}$) obtained with BOES at BOAO, while applying the automatic spectrum fitting algorithm (Takeda 1995) which efficiently enables determinations of the abundances (for selected six elements of C, O, Si, Ti, Fe, and Ba) even for rapid rotators showing considerably merged spectra. Our ultimate aim is to clarify the following questions of interest:\\\\ --- (1) Is there any systematic rotation-dependent tendency between slow and rapid rotators in terms of the abundance anomaly? If so, what is the critical value of $v_{\\rm e} \\sin i$, above which stars may be regarded as normal? \\\\ --- (2) What would the abundance characteristics of ``normal A-type stars'' like, which we may consider as retaining the composition of the galactic gas from which they formed? \\\\  We will show that reasonable answers to these points are provided from this study (cf. Sect. V). ",
        "conclusions": "\\subsection{Rotation--Abundance Connection}  Figures 9a--f display the resulting [X/H] vs. [Fe/H] correlations (for X = C, O, Si, Ca, Ti, and Ba), from which we can roughly divide these elements into three groups.\\\\ (i) Si and Ti: almost scaling in accordance with Fe.\\\\ (ii) C, O, and Ca: showing an anti-correlation trend with Fe.\\\\ (iii) Ba: positive correlation with Fe, though its range of peculiarity is much more conspicuous than that of Fe.  Besides, [C/H], [O/H], [Si/H], [Ca/H], [Ti/H], [Ba/H], and [Fe/H] are plotted against $v_{\\rm e} \\sin i$ in Figures 10a--g. We can recognize from these figures that [C/H], [O/H], [Ca/H] and [Ba/H] are systematically $v_{\\rm e} \\sin i$-dependent in the sense that the peculiarity (overabundance for Ba, underabundance for C/O/Ca) tends to decrease with an increase in $v_{\\rm e} \\sin i$. While such a convincing tendency is not apparent for the remaining elements (Si, Ti, and Fe), [Fe/H] appears to weakly conform to this trend (i.e., decreasing tendency with $v_{\\rm e} \\sin i$).  Combining these observational fact, we may conclude as follows:\\\\ --- (a) All the seven elements exhibit some kind of abundance peculiarities, which are more conspicuously seen in slow rotators ($v_{\\rm e}\\sin i \\la 50$~km~s$^{-1}$) and characterized by the deficiency of C, O, and Ca and the enrichment of Si, Fe, and (especially) Ba.\\\\ --- (b) These anomalies tend to diminish progressively with an increase in $v_{\\rm e} \\sin i$ (at least in the range of slow/moderate rotators of $\\la 100$~km~s$^{-1}$). \\\\ --- (c) The stellar rotational velocity must thus be the most important key factor in the sense that the extent of abundance peculiarity tends to be larger as a star rotates more slowly, which is presumably because some counter-acting mechanism of diluting the built-up anomaly (most probably due to the element segregation in a stable atmosphere/envelope) takes place in rapid rotators.  We also point out these tendencies seen in Figures 9 and 10 are more or less consistent with the results of recently published papers focused on the abundance trends of A-type stars (including Am stars) for a wide range of $v_{\\rm e} \\sin i$ values: e.g., Lemke (1990, 1993) [field stars;  C, Ba (elements in common with this study)], Savanov (1995a,b) [field stars; C, O, Si, Ca, Fe, Ba], Takeda \\& Sadakane (1997) [Hyades and field stars; Fe, O], Gebran, Monier, \\& Richard (2008) [Coma Berenices; C, O, Si, Ca, Fe, Ba], Gebran \\& Monier (2008) [Pleiades; C, O, Si, Ca, Fe, Ba], and Fossati et al. (2008)[Praesepe; C, O, Si, Ca, Fe, Ba].  \\subsection{Implication of Subsolar Compositions in Normal A Stars}  According to what we learned in Sect. V-a, we may assume that the abundance peculiarities of A-type stars (conspicuously seen slow rotators) tend to disappear for rapid rotators at $v_{\\rm e} \\sin i \\ga 100$~km~s$^{-1}$ (cf. Figure 10). If so, we would be able to gain information of the galactic gas $\\la 10^{9}$~yr ago by inspecting the photospheric abundances of such rapidly-rotating A-type stars, since they are considered to retain the composition of the gas from which they formed.  From this point of view, it is interesting to note in Figure 10 that the [X/H] values at the high-$v_{\\rm e} \\sin i$ range tend to be somewhat negative or ``subsolar'' for many elements such as C, O, Ti, Fe, and Ba; i.e., by several tenths dex below the solar (or Procyon) abundances on the average.  Here we recall Takeda, Sato, \\& Murata's (2008) conclusion that the [Fe/H] values (as well as those of other elements whose abundances almost scale with Fe) of evolved G giants, many of which have mass values around $\\sim 2 M_{\\odot}$ like A-type dwarfs, spread in a range of $-0.8 \\la$~[Fe/H]~$\\la +0.2$ around an average value of [Fe/H]~$\\sim -0.3$.  Considering these two observational consequences, we would conclude that the metallicities of the galactic gas $\\la 10^{9}$~yr ago had really a subsolar tendency (though with a rather large diversity). If this is the case, the gas metallicity of [Fe/H] $\\sim 0$ ($\\sim 5 \\times 10^{9}$ ago when our Sun was born) must have decreased by several tenths dex with an elapse of time until $\\la 10^{9}$~yr ago when A dwarfs (progenitors of G giants) were born. Although this trend does not seem to have been taken very seriously so far \\footnote{Meanwhile, a completely different solution to this problem has also been proposed, arguing the necessity of downward revision of the solar abundances as a result of the application of sophisticated 3D line formation theory; (cf. Asplund et al. 2004). While this possibility may be worth considering, it can not yet be regarded as reliable in our opinion, since it causes serious discrepancies between theory and observation in the solar interior model (see, e.g., Young 2005 and the references therein). Besides, some questionable points still remain in their line-formation treatment (see also Appendix 1 in Takeda \\& Honda 2005).} in spite of not a few supportive evidences\\footnote{ Actually, the apparent subsolar tendency in the photospheric abundances of comparatively young stars has often been reported; e.g., C/N/O in early B main-sequence stars (Gies \\& Lambert 1992, Kilian 1992, see also Nissen 1993); C/N/O/Si/Mg/Al in early B stars (Kilian 1994); [Fe/H] in superficial normal late B and A stars (Sadakane 1990); [Fe/H] of B stars from UV spectra (Niemczura 2003); O in supergiants (Luck \\& Lambert 1985; Takeda \\& Takada-Hidai 1998).} since it contradicts the conventional scenario of galactic chemical evolution (where elemental abundances are generally believed to increase with time), we tend to regard this tendency as real, which means that the gas metallicity actually {\\it decreased} in an elapse of time between the formation of our Sun ($\\sim 5\\times 10^{9}$~yr ago) and the formation of $\\sim 2~M_{\\odot}$ stars ($\\la 10^{9}$~yr ago). Of course, in order to make this hypothesis more convincing, a reasonable explanation has to be done why such a reduction of the gas metallicity had occurred against the intuitive chemical evolution picture of increasing metallicity. While one such interpretation might be the dilution of the metallicity caused by an substantial infall of metal-poor primordial galactic gas speculated by Takeda et al. (2008), further observations and extensive abundance analyses on a much larger number of rapidly-rotating A dwarfs (as well as evolved G giants) would be required until we can say something about it with confidence."
    },
    "0808/0808.3201_arXiv.txt": {
        "abstract": "We present sensitive, high angular resolution ($0\\rlap.{''}05$) VLA continuum observations made at 7 mm of the core of the HH~111/121 quadrupolar outflow. We estimate that at this wavelength the continuum emission is dominated by dust, although a significant free-free contribution ($\\sim$30\\%) is still present. The observed structure is formed by two overlapping, elongated sources approximately perpendicular to each other as viewed from Earth. We interpret this structure as either tracing two circumstellar disks that exist around each of the protostars of the close binary source at the core of this quadrupolar outflow or a disk and a jet perpendicular to it. Both interpretations have advantages and disadvantages, and future high angular resolution spectroscopic millimeter observations are required to favor one of them in a more conclusive way. ",
        "introduction": "It is generally accepted that most stars form in binary or multiple systems (Lada \\& Lada 2003). Furthermore, in the case of low and intermediate-mass stars it is also known that the process occurs with the presence of an accretion disk and a collimated outflow. From these two facts it follows that binary disk-jet systems should be common in regions of low-mass star formation. However, in practice these systems are hard to detect and identify and only a few have been studied in detail (e. g. Rodr\\'\\i guez et al. 1998; Anglada et al. 2004; Monin et al. 2007). Furthermore, it is still unclear if the members of a binary system will both be able to maintain the disks and outflows that characterize the formation of single stars. For example, the binary stars that form L1551~IRS5 are both believed to possess disks and jets (Rodr\\'\\i guez et al. 2003; Lim \\& Takakuwa 2006), while it is known that only one of the stars that forms the SVS 13 close binary system is associated with detectable circumstellar dust emission that is probably tracing a disk (Anglada et al. 2004).   One of the most interesting cases of a binary source that is known to exhibit independent outflows, most probably associated with each of the components of the binary, is HH~111 (Reipurth et al. 1999). Located in Orion at a distance of 414 pc (Menten et al. 2007), the optical HH 111 jet was discovered by Reipurth (1989). It is an extremely well-collimated jet, aligned approximately in the east-west direction (at a PA of $\\sim 97^\\circ$) and whose knots move in the plane of the sky with velocities of the order of several hundred km s$^{-1}$ (Reipurth, Raga \\& Heathcote 1992). Reipurth, Bally \\& Devine (1997) found that this optical jet is part of a giant HH complex extending over 7.7 pc. This giant HH complex is very straight, suggesting great stability over the $10^4$ years of its lifetime.  Near-infrared observations (Gredel \\& Reipurth 1993; 1994) revealed a second bipolar flow, named HH 121, that emerges from about the same position as the optical outflow, and is aligned approximately in the north-south direction (at a PA of $\\sim 35^\\circ$). This result suggested the presence of a close binary source in this region. Both the optical and the infrared outflows are detected as bipolar molecular outflows (Cernicharo \\& Reipurth 1996; Nagar et al. 1997; Lefloch et al. 2007).  At the center of the quadrupolar outflow is the source IRAS 05491+0247 = VLA 1, a suspected class I binary with a total luminosity of about 25 L$_\\odot$ (e.g. Stapelfeldt \\& Scoville 1993, Yang et al. 1997). To advance in our understanding of this source a high angular resolution image at millimeter wavelengths was needed to compare with the information available for the quadrupolar outflow. ",
        "conclusions": "Our main conclusions follow:  1) Our high angular resolution ($0\\rlap.{''}05$) VLA continuum 7 mm observations of the core of the HH~111/121 quadrupolar outflow reveal a structure that can be described as two overlapping, elongated sources that appear approximately perpendicular to each other in the plane of the sky.  2) We discuss possible interpretations for this structure and conclude that the most viable ones are that we are observing two orthogonal disks around separate protostars or a disk with a perpendicular jet. Both intepretations have advantages and disadvantages, and high angular resolution spectroscopic millimeter observations (possible only in the future with the Atacama Large Millimeter Array) are required to disentangle what is going on at the core of this quadrupolar outflow.           \\ack  We thank an anonymous referee for valuable suggestions. LFR acknowledges the support of CONACyT, M\\'exico and DGAPA, UNAM. JMT and GA are supported by the MEC AYA2005-05823-C03 grant (co-funded with FEDER funds). GA also acknowledges support from Junta de Andaluc\\'{\\i}a."
    },
    "0808/0808.0344_arXiv.txt": {
        "abstract": "The study of  extragalactic sources of high energy radiation via the direct measurement of  the  proton and neutrino fluxes that they are likely to emit is one of the main goals for  the   future observations   of  the recently developed air showers detectors and neutrino telescopes. In this work we discuss the relation between the inclusive proton and neutrino signals from the ensemble of all sources in the universe, and the ``resolved'' signals from the closest and  brightest  objects. We also  compare the sensitivities  of proton and neutrino  telescopes  and comment on the relation between these two new astronomies. ",
        "introduction": "There is a general consensus that   the highest energy cosmic  rays  (CR) are  of  extragalactic  origin   because  they reach the Earth  with an approximately isotropic  angular distribution.  The magnetic fields of the Milky Way are not sufficiently strong and extended to  randomize the directions of particles  produced   by  our own Galaxy,  and the   this isotropy   of these CR is likely to  reflect the  large scale    homogeneity of the universe. It is  natural to expect that the  extragalactic    Ultra High Energy   Cosmic Rays (UHECR) are produced in  ``point--like'' astrophysical  sources. The identification of these  sources, and the clarification of the  mechanisms that  accelerate  particles to these  very large energies  is clearly one of the crucial goals of  the   new large acceptance  air shower  detectors like the Pierre Auger Observatory.  The  trajectories  of  charged  particles is bent by the presence of astrophysical  magnetic  fields, however  at sufficiently  large rigidity $E/Z$ (with $Z$ the  particle electric  charge)  the magnetic deviations   should become sufficiently   small  to allow the direct  imaging of sources with cosmic rays. Unfortunately  the  intensity and  structure of   the intergalactic  magnetic  field are  very poorly known, and therefore   the   region  (in source distance and particle energy)  where  CR astronomy is possible remains  very uncertain. It is possible (and there  are  observational hints)  that  the CR above  $E \\gtrsim {\\rm few} \\times 10^{19}$~eV are already propagating in quasi--linear mode  for  source distances  of  order 100~Mpc or more, however the identification of the sources remains difficult because of the very small number of events available, with  most of the   sources contributing  with not more than a  single  event.  The analysis of the ``clustering'' of the observed events can give information about  the luminosity of the individual sources. After  the  data of AGASA  \\cite{Hayashida:1996bc,Takeda:1999sg}    gave  hints of possible clustering. this question   has received  a  significant amount of attention  \\cite{Uchihori:1999gu,Dubovsky:2000gv,Sommers:2008ji}. In this  work  we want to  make a more quantitative   and detailed  study on how to interpret the results on the clustering.   The nature  of the  particles   in the UHECR remains  a central open questions in the field. In this  work we   will assume that  most of these particles are protons. This   hypothesis  is consistent  with  the existing  data (if one takes into account the  systematic  uncertainties on  the modeling of hadronic showers) and it is also  favored  in most theoretical models. Here this   assumption is also  made to  allow a detailed quantitative description of particle  propagation in  intergalactic space, which is  essential for the problem we are considering. It is  possible  and  reasonably straightforward to generalize the  discussion to the case of nuclei.  All cosmic ray sources are  unavoidably also sources of neutrinos \\cite{Gaisser:1994yf,Learned:2000sw} because some   fraction of  a population of relativistic hadrons  will necessarily interact   with   ordinary matter or radiation fields  targets   inside or near their acceleration site creating pions and other  weakly decaying particles   that can (chain) decay into  neutrinos. Viceversa,  the emission of neutrinos  imply the existence of  relativistic hadrons, and therefore the  acceleration of cosmic rays. The  catalogues of the CR and neutrino sources  are in principle  identical. In practice   the situation could more  complicated, because  the  energy range studied by  the CR  and neutrino  telescopes   differ by several  orders of magnitude,  and the  relation  between the  neutrino and cosmic   ray emission  from a source can  vary significantly  depending on the structure of the source, and it is certainly possible  to  have  bright  CR sources that emit a relatively small  amount of neutrinos, and  viceversa. Nonetheless, it is  reasonable to expect  that the   brighest extragalactic sources of both  CR  and neutrinos  will coincide, and it is interesting to discuss in parallel these new  astronomies.  Also  in the case of neutrinos  one  can  measure an  inclusive  flux, that sums  the  contributions from all sources in the universe,   while the brightest  and most powerful  neutrino sources   should  be identifiable as  individual objects.  Therefore also  for neutrinos the ``clustering'' of the detected events is  a useful  method of study. An important  advantage is that   for neutrinos linear propagation  is   certain,  but one has to deal with the existence  of   the foreground of atmospheric  neutrinos. The  theoretical  framework   needed to discuss the clustering of neutrino  events is  essentially identical  to the  one  needed for  protons. A  quantitative difference is  that for  neutrinos  extragalactic  space is perfectly  transparent, and the  inclusive flux  receives most of its contribution from  very distant and very faint sources, and  the fraction of this  flux that can be ``resolved'' in     the contribution of identified  sources is likely to be   small. More that the interpretation of the  existing data, the goal of this  work is the development of some general analysis instruments  that can be used in  the study of future, higher statistics results.  This  work is organized as follows: in the next  section we   make a preliminary discussion of the so called ``Olbers Paradox''. The discussion of this  celebrated puzzle allows to introduce  the  key concepts needed  to compute the signal from the  ensemble of all sources in  the universe. In section~3  we  collect the results on particle propagation in intergalactic space that are needed in the following. Section~4 discusses   how the flux   received   from  one astrophysical source depends on  its  redshift. Section~5  discusses   possible forms of the luminosity function of the proton and neutrino sources. Section~6 compute the inclusive  particle flux from the combined  emission  of all sources in  the universe. Section~7 discusses  what  part of the inclusive flux can  be  resolved in the contribution  of individual sources. Section~8  applies these  analysis  instruments to the  interpretation of the recent Auger data. Section~9  discusses  extragalactic  neutrino astronomy. Section~10  gives  some conclusions. ",
        "conclusions": ""
    },
    "0808/0808.1081_arXiv.txt": {
        "abstract": "A common belief about big-bang cosmology is that the cosmological redshift cannot be properly viewed as a Doppler shift (that is, as evidence for a recession velocity), but must be viewed in terms of the stretching of space. We argue that, contrary to this view, the most natural interpretation of the redshift is as a Doppler shift, or rather as the accumulation of many infinitesimal Doppler shifts. The stretching-of-space interpretation obscures a central idea of relativity, namely that it is always valid to choose a coordinate system that is locally Minkowskian. We show that an observed frequency shift in any spacetime can be interpreted either as a kinematic (Doppler) shift or a gravitational shift by imagining a suitable family of observers along the photon's path. In the context of the expanding universe the kinematic interpretation corresponds to a family of comoving observers and hence is more natural. ",
        "introduction": "\\label{sec:intro}  Many descriptions of big-bang cosmology declare that the observed redshift of distant galaxies is not a Doppler shift but is due to the ``stretching of space.'' The purpose of this paper is to examine the meaning of such statements and to assess their validity. We wish to make clear at the outset that we are not suggesting any doubt about either the observations or the general-relativistic equations that successfully explain them. Rather, our focus is on the interpretation: given that a photon does not arrive at the observer conveniently labeled ``Doppler shift,'' ``gravitational shift,'' or ``stretching of space,'' when can or should we apply these labels?  Arguably an enlightened cosmologist never asks this question. In the curved spacetime of general relativity, there is no unique way to compare vectors at widely separated spacetime points, and hence the notion of the relative velocity of a distant galaxy is almost meaningless. Indeed, the inability to compare vectors at different points is the definition of a curved spacetime.\\cite{baezbunn,carroll,schutz,mtw} In practice, however, the enlightened view is far from universal. The view presented by many cosmologists and astrophysicists, particularly when talking to nonspecialists, is that distant galaxies are ``really'' at rest, and that the observed redshift is a consequence of some sort of ``stretching of space,'' which is distinct from the usual kinematic Doppler shift. In these descriptions, statements that are artifacts of a particular coordinate system are presented as if they were statements about the universe, resulting in misunderstandings about the nature of spacetime in relativity.  In this paper we will show that the redshifts of distant objects in the expanding universe may be viewed as kinematic shifts due to relative velocities, and we will argue that if we are forced to interpret the redshift, this interpretation is more natural than any other.  We begin with examples of the description of the cosmological redshift in the first three introductory astronomy textbooks chosen at random from the bookshelf of one of the authors. \\begin{itemize} \\item The cosmological redshift ``is {\\it not} the same as a Doppler shift. Doppler shifts are caused by an object's {\\it motion through space}, whereas a cosmological redshift is caused by the {\\it expansion of space}.''\\cite{kaufmann} (Emphasis in original.)  \\item ``A more accurate view [than the Doppler effect] of the redshifts of galaxies is that the waves are stretched by the stretching of space they travel through \\ldots\\ If space is stretching during all the time the light is traveling, the light waves will be stretched as well.''\\cite{fraknoi}  \\item ``Astronomers often express redshifts as if they were radial velocities, but the redshifts of the galaxies are not Doppler shifts \\ldots\\ Einstein's relativistic Doppler formula applies to motion through space, so it does not apply to the recession of the galaxies.''\\cite{seeds} \\end{itemize}  More advanced textbooks often avoid this language. For instance, the books by Peacock\\cite{peacock} and Linder\\cite{linder} give particularly careful and clear descriptions of the nature of the cosmological redshift. However, statements similar to those we have cited can be found even in some advanced textbooks. For example, a leading advanced undergraduate level text states that Doppler shifts ``are produced by peculiar and not by recession velocities.''\\cite{harrison} In this paper we argue, as others have before us,\\cite{chodorowskimilne,chodorowski,peacocknotes,whiting} that statements such as these are misleading and foster misunderstandings about the nature of space and time.  In general relativity the ``stretching of space'' explanation of the redshift is quite problematic. Light is governed by Maxwell's equations (or their general relativistic generalization), which contain no ``stretching of space term'' and no information on the current size of the universe. On the contrary, one of the most important ideas of general relativity is that spacetime is always locally indistinguishable from the (non-stretching) spacetime of special relativity, which means that a photon doesn't know about the changing scale factor of the universe.\\cite{footnote1}  The emphasis in many textbooks on the stretching-of-spacetime interpretation of the cosmological redshift causes readers to take too seriously the stretching-rubber-sheet analogy for the expanding universe. For example, it is sometimes stated as if it were obvious that ``it follows that all wavelengths of the light ray are doubled'' if the scale factor doubles.\\cite{harrison} Although this statement is correct, it is not obvious. After all, solutions to the Schr\\\"odinger equation, such as the electron orbitals in the hydrogen atom, don't stretch as the universe expands, so why do solutions to Maxwell's equations?  A student presented with the stretching-of-space description of the redshift cannot be faulted for concluding, incorrectly, that hydrogen atoms, the Solar System, and the Milky Way Galaxy must all constantly ``resist the temptation'' to expand along with the universe. One way to see that this belief is in error is to consider the problem sometimes known as the ``tethered galaxy problem,''\\cite{harrisontethered,davistethered} in which a galaxy is tethered to the Milky Way, forcing the distance between the two to remain constant. When the tether is cut, does the galaxy join up with the Hubble flow and start to recede due to the expansion of the universe? The intuition that says that objects suffer from a temptation to be swept up in the expansion of the universe will lead to an affirmative answer, but the truth is the reverse: unless there is a large cosmological constant and the galaxy's distance is comparable to the Hubble length, the galaxy falls toward us.\\cite{whiting,peacocknotes} Similarly, it is commonly believed that the Solar System has a very slight tendency to expand due to the Hubble expansion (although this tendency is generally thought to be negligible in practice). Again, explicit calculation shows this belief not to be correct.\\cite{sereno,cooperstock} The tendency to expand due to the stretching of space is nonexistent, not merely negligible.  The expanding rubber sheet is quite similar to the ether of pre-relativity physics in that although it is intuitively appealing, it makes no correct testable predictions, and some incorrect ones such as the examples we have given. It therefore has no rightful place in the theory. (Some authors\\cite{barnes,francis} have argued that considerations such as these do not refute the notion that space is really expanding. We agree with the calculations in these papers but differ regarding the most useful language to use to describe the relevant phenomena.)  In one set of circumstances the proper interpretation of the redshift seems clear. When the curvature of spacetime is small over the distance and time scales traveled by a photon, it is natural to interpret the observed frequency shift as a Doppler shift. This interpretation is the reason that a police officer can give you a speeding ticket based on the reading on a radar gun. As far as we know, no one has successfully argued in traffic court that there is an ambiguity in interpreting the observed frequency shift as a Doppler shift.\\cite{footnote3} In the expanding universe, spacetime curvature is small over regions encompassing nearby objects, specifically those with $z=\\Delta\\lambda/\\lambda\\ll 1$. There should be no hesitation about calling the observed redshifts Doppler shifts in this case, just as there is none in traffic court. Surprisingly, however, many people seem to believe that the ``stretching of space'' interpretation of the redshift is the only valid one, even in this limit. We will examine the interpretation of redshifts of nearby objects more carefully in Sec.~\\ref{sec:lowz}.  Aside from low-redshift sources, there is another case in which spacetime curvature can be neglected in considering cosmological redshifts, namely low-density cosmological models. An expanding universe with density $\\Omega=0$ (often known as the Milne model\\cite{milne}) is merely the flat Minkowski spacetime of special relativity expressed in nonstandard coordinates.\\cite{footnote4} In an $\\Omega=0$ universe there are no gravitational effects at all, so any observed redshift, even of a very distant galaxy, must be a Doppler shift. Furthermore, for low but nonzero density ($\\Omega \\ll 1$), the length scale associated with spacetime curvature is much longer than the horizon distance. Hence, spacetime curvature effects (that is, gravitational effects) are weak throughout the observable volume, and the special-relativistic Doppler shift interpretation remains valid even for galaxies with arbitrarily high redshifts.  The more interesting cases are when $z$ and $\\Omega$ are not small; that is, when the source is sufficiently distant that gravitational effects are important over the photon's trajectory. The consensus is that the Doppler shift language must be eschewed in this setting. In Sec.~\\ref{sec:highz} we review a standard argument\\cite{peacock,peacocknotes} that even in this case the redshift can be interpreted as the accumulation of infinitesimal Doppler shifts along the line of sight, and we further argue that there is a natural way to interpret the redshift as a single (non-infinitesimal) Doppler shift.  A common objection to this claim is that the coordinate velocity is not related to the redshift in accordance with the special-relativistic Doppler formula.\\cite{davissuperluminal,davisconfusion} However, the velocity referred to in this claim is a mere artifact of a particular choice of time coordinate. Specifically, it is the rate of change of the distance to the object with respect to the cosmic time coordinate, as measured at the present cosmic time. This coordinate velocity is an unnatural quantity to discuss, because it depends on data outside of the observer's light cone. The more natural velocity is the velocity of the object at the time it crossed our past light cone, relative to us today. This velocity is also a coordinate-dependent concept, but as we will show in Sec.~\\ref{sec:highz}, the most natural way to specify it operationally\\cite{synge,narlikar} leads to a result that is consistent with the special-relativistic Doppler formula.  In Sec.~\\ref{sec:observers} we widen our focus to consider frequency shifts in arbitrary curved spacetimes. In any curved spacetime the observed frequency shift in a photon can be interpreted as either a kinematic effect (a Doppler shift) or as a gravitational shift. The two interpretations arise from different choices of coordinates, or equivalently from imagining different families of observers along the photon's path. We will describe this construction explicitly, and show that the comoving observers who are usually used to describe phenomena in the expanding universe are the ones that correspond to the Doppler shift interpretation. ",
        "conclusions": ""
    },
    "0808/0808.2611_arXiv.txt": {
        "abstract": "We present new evolution sequences for very low mass stars, brown dwarfs and giant planets and use them to explore a variety of influences on the evolution of these objects.  While the predicted adiabatic evolution of luminosity with time is very similar to results of previous work, the remaining disagreements reveal the magnitude of current uncertainty in brown dwarf evolution theory. We discuss the sources of  those differences and argue for the importance of the surface boundary condition provided by atmosphere models including clouds.  The L- to T-type ultracool dwarf transition can be accommodated within the \\citet{am01} cloud model by varying the cloud sedimentation parameter. We develop a simple model for the evolution across the L/T transition. By combining the evolution calculation and our atmosphere models, we generate colors and magnitudes of synthetic populations of ultracool dwarfs in the field and in galactic clusters. We focus on near infrared color-magnitude diagrams (CMDs) and on the nature of the ``second parameter'' that is responsible for the scatter of colors along the $\\teff$ sequence.  Instead of a single second parameter  we find that variations in metallicity and cloud parameters, unresolved binaries and possibly a relatively young population all play a role in defining the spread of brown dwarfs along the cooling sequence.  We also find that the transition from cloudy L dwarfs to cloudless T dwarfs slows down the evolution and causes a pile up of substellar objects in the transition region, in contradiction with previous studies. The same model is applied to the Pleiades brown dwarf sequence with less success, however.  Taken at face value, the present Pleiades data suggest that the L/T transition occurs at lower $\\teff$ for lower gravity objects, such as those found in young galactic clusters.  The simulated populations of brown dwarfs also reveal that the phase of deuterium burning produces a distinctive feature in CMDs that should be detectable in $\\sim 50$--100$\\,$Myr old clusters. ",
        "introduction": "There are now approximately 450 L dwarfs and 100 T dwarfs known (see \\citet{kirk05} for a review of these spectral classes). They span effective temperatures from about 2400 to 700 K and exhibit a range of gravities, metallicities, and atmospheric condensate contents.  After more than a decade of intense study, the modeling of the complex atmospheres and synthetic spectra of brown dwarfs has reached a rather high degree of sophistication, including the chemistry of a very large number of gas and condensate species \\citep{allard01,lod02}, increasingly complete molecular opacity databases \\citep{fml08,sb07}, extreme resonance line broadening \\citep{bms00,aak05}, and particulate cloud models \\citep{allard01,am01,tsu02,helling08}.  While a few conspicuous problems remain, the synthetic spectra and colors reproduce the observations fairly well and the determination of the basic astrophysical properties of brown dwarfs has begun in earnest. The full astrophysical benefit of synthetic spectra and colors is obtained when atmosphere calculations are coupled with evolution models that provide the surface atmospheric parameters $(\\teff,g)$ as a function of mass and age. The time evolution of the spectrum and colors, as well as absolute fluxes can be computed and directly compared with observations to estimate astrophysical parameters that are not easily amenable to direct observation.  To enable such comparisons using our own model atmosphere effort and to pursue a more complete analysis of spectroscopic and photometric data, we have developed a code to compute evolution sequences of low mass stars, brown dwarfs and giant planets.  These evolution sequences have already been applied extensively to the analysis of brown dwarf observations \\citep{roellig04,saumon06,saumon07,leggett07y, leggett07p,cushing08} but have not been discussed in any detail. Here, we describe the input physics and assumptions of our particular approach as well as the unique aspects of our atmospheric boundary condition. The evolution of isolated brown dwarfs---a relatively simple case of stellar evolution---has been studied extensively for the past 20 years and is well understood \\citep{dm85,nrj86,bhl89,bl93,bur97,cbah00a,cbah00b,bcah02}. Our evolution model is similar to the more recent work and the result are very similar to previous work. A detailed comparison with other published calculations of the evolution of brown dwarfs reveals small differences that we quantify and, to the extent possible, attribute to the different assumptions and approximations in each model. We highlight the application of the evolution sequences to the calculation of synthetic near infrared color-magnitude diagrams (CMDs) to explore the nature of the ``second parameter'' responsible for the spread of observed objects around the main trends along the L and T spectral sequences. We develop a simple parametric model for the L/T transition to reproduce with good success the CMD of field brown dwarfs.  We apply the same model to synthesize the population of the Pleiades cluster (110$\\,$Myr) and compare with the latest deep survey data for this galactic cluster.  We discuss several potentially observable features in near-infrared CMDs that would illuminate the evolution of brown dwarfs as well as the nature of L/T transition. ",
        "conclusions": "Our calculation of the evolution of very low mass stars, brown dwarfs and planetary mass objects produces models that are quantitatively in very good agreement with published calculations.  A detailed comparison shows some systematic differences that can be attributed to conductive energy transport, different choices of composition for the interior, and for the initial state but the primary source of discrepancy is the surface boundary condition.  The cloudless and cloudy model atmospheres from which we extract the surface boundary condition have been validated with extensive comparisons with spectroscopic and photometric data.  Thus, our boundary condition is quite realistic and we do not expect that the foreseeable improvements in atmosphere models will have much effect on our modeled evolution of brown dwarfs. By using the boundary condition from our atmosphere models, we can compute self-consistently the evolution, absolute fluxes, absolute magnitudes, and colors for a variety of cloud properties, metallicities and eventually including vertical mixing \\citep{saumon06}. Only the evolution across the L/T transition region cannot be modeled in this self-consistent fashion because a transition cloud model is not yet available.  We have developed a simple model for the cooling and color evolution of brown dwarfs across the L/T transition. This hybrid model predicts an excess of brown dwarfs in the $\\teff$ range of the transition by about a factor of 2 compared to purely cloudy or cloudless evolution. We have applied this hybrid evolution model, combined with the near-infrared magnitudes predicted by our atmosphere models, to generate synthetic CMDs that can be compared with samples of brown dwarfs in the field and in galactic clusters.  Our primary focus is the ``second parameter'' responsible for the dispersion about the brown dwarf sequence in the CMD. Population synthesis is a potentially powerful tool but the results can be quite sensitive to the input assumptions for the IMF, SFR, metallicity, and the binary frequency and mass ratio distribution. Our knowledge of the brown dwarf population in the solar neighborhood cannot yet provide these inputs {\\it a priori}.  The relatively small number of brown dwarfs with known parallax and the heterogeneous nature of the sample imply that the observed distribution of field substellar dwarfs in the CMD is not yet fully characterized.  Nevertheless, both observations and models have reached a stage where general trends in the ``second parameter'' along the L-T spectral sequence can be interpreted.  We find that for our fiducial assumptions (power law IMF with $\\alpha=1$, constant SFR over the past 10$\\,$Gyr, single brown dwarfs only and [M/H]=0), the hybrid sequence reproduces the overall sequence from late M through late T dwarfs rather well, but not the dispersion along the sequence.  Based on the near-infrared CMDs, we find that the L/T transition occurs between $\\teff \\sim 1400-1200\\,$K in field brown dwarfs, in agreement with previous estimates. While a transition over such a narrow range of $\\teff$ appears to be ``fast'' considering the rather dramatic change in $JHK$ colors across the transition, these values of $\\teff$ correspond to ages of 2 and 4$\\,$Gyr for a 0.06$\\,M_\\odot$ brown dwarf. The duration of the transition decreases rapidly with mass however, lasting only 0.15$\\,$Gyr for a 0.03$\\,M_\\odot$ brown dwarf.  Better agreement can be obtained from late M to late L spectral types if the population is younger, such as with a constant SFR that started only 5$\\,$Gyr ago, by including binaries, or assuming that there is a wider range of cloud properties for later L spectral types. For a fixed metallicity, all simulations predict that the distribution of brown dwarfs in the CMD will have a sharp edge formed by old brown dwarfs of all masses ($\\gtrsim 3\\,$Gyr). This edge is on the blue side of the distribution for $M_K \\lesssim 12.5$ and to the red side after the $J-K$ color of the sequence turns over, corresponding to $M_K \\gtrsim 13$.  This feature is not visible in the data, however, most likely because it is blurred by variations in metallicity within the sample.  We are not able to include metallicity variations in simulations of cloudy brown dwarfs, except in a very approximate way (Fig. \\ref{fig:metal_f2}).  We find that it could be a significant contributor to the second parameter.  Detailed spectral analysis of brown dwarfs with unusual $J-K$ colors for their spectral types have more extreme cloud parameters \\citep{burgasser08,cushing08,stephens08}, which strongly suggest that cloudiness is the second parameter. Spectral analysis with models of non-solar metallicities have barely begun, however, so it would be premature to attribute the dispersion along the L sequence entirely to cloud characteristics. On the other hand, a simulation of cloudless models with an empirical metallicity distribution shows a good match to the dispersion of the late T dwarfs. For those coolest dwarfs, we find that the second parameter is a combination of metallicity variations (which dominate) and binaries.  Gravity is not important as the with of the distribution is not affected by the choice of SFR or IMF. To summarize, we find that there is no single second parameter that accounts for the dispersion of brown dwarfs around the sequence seen in the CMD. A young age distribution, a range of metallicities and cloud properties as well as binaries all contribute to the dispersion.  The challenge will be to untangle their contributions.  The hybrid model fares somewhat worse when compared to the much younger brown dwarf population of the Pleiades. If the two faintest Pleiads reported are indeed T dwarf members of the cluster, then they provide strong evidence that the L/T transition occurs at lower $\\teff$ in lower gravity objects (i.e. younger or less massive). Finally, isochrones in CMDs clearly reveal the phase of deuterium burning at young ages, a feature that should be observable in young clusters with ages between of $\\sim 50-100\\,$Myr.  At this time, the modest size of the sample of L and T dwarfs with known parallax and the lingering problems in modeling the atmospheres of cloudless and cloudy brown dwarfs restrict how much we can learn from the study of CMDs.  Model limitations will eventually be overcome as new moelcular line lists are being developed for key molecules and cloud models become more sophisticated. We anticipate a rich harvest of brown dwarf parallaxes from the volume limited solar neighborhood census component of the Panoramic Survey Telescope \\& Rapid Response System (Pan-STARRS) \\footnote{\\tt http://pan-starrs.ifa.hawaii.edu/project/reviews/PreCoDR/documents/scienceproposals/sol.pdf}. Color-magnitude diagrams are a potentially powerful tool for the study of brown dwarf evolution and of the L/T transition. Statistical comparisons with synthetic populations in two-dimensional parameter space will become an important complement to the detailed studies of the spectra of individual transition objects and of brown dwarf binaries."
    },
    "0808/0808.0614_arXiv.txt": {
        "abstract": "Fermionic condensate and the vacuum expectation values of the energy-momentum tensor are investigated for twisted and untwisted massive spinor fields in higher-dimensional de Sitter spacetime with toroidally compactified spatial dimensions. The expectation values are presented in the form of the sum of corresponding quantities in the uncompactified de Sitter spacetime and the parts induced by non-trivial topology. The latter are finite and renormalizations are needed for the first parts only. Closed formulae are derived for the renormalized fermionic vacuum densities in uncompactified odd-dimensional de Sitter spacetimes. It is shown that, unlike to the case of 4-dimensional spacetime, for large values of the mass, these densities are exponentially suppressed. Asymptotic behavior of the topological parts in the expectation values are investigated in the early and late stages of the cosmological expansion. When the comoving lengths of compactified dimensions are much smaller than the de Sitter curvature radius, the leading term in the topological parts coincide with the corresponding quantities for a massless fermionic field and are conformally related to the corresponding flat spacetime results. In this limit the topological parts dominate the uncompactified de Sitter part and the back-reaction effects should be taken into account. In the opposite limit, for a massive field the asymptotic behavior of the topological parts is damping oscillatory. ",
        "introduction": "De Sitter (dS) spacetime is one of the simplest and most interesting spacetimes allowed by general relativity. Quantum field theory in this background has been extensively studied during the past two decades. Much of early interest to dS spacetime was motivated by the questions related to the quantization of fields propagating on curved backgrounds. This spacetime has a high degree of symmetry and numerous physical problems are exactly solvable on this background. The importance of this theoretical work increased by the appearance of \\ the inflationary cosmology scenario \\cite% {Lind90}. In most inflationary models, an approximately dS spacetime is employed to solve a number of problems in standard cosmology. During an inflationary epoch, quantum fluctuations in the inflaton field introduce inhomogeneities and may affect the transition toward the true vacuum. These fluctuations play a central role in the generation of cosmic structures from inflation. More recently astronomical observations of high redshift supernovae, galaxy clusters and cosmic microwave background \\cite{Ries07} indicate that at the present epoch, the Universe is accelerating and can be well approximated by a world with a positive cosmological constant. If the Universe would accelerate indefinitely, the standard cosmology would lead to an asymptotic dS universe. Hence, the investigation of physical effects in dS spacetime is important for understanding both the early Universe and its future.  Many of high energy theories of fundamental physics are formulated in higher-dimensional spacetimes. In particular, the idea of extra dimensions has been extensively used in supergravity and superstring theories. It is commonly assumed that the extra dimensions are compactified. From an inflationary point of view, universes with compact spatial dimensions, under certain conditions, should be considered a rule rather than an exception \\cite{Lind04}. The models of a compact universe with non-trivial topology may play important roles by providing proper initial conditions for inflation. As it was argued in Refs. \\cite{McIn04}, there is no reason to believe that the version of dS spacetime which may emerge from string theory, will necessarily be the most familiar version with symmetry group $% O(1,4)$ and there are many different topological spaces which can accept the dS metric locally. There are many reasons to expect that in string theory the most natural topology for the universe is that of a flat compact three-manifold. The quantum creation of the universe having toroidal spatial topology is discussed in \\cite{Zeld84} and in references \\cite{Gonc85} within the framework of various supergravity theories.  The compactification of spatial dimensions leads to a number of interesting quantum field theoretical effects which include instabilities in interacting field theories \\cite{Ford80a}, topological mass generation \\cite{Ford79} and symmetry breaking \\cite{Toms80b}. In the case of non-trivial topology, the boundary conditions imposed on fields give rise to the modification of the spectrum for vacuum fluctuations and, as a result, to the Casimir-type contributions in the vacuum expectation values of physical observables (for the topological Casimir effect and its role in cosmology see \\cite% {Grib94,Most97} and references therein). In the Kaluza-Klein-type models, the Casimir effect has been used as a stabilization mechanism for moduli fields which parametrize the size and the shape of the extra dimensions. The Casimir energy can also serve as a model for dark energy needed for the explanation of the present accelerated expansion of the universe (see \\cite% {Milt03} and references therein). One-loop quantum effects for various spin fields on the background of dS spacetime, have been discussed by several authors (see, for instance, \\cite{Cher68}-\\cite{Birr82} and references therein). The effects of the toroidal compactification of spatial dimensions in dS spacetime on the properties of quantum vacuum for a scalar field with general curvature coupling parameter are investigated in Refs. \\cite% {Saha07,Bell08} (for quantum effects in braneworld models with dS spaces and in higher-dimensional brane models with compact internal spaces see, for instance, Refs. \\cite{dSbrane,Flac03}). The one-loop quantum effects for a fermionic field on background of 4-dimensional dS spacetime with spatial topology $\\mathrm{R}^{p}\\times (\\mathrm{S}^{1})^{q}$ are studied in \\cite% {Saha08}.  In the present paper, we investigate one-loop quantum effects arising from vacuum fluctuations of a fermionic field on background of higher-dimensional dS spacetime with toroidally compactified spatial dimensions. The important quantities that characterize the quantum fluctuations during the dS expansion are the fermionic condensate and the expectation value of the energy-momentum tensor. In the next section, by using the dimensional regularization procedure, we evaluate these quantities in uncompactified odd-dimensional dS spacetimes. The plane wave fermionic eigenfunctions in $% (D+1)$-dimensional dS spacetime with an arbitrary number of toroidally compactified dimensions are constructed in section \\ref{sec:EigFunc}. In section \\ref{sec:FermCond} these eigenfunctions are used for the evaluation of the fermionic condensate in both cases of the fields with periodicity and antiperiodicity conditions along compactified dimensions. The behavior of these quantities are investigated in asymptotic regions of the parameters. The topological parts in the vacuum expectation values of the energy-momentum tensor are investigated in section \\ref{sec:EMT}. In the last section we summarize the main results of the paper. ",
        "conclusions": "\\label{sec:Conc}  In the present paper we have investigated the fermionic condensate and the VEV of the energy-momentum tensor for a massive fermionic field in higher-dimensional dS spacetime with toroidally compactified spatial dimensions. In Section \\ref{sec:UncompdS} we have considered the corresponding quantities in uncompactified odd-dimensional dS spacetime assuming that the field is prepared in the Bunch-Davies vacuum state. The renormalization is done by using the dimensional regularization procedure. Closed expressions, formulae (\\ref{FermConddSRen}) and (\\ref{TkldSren}), are derived for the renormalized fermionic condensate and the VEV of the energy-momentum tensor respectively. For large values of the mass these quantities are exponentially suppressed. Note that in even-dimensional dS spacetime for large mass the suppression is power-law.  Further, we have investigated one-loop quantum effects on the fermionic vacuum induced by the non-trivial topology of spatial dimensions. Specifically, we have considered the dS spacetime with toroidally compactified dimensions having the spatial topology $\\mathrm{R}^{p}\\times (% \\mathrm{S}^{1})^{q}$. For the evaluation of the vacuum densities, the mode-summation procedure is employed. In this procedure we need to know the corresponding eigenspinors satisfying appropriate boundary conditions along the compactified dimensions. These eigenspinors are constructed in section % \\ref{sec:EigFunc} for both fields obeying periodicity and antiperiodicity boundary conditions. By using these eigenfunctions and applying to the mode-sums the Abel-Plana formula, the VEVs for the spatial topology $\\mathrm{% R}^{p}\\times (\\mathrm{S}^{1})^{q}$ are presented in the form of the sum of the corresponding quantity in the topology $\\mathrm{R}^{p+1}\\times (\\mathrm{S% }^{1})^{q-1}$ and of the part which is induced by the compactness of $(p+1)$% th dimension. For fields obeying periodicity conditions, the topological parts are given by formulae (\\ref{DeltCond2}) and (\\ref{TopTll}) for fermionic condensate and energy-momentum tensor, respectively. The corresponding formulae for the field with antiperiodicity conditions are obtained from those for the field obeying periodicity conditions inserting the factor $(-1)^{n}$ in the summation over $n$ and replacing the definition for $k_{\\mathbf{n}_{q-1}}^{2}$ by (\\ref{knq-1Tw}).  The topological parts are finite and the renormalization procedure is needed only for the uncompactified dS spacetime. These parts are time-dependent and break the dS symmetry. The corresponding vacuum stresses along the uncompactified dimensions coincide with the energy density and, hence, in the uncompactified subspace the equation of state for the topological part of the energy-momentum tensor is of the cosmological constant type. For a massless fermionic field the problem under consideration is conformally related to the corresponding problem in the Minkowski spacetime with spatial topology $\\mathrm{R}^{p}\\times (\\mathrm{S}^{1})^{q}$ and the topological part of the fermionic condensate vanishes. For the VEV of the energy-momentum tensor we have the standard relation $\\langle T_{k}^{l}\\rangle _{c}=a^{-(D+1)}(\\eta )\\langle T_{k}^{l}\\rangle _{c}^{% \\mathrm{(M)}}$ between the topological contributions.  For a massive fermionic field, in the limit when the comoving length of a compactified dimension is much smaller than the dS curvature radius, the topological part in the VEV of the energy-momentum tensor coincides with the corresponding quantity for a massless field and is conformally related to the VEV in toroidally compactified Minkowski spacetime. In particular, the topological part in the vacuum energy density is positive for an untwisted fermionic field. This limit corresponds to the early stages of the cosmological evolution and the topological parts dominate over the uncompactified dS parts. At these stages the back-reaction effects of the topological terms are important and these effects can essentially change the dynamics of the model. In the opposite limit, when the comoving lengths of the compactified dimensions are large with respect to the dS curvature radius, in the case of a massive field the asymptotic behavior of the topological parts are oscillatory damping for both fermionic condensate and the energy-momentum tensor and their respectively leading term are given by formulae (\\ref{CondLate}) and (\\ref{TllSmall}). These formulae describe the behavior of the topological parts in the late stages of the cosmological expansion. As the corresponding uncompactified dS parts are time-independent, we have similar oscillations in the total VEVs as well. Note that this type of oscillatory behavior is absent for a massless fermionic field."
    },
    "0808/0808.0939_arXiv.txt": {
        "abstract": "A general phenomenological theory is presented for the phase behavior of ferromagnetic superconductors with spin-triplet electron Cooper pairing. The theory accounts in detail for the temperature-pressure phase diagram of ZrZn$_2$, while the main features of the diagram for UGe$_2$ are also described. Quantitative criteria are deduced for the U-type (type I) and Zr-type (type II) unconventional ferromagnetic superconductors with spin-triplet Cooper electron pairing. Some basic properties of quantum phase transitions are also elucidated. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.0422_arXiv.txt": {
        "abstract": "{Below 1 mHz, the power spectrum of  helioseismic velocity measurements is dominated by the spectrum of convective motions (granulation and supergranulation) making it difficult to detect the low-order acoustic modes and the gravity modes.} {We want to better understand the behavior of solar granulation as a function of the observing height in the solar atmosphere and with magnetic activity during solar cycle 23. } {We analyze the Power Spectral Density (PSD) of eleven years of GOLF/SOHO velocity-time series using a Harvey-type model to characterize the properties of the convective motions in the solar oscillation power spectrum. We study then the evolution of the granulation with the altitude in the solar atmosphere and with the solar activity.} {First, we show that the traditional use of a lorentzian profile to fit the envelope of the $p$ modes is not well suitable for GOLF data. Indeed, to properly model the solar spectrum, we need a second lorentzian profile. Second, we show that the granulation clearly evolves with the height in the photosphere but does not present any significant variation with the activity cycle.} {} ",
        "introduction": "\\label{Intro}  The photosphere, the visible layer of the Sun, is the location where the energy transport previously dominated by convection and turbulence is largely done by radiation with an optical depth of $2/3$. The gas is visible in the form of granules, that penetrate inside the stable photosphere. These granules, as well as other larger structures like the mesogranules or the supergranules, are the manifestation of  the different spatial scales of the convective motions occurring in this region of the Sun \\citep{Zahn87,Roudier91,Espagnet93}.  The study of the granulation is particularly important in helioseismology because, on the one hand, it excites the so-called 5-min oscillations, i.e. the acoustic (p) modes, and, on the other hand, it dominates the power spectrum at low frequencies  preventing the detection of low-order p modes. Indeed, in the case of velocity measurements, the lower detection limit of acoustic modes is established around 1 mHz \\citep{Garcia01,Garcia04a,Broomhall07}. To progress in the detection of such modes as well as to increase the detection probability of gravity modes  \\citep{Appourchaux00,Gabriel02,Turck04,Garcia07,Mathur07} we need to better characterize the properties of the granulation in order to reduce, if possible, their impact on the helioseismic  measurements. The coming GOLF-NG instrument will soon address this problem \\citep{Turck06}. This new-generation instrument should improve the signal-to-noise (S/N) ratio of low-frequency modes by measuring the Doppler velocity at different heights in the solar atmosphere. It would thus benefit from the reduction in the coherence of the granulation with the atmospheric altitude \\citep{Garcia04b}.  In this paper, we analyze the  Power Spectral Density (PSD) of velocity time sub-series from the GOLF\\footnote{Global Oscillation at Low Frequencies \\citep{Gabriel95}} instrument on board SOHO\\footnote{SOlar and Heliospheric Observatory \\citep{Domingo95}} which is a solar disk integrated resonant spectrometer. With such an instrument, we can already study the mean behaviour of the solar granulation in some range of the atmosphere.  So this work  comes in complement to  the study of \\citet{Espagnet95} where the authors found that the photosphere is highly structured with two distinct layers below and above about 90 km. With GOLF, and the technique used in \\citet{Jimenez07}, we are able to study, without spatial resolution, a region located between 250 up to 550 km above the photosphere.  We show that the granulation evolves with the height in the photosphere and that the granules tend to have shorter lifetimes with a weaker velocity when higher in the atmosphere.  Section \\ref{Analysis} is devoted to the data analysis, with first a brief summary of the velocity calibration procedure of the GOLF signal and secondly a description of the fitting procedure of the power spectra in more details. Section \\ref{Results} is dedicated to a detailed study of the granulation motions and its evolution with time during the solar cycle. Finally, last section (Sect. \\ref{Discussion}) will emphasize the main results of this paper and will anticipate on future incoming works. ",
        "conclusions": "\\label{Discussion}  In this paper we have investigated the vertical structure and time evolution of the solar granulation by means of a novel methodology based on the analysis of the full-disk Sun-as-a-star Doppler velocity observations. Thus we have been able to study the vertical velocity fluctuations and lifetimes of the solar granulation. We have shown that the GOLF PSD can be correctly characterized by our model, a Harvey function with two Lorentzian profiles.  This work extends the study of \\citet{Espagnet95} where they found that the photosphere is highly structured with two distinct layers below and above about 90 km. With GOLF, we  study the photosphere above $\\approx$ 280 km and we showed that granules tend to live longer with a weaker velocity when higher in the atmosphere.   Following the results of \\citet{Title89}, there is a strong correlation between the granule sizes and lifetime. Therefore, we can conclude that larger granules reach the top of the photosphere, while penetration of small granules decreases with the size. To be more precise, as Figure \\ref{fig3} shows a lifetime of about 400-550 s for granules between 250 and 550 km, we can estimate from the article of \\citet{Title89} and their figure 21 that the granules in these altitudes have a lifetime-average size of about 1.2 arcsec and a maximum size of about 1.4-1.5 arcsec.  We have found that the granulation rms velocities ($\\sigma_{gr}$) are between 0.28 and 0.4 $ms^{-1}$. These values are very different from those obtained from high-resolution measurements which are about 1 $kms^{-1}$. The difference in spatial resolution is the most likely cause of this difference of three orders of magnitude.  However we did not put in evidence a change with the solar activity. This is consistent with the result of \\citet{Jimenez03} that the solar cycle effects are very small compared to the change in the observing height in the photosphere due to the orbital motion. It could be due to the fact that GOLF observes in the more stable part of the sodium lines. However if we do not find any significant variation with the cycle, an other study \\citep{Muller07} found a cyclic variation of the contrast of the granules, nearly in phase with the solar cycle, the contrast being smaller at the periods of solar maximum, but no corresponding variation in the scale.  This work opens two perspectives:  (1) to better understand the evolution of the whole solar convective background during the activity cycle, and specially the evolution of the acoustic-mode envelope presently characterized by the presence of two Lorentzians. We are likely to think that the excess power characterized by the second Lorentzian corresponds to the presence of chromospheric modes, (2) to extend this study to our new instrument GOLF-NG observation. It will cover a larger part of the solar atmosphere thanks to the 8 extractions of the Doppler velocity with a proper determination of their location due to the measurement of both wings and the use of only the D1 line  \\citep{Jimenez07, Turck06, Turck08}."
    },
    "0808/0808.0570_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1} The original classification of elliptical galaxies, which goes back to Edwin Hubble, is based on their apparent ellipticity. The projected ellipticity, however, is also a function of inclination, and not only of the intrinsic ellipticity. \\cite{B88} showed that the isophotes of elliptical galaxies deviate significantly from perfect ellipses. The deviations can be quantified by the $a_4$-parameter, where negative $a_4$ values signify boxy deviations and positive $a_4$ values discy deviations from a perfect ellipse. But it was also found that ellipticals which have boxy isophotes, are also X-ray bright, more luminous, rotate slowly and have cored inner surface density slopes, while discy ellipticals are less luminous, rotate fast and have power-law surface density profiles \\cite{BDM88}. This led \\cite{KB96} to revise the Hubble classification of ellipticals using isophotal shape instead of ellipticity as a galaxy family membership criterium. One of the most successful models of elliptical galaxy formation is the merging of disc galaxies. \\cite{BA98} proposed that the isophotal shape depends on the violence of the merging and subsequently \\cite{N99} showed that mergers of galaxies with comparable mass form more likely boxy ellipticals, while mergers of galaxies with differing mass form discy ellipticals. However, disc-disc mergers do not form a perfect dichtomy \\cite{N03}, while elliptical-elliptical mergers seem to disconnect merging ratio and isophotal shape, i.e. they are always boxy \\cite{NKB}. In contrast to this gas seems to play an important role in fast-rotating, discy ellipticals \\cite{N06},\\cite{J07}. However, recently the SAURON sample, a survey of 48 elliptical and lenticular galaxies, \\cite{E04} showed that isophotal shapes are not well connected to rotation, i.e. fast rotating galaxies can have boxy isophotes \\cite{E07}. In the following we want to shed some light  on the connection of photometric and kinematic parameters in N-body merger remnants and how they relate to the internal orbital structure. ",
        "conclusions": "The trends seen here are typical for disc-disc merger remnants, while they are probably representative for low to intermediate luminosity elliptical galaxies, they cannot explain the whole population of elliptical galaxies. The analysis of the orbital fine structure should be extended to a broader range of formation mechanisms, e.g. E/S0 galaxies with boxy isophotes, will certainly have bars \\cite{A05} which have not been covered by our present analysis. Dry merging will have certainly played a role in the formation of the most massive stellar systems in the universe and will lead to round or boxy systems \\cite{NKB},\\cite{KB05}, while gas physics is important to form remnants with realistic LOSVDs \\cite{N06},\\cite{J07}. Monolithic collapse can also form triaxial self-gravitating systems which can have higher fractions of semi-stochastic orbits than our remnants \\cite{Aqui07}. Finally binary merger remnants can be compared to elliptical galaxies which formed in cosmological simulations and will probably have very different orbital structures\\cite{N07}."
    },
    "0808/0808.0093_arXiv.txt": {
        "abstract": "{Despite many studies of star formation in spiral galaxies, a complete and coherent understanding of the physical processes that regulate the birth of stars has not yet been achieved, nor has unanimous consensus been reached, despite the many attempts, on the effects of the environment on the star formation in galaxy members of rich clusters.} {We focus on the local and global Schmidt law and we investigate how cluster galaxies have their star formation activity perturbed.} {We collect multifrequency imaging for a sample of spiral galaxies, members of the Virgo cluster and of the local field; we compute the surface density profiles for the young and for the bulk of the stellar components,  for the molecular and for the atomic gas.} {Our analysis shows that the bulk of the star formation correlates with the molecular gas, but the atomic gas is important or even crucial in supporting the star formation activity in the outer part of the disks. Moreover, we show that cluster members that suffer from a moderate HI removal have their molecular component and their SFR quenched, while highly perturbed galaxies show an additional truncation in their star forming disks.} {Our results are consistent with a model in which the atomic hydrogen is the fundamental fuel for the star formation, either directly or indirectly through the molecular phase; therefore galaxies whose HI reservoirs have been depleted suffer from starvation or even from truncation of their star formation activity.} ",
        "introduction": "A satisfactory understanding of the physical processes governing the formation of giant molecular clouds (GMC) from primeval atomic hydrogen (HI) and their instability leading to the formation of stars is still far from achieved. This is not surprising due to the complex mix of hydrodynamical and gravitational physics involved in the process of star formation. Beside the theoretical difficulties, what surprisingly limits our present understanding of this issue is the lack of a robust phenomenology of star formation in nearby galaxies. Despite the many recipes that have been proposed\\footnote{See for instance the dynamically modified Schmidt law in \\citet{won02}}, the simplest and most commonly used parametrization is the Schmidt law \\citep{sch59} that establishes a correlation between the star formation rate density $\\rho_{sfr}$ and the gas density $\\rho_{gas}$: \\begin{equation} \\rho_{sfr}=a\\rho_{gas}^\\alpha\\:; \\end{equation} derived for our Galaxy, the validity of this law has been tested also for external galaxies \\citep[e.g.][]{ken83}, for which a similar relation, usually written in terms of surface densities, holds: \\begin{equation}\\label{Slaw} \\Sigma_{sfr}=A\\Sigma_{gas}^n\\:. \\end{equation} Many studies of the global Schmidt law, i.e. the relationship between the disk--averaged star formation rate (SFR) and the total gas content, have been carried out on large samples of galaxies. A milestone paper on this issue is by \\citet{ken98}, who studied the connection between the integrated SFR surface density and the integrated gas surface density in a sample of normal and starburst galaxies. Using integrated quantities, he found that the same Schmidt law with an index $n=1.4$ holds for both normal and starburst galaxies, covering a range of 5 dex in the gas density and over 6 dex in the SFR; he also found that this relation is mainly driven by the correlation between the SFR and the HI, while there is a poorer correlation between the SFR and the H$_2$, significant only when massive (L$_B>10^{10}$ L$_\\odot$) objects are considered. Owing to the fact that star formation takes place in molecular regions on parsec scales, this result was somewhat unexpected, but Kennicutt argued that the poor knowledge of the CO--to--H$_2$ conversion factor might be responsible for the scatter in the H$_2$/SFR correlation. Using seven galaxies from the BIMA survey \\citep{hel03}, \\citet{won02} studied the applicability of the Schmidt law on local scales, i.e. through a comparison between the SFR and the gas surface density profiles. They found  that a Schmidt law between the SFR and the gas content holds also on local scales and, even if the molecular gas content alone does not always control the SFR, they concluded that the correlation found for the total gas is entirely driven by the molecular component. Consistent results were found by \\citet{boi03}, who used low resolution CO data but included an $X$ factor that varies as a function of the metallicity. Another open issue is to asses the influence of the environment on the star formation. Galaxies in clusters suffer from a variety of environmental perturbations, mainly due to gravitational interactions between  galaxies themselves or between a galaxy and the cluster potential well \\citep[see the review by][]{bos06} or to various hydrodynamical interactions between the ISM and the intergalactic medium (IGM). Tidal interactions between galaxies or between individual galaxies and the cluster potential \\citep{mer84, byr90} produce both the removal of the outer and looser components and the gas infall towards the galaxy center; harassment \\citep{moo96} can induce sinking of gas towards the galaxy center and can shape the stellar profiles; the ram pressure stripping \\citep{gun72} or the viscous stripping \\citep{nul82} can effectively remove the gas components, mainly from the outer part of the disks. Furthermore, a combination of gravitational and hydrodynamic effects can cause galaxy starvation or strangulation \\citep{lar80,bek02}. According to starvation, the gas that feeds the star formation in the local universe probably comes from the infall of an extended gas reservoir, therefore the effect of removing the outer galaxy halo would be that of preventing further gas infall; as a consequence,  the star formation exhausts the available gas, quenching further activity. The literature about the classification and the effects produced by the environment on the evolution of cluster galaxies is very rich \\citep[e.g.][]{bos06}, but, despite many studies on nearby and distant clusters, it is not clear what the dominant processes are, nor how the star formation activity is affected \\citep[and reference therein]{koo04}. Unanimous consent seems only to hold on the statement that environmental processes are effective at removing the outer ISM and therefore affecting the HI component \\citep{cay94,gio85}, leaving the molecular component, well--bounded inside the galaxy potential well, undisturbed \\citep{ken89, bos02c}. In order to study the environmental effects on star formation, we collect both objects that are members of the Virgo cluster and local field galaxies. Since the different components involved in the star formation are observable along a broad stretch of the electromagnetic spectrum, we adopt a multifrequency analysis. To quantify the atomic hydrogen we collect observations at 21 cm; the molecular hydrogen content is estimated indirectly via CO observations at 2 mm; we quantify the stellar mass with observations at the near--infrared (NIR) or visible bands and we study the presence of new stars indirectly, observing the hydrogen recombination line (H$\\alpha$) at 6563 \\AA. The sample is presented in Sec. \\ref{data}, together with the description of the data reduction procedures; the analysis is given in Sec. \\ref{analysis}, while we discuss our results and conclusions in Sec. \\ref{results} and \\ref{concl}. ",
        "conclusions": "With the aim of exploring the relations between the process of star formation and the physical properties of the ISM in various galaxy environments, we collected state-of-the-art imaging material and maps for 28 massive spiral galaxies belonging to the Virgo cluster and to the local field. The observational material includes: images of the stellar continuum (taken in the red/NIR bands); H$\\alpha$ images of the young ($<4\\times 10^6$ yrs) stars; radio maps at 2mm from the recent {\\it Nobeyama CO atlas of nearby spiral galaxies} that combines good sensitivity to large-scale CO emission with a 15 arcsec spatial resolution, and sensitive radio maps at 21 cm from the ongoing, yet unpublished, VIVA and THINGS surveys carried out at the VLA. Physical parameters have been derived homogeneously and carefully for the individual galaxies applying corrections based on measured quantities, rather than average values. We hope that the effort we put in the analysis of the data reduces the statistical limitations that arise from the paucity of imaging material currently available.  \\\\ The present analysis indicates that the bulk of the star formation in spiral galaxies is supported by a diffuse molecular medium which forms through the conversion of the atomic hydrogen, due to the pressure exerted by the stellar potential. However, at the edge of the star forming disks, the HI plays a more important role than previously believed in directly sustaining the star formation activity. When environmental processes cause significant removal of the outer part of the HI disks, galaxies suffer from starvation; the replenishment of the atomic gas inside the optical disks is reduced, leading to a depletion of the molecular component that is consumed during the star formation activity, thus causing the quenching of the star formation activity itself. When the HI removal is so severe that the HI disk shrinks inside the optical radius, there is also a truncation of both the molecular and the star forming disk."
    },
    "0808/0808.0746_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{ch1:sec1}  \\subsection{The Current Paradigm of Early Universe Cosmology}  According to the inflationary universe scenario \\cite{Guth} (see also \\cite{Sato,Starob1,Brout}), there was a phase of accelerated expansion of space lasting at least 50 Hubble expansion times during the very early universe. This accelerated expansion of space can explain the overall homogeneity of the universe, it can explain its large size and entropy, and it leads to a decrease in the curvature of space. Most importantly, however, it includes a causal mechanism for generating the small amplitude fluctuations which can be mapped out today via the induced temperature fluctuations of the cosmic microwave background (CMB) and which develop into the observed large-scale structure of the universe \\cite{Mukhanov} (see also \\cite{Press,Sato,Starob2,Lukash}). The accelerated expansion of space stretches fixed co-moving scales beyond the Hubble radius. Thus, it is possible to have a causal mechanism which generates the fluctuations on microscopic sub-Hubble scales. The wavelengths of these inhomogeneities are subsequently inflated to cosmological scales which are super-Hubble until the late universe. The generation mechanism is based on the assumption that the fluctuations start out on microscopic scales at the beginning of the period of inflation in a quantum vacuum state. If the expansion of space is almost exponential, an almost scale-invariant spectrum of cosmological perturbations results, and the squeezing which the fluctuations undergo while they evolve on scales larger than the Hubble radius predicts a characteristic oscillatory pattern in the angular power spectrum of the CMB anisotropies \\cite{Sunyaev}, a pattern which has now been confirmed with great accuracy \\cite{Boomerang,WMAP} (see e.g \\cite{MFB} for a comprehensive review of the theory of cosmological fluctuations, and \\cite{RHBrev2} for an introductory overview).  To establish our notation, we write the metric of a homogeneous, isotropic and spatially flat four-dimensional universe in the form \\be \\label{metric} ds^2 \\, = \\, dt^2 - a(t)^2 d{\\bf x}^2 \\, , \\ee where $t$ is physical time, ${\\bf x}$ denote the three co-moving spatial coordinates (points at rest in an expanding space have constant co-moving coordinates), and the scale factor $a(t)$ is proportional to the size of space. The expansion rate $H(t)$ of the universe is given by \\be H(t) \\, = \\, \\frac{{\\dot a}}{a} \\, , \\ee where the overdot represents the derivative with respect to time.  \\begin{figure} \\begin{center} \\includegraphics[height=9cm]{canc1.eps} \\caption{Space-time diagram (sketch) of inflationary cosmology. Time increases along the vertical axis. The period of inflation begins at time $t_i$, ends at $t_R$, and is followed by the radiation-dominated phase of standard big bang cosmology. If the expansion of space is exponential, the Hubble radius $H^{-1}$ is constant in physical spatial coordinates (the horizontal axis), whereas it increases linearly in time after $t_R$. The physical length corresponding to a fixed co-moving length scale is labelled by its wave number $k$ and increases exponentially during inflation but increases less fast than the Hubble radius (namely as $t^{1/2}$), after inflation. Hence, the wavelength crosses the Hubble radius twice. It exits the Hubble radius during the inflationary phase at the time $t_i(k)$ and re-enters during the period of standard cosmology at time $t_f(k)$.} \\end{center} \\label{fig:1} \\end{figure}  A space-time sketch of inflationary cosmology is shown in Fig. 1. The vertical axis is time. The inflationary phase begins at the time $t_i$ and lasts until the time $t_R$, the time of ``reheating\". At that time, the energy which is driving inflation must change its form into regular matter. The Hubble radius is labelled by $H^{-1}(t)$ and divides scales into those where micro-physics dominates and thus the generation of fluctuations by local physics is possible (sub-Hubble scales) and those where gravity dominates and micro-physical effects are negligible (super-Hubble). As shown in the sketch, during inflation fixed co-moving scales (labelled by $k$ in the sketch) are inflated from microscopic to cosmological. Note also that the horizon, the forward light cone, becomes exponentially larger than the Hubble radius during the inflationary phase.  \\subsection{Challenges for String Cosmology}  Working in the context of General Relativity as the theory of space-time, inflationary cosmology requires the presence of a new form of matter with a sufficiently negative pressure $p$ ($p < - 2/3 \\rho$, where $\\rho$ denotes the energy density). In order to obtain such an equation of state, in general the presence of scalar field matter must be assumed. In addition, it must be assumed that the scalar field potential energy dominates over the scalar field spatial gradient and kinetic energies for a sufficiently long time period. The Higgs field used for the spontaneous breaking of gauge symmetries in particle physics has a potential which is not flat enough to sustain inflation. Models beyond the Standard Model of particle physics, in particular those based on supersymmetry, typically have many scalar fields. Nevertheless, it has proven to be very difficult to construct viable inflationary models. The problems which arise when trying to embed inflation into the context of effective field theories stemming from superstring theory are detailed in the contribution to this book by Burgess.  If inflationary cosmology is realized in the context of classical General Relativity coupled to scalar field matter, then an initial cosmological singularity is unavoidable \\cite{Borde}. Resolving this initial singularity is one of the challenges for string cosmology.  The energy scale during inflation is set by the observed amplitude of the CMB fluctuations. In simple single field models of inflation, the energy scale is of the order of the scale of Grand Unification, i.e. many orders of magnitude larger than scales for which field theory has been tested experimentally, and rather close to the string and Planck scales, scales where we know that the low energy effective field theory approach will break down. It is therefore a serious concern whether the inflationary scenario is robust towards the inclusion of non-perturbative stringy effects, effects which we know must not only be present but in fact must dominate at energy scales close to the string scale.  The problem for cosmological fluctuations is even more acute: provided that the inflationary phase lasts for more than about 70 Hubble expansion times, then all scales which are currently probed in cosmological observations had a wavelength smaller than the Planck length at the beginning of the inflationary phase. Thus, the modes definitely are effected by trans-Planckian physics during the initial stages of their evolution. The ``trans-Planckian problem\" for fluctuations \\cite{RHBrev1,Jerome} is whether the stringy effects which dominate the evolution in the initial stages leave a detectable imprint on the spectrum of fluctuations. To answer this question one must keep in mind that the expansion of space does not wash out specific stringy signatures, but simply red-shifts wavelengths. For string theorists, the above ``trans-Planckian problem\" is in fact a window of opportunity: if the universe underwent a period of inflation, this period will provide a microscope with which string-scale physics can be probed in current cosmological observations.  Some of the conceptual problems of inflationary cosmology are highlighted in Figure 2, a space-time sketch analogous to that of Figure 1, but with the two zones of ignorance (length scales smaller than the Planck (or string) length and densities higher than the Planck (or string) density) are shown. As the string scale decreases relative to the Planck scale, the horizontal line which indicates the boundary of the super-string density zone of ignorance approaches the constant time line corresponding to the onset of inflation. This implies that the inflationary background dynamics itself might not be robust against stringy corrections in the dynamical equations.  The sketch in Figure 2 also shows the exponential increase of the horizon compared to the Hubble radius during the period of inflation.  \\begin{figure} \\begin{center} \\includegraphics[height=9cm]{infl3.eps} \\caption{Space-time diagram (sketch) of inflationary cosmology including the two zones of ignorance - sub-Planckian wavelengths and trans-Planckian densities. The symbols have the same meaning as in Figure 1. Note, specifically, that - as long as the period of inflation lasts a couple of e-foldings longer than the minimal value required for inflation to address the problems of standard big bang cosmology - all wavelengths of cosmological interest to us today start out at the beginning of the period of inflation with a wavelength which is in the zone of ignorance.} \\end{center} \\label{fig:2} \\end{figure}  \\subsection{Preview}  The conceptual problems of inflationary cosmology discussed in the previous subsection motivate a search for a new paradigm of early universe cosmology based on string theory. Such a new paradigm may provide the initial conditions for a robust inflationary phase. However, it may also lead to an alternative scenario. In the following, we will explore this second possibility.  In the best possible world, the initial phase of string cosmology will eliminate the cosmological ``Big Bang\" singularity, it will provide a unified description of space, time and matter, and it will allow a controlled computation of the induced cosmological perturbations. The development of such a consistent framework of string cosmology will, however, have to be based on a consistent understanding of non-perturbative string theory. Such an understanding is at the present time not available.  Given the lack of such an understanding, most approaches to string cosmology are based on treating matter using an effective field theory description motivated by string theory. However, in such approaches key features of string theory which are not present in field theory cannot be seen. The approach to string cosmology discussed below is, in contrast, based on studying effects of new degrees of freedom and new symmetries which are key ingredients to string theory, which will be present in any non-perturbative formulation of string theory. ",
        "conclusions": "The string gas scenario is an approach to early universe cosmology based on coupling a gas of strings to a classical background. It includes string degrees of freedom and string symmetries which are hard to implement in an effective field theory approach.  The background of string gas cosmology is non-singular. The temperature never exceeds the limiting Hagedorn temperature. If we start the evolution as a dense gas of strings in a space in which all dimensions are string-scale tori, then there are dynamical arguments according to which only three of the spatial dimensions can become large \\cite{BV}. Thus, string gas cosmology yields the hope of understanding why - in the context of a theory with more than three spatial dimensions - exactly three are large and visible to us.  If the Hagedorn phase (the phase during which the temperature is close to the Hagedorn temperature and both the scale factor and the dilaton are static) is sufficiently long to establish thermal equilibrium on length scales of about 1 mm, then string gas cosmology can provide an alternative to cosmological inflation for explaining the origin of an almost scale-invariant spectrum of cosmological fluctuations \\cite{NBV}. A distinctive signature of the scenario is the slight blue tilt in the spectrum of gravitational waves which is predicted \\cite{BNPV1}.  The inflationary universe scenario has successes beyond the fact that it successfully predicted a scale-invariant spectrum of fluctuations - it also explains why, starting from a hot Planck scale space, an extremely low entropy state - one can obtain a universe which is large enough and contains enough entropy to correspond to our observed universe. In addition, it explains the observed spatial flatness. However, if the Hagedorn phase of string gas cosmology is realized as a long bounce phase in a universe which starts out large and cold, then the horizon, flatness, size and entropy problems do not arise.  A serious concern for the current realization of string gas cosmology, however, is the gravitational Jeans instability problem. This problem was first raised in \\cite{stability}. In the context of dilaton gravity, it can be shown \\cite{dilflucts} that gravitational fluctuations do not grow. However, dilaton gravity is not a consistent background for the Hagedorn phase of string gas cosmology. One might hope that since the string states are relativistic, the gravitational Jeans length will be comparable to the Hubble radius, as it is for a gas of regular radiation. However, a recent computation of the speed of sound in string gas cosmology \\cite{Nima} has shown that in a background space sufficiently large to evolve into our present universe the overall speed of sound is very small. Further work needs to be done on this issue. This is complicated by fact that string thermodynamics is non-extensive (see e.g. \\cite{Cobas2}), which leads to problems in using the usual thermodynamic intuition.  Note that the background space does not need to be toroidal. Crucial for string gas cosmology to yield the predictions summarized above is the existence and stability (or quasi-stability) of string winding modes. Certain orbifolds \\cite{Col1} have also been shown to yield good backgrounds for string gas cosmology. Non-trivial one cycles will ensure the existence and stability of string winding modes.  If the background space does not have have any non-trivial one-cycles, then it might be possible to construct a cosmological scenario based on stable branes rather than strings. The cosmology of brane gases has been considered in \\cite{branes}. If there are stable winding strings, and if the string coupling constant is small such that the fundamental strings are lighter than branes, then \\cite{ABE} it is the fundamental strings which will dominate the thermodynamics in the Hagedorn phase and which will be the most important degrees of freedom for cosmology. However, if there are no stable winding strings, then winding branes would become important.  It appears at the present time that Heterotic string theory is most suited for string gas cosmology since this theory admits the enhanced symmetry states which have been shown to yield a very simple way to stabilize the size and shape moduli of the extra spatial dimensions. It will be interesting to study if string gas cosmology can be embedded into particular models of Heterotic string theory which yield reasonable particle phenomenology.  The presentation we gave of string gas cosmology is based on minimal input. In particular, we did not include fluxes since we assume that the net fluxes should cancel for a situation with the most symmetric initial conditions. The role of fluxes in string gas cosmology has been studied in \\cite{Campos}. Whereas the primary application of string gas cosmology will be to the cosmology of the very early universe, it is also interesting to consider applications of string gas cosmology to later time cosmology. The late time dynamics of string and brane gases has been considered in \\cite{late}. In particular, in \\cite{Ferrer} applications of string gases to the dark energy problem has been considered (see also \\cite{McInnes}). String and brane gases have also been studied as a way to obtain inflation \\cite{Turok,Parry,Anupam} (see also \\cite{Freese}), or as a way to obtain non-inflationary bulk expansion which may provide a way to solve the size problem in string gas cosmology if one starts with a spatial manifold of string scale in all directions \\cite{Natalia}.  \\vskip0.3cm  \\centerline{\\bf Acknowledgements}  \\vskip0.2cm  I am grateful to all of my present and former collaborators with whom I have had the pleasure of working on string gas cosmology. For comments on the draft of this manuscript I wish to thank Nima Lashkari and Subodh Patil. This work is supported in part by an NSERC Discovery Grant and by funds from the Canada Research Chairs Program."
    },
    "0808/0808.1740_arXiv.txt": {
        "abstract": "The broad-line radio galaxy 3C\\,111 has been suggested as the counterpart of the $\\-$--ray source 3EG\\,J0416+3650.  While 3C\\,111 meets most of the criteria for a high-probability identification, like a bright flat-spectrum radio core and a blazar-like broadband SED, in the Third EGRET Catalog, the large positional offset of about 1.5$^\\circ$ put 3C\\,111 outside the 99\\% probability region for 3EG\\,J0416+3650, making this association questionable.  We present a re-analysis of all available data for 3C\\,111 from the EGRET archives, resulting in probable detection of high-energy $\\gamma$--ray emission above 1000\\,MeV from a position close to the nominal position of 3C\\,111, in two separate viewing periods (VPs), at a 3$\\sigma$ level in each.  A new source, GRO\\,J0426+3747, appears to be present nearby, seen only in the $>$1000\\,MeV data.  For $>$100\\,MeV, the data are in agreement with only one source (at the original catalog position) accounting for most of the EGRET-detected emission of 3EG\\,J0416+3650.  A follow-up \\textit{Swift} UVOT/XRT observation reveals one moderately bright X--ray source in the error box of 3EG\\,J0416+3650, but because of the large EGRET position uncertainty, it is not certain that the X--ray and $\\gamma$--ray sources are associated.  A \\textit{Swift} observation of GRO\\,J0426+3747 detected no X--ray source nearby. ",
        "introduction": "One of the main scientific goals of the recently-launched $\\gamma$--ray astronomy satellite mission GLAST and its Large Area Telescope (LAT) \\citep{Ri07, Mi07} is to shed light on the nature of powerful relativistic extragalactic jets, which are ejected from the nuclei of some active galaxies (AGN).  Originally, this class of AGN was defined based on the bright and prominent radio emission from these jets: the radio-loud population of active galaxies.  Based on data from EGRET, the high-energy $\\gamma$--ray telescope on the {\\it Compton Gamma Ray Observatory}, it was realized \\citep{Fi94, Th94, vM95} that the largest population of extragalactic $\\gamma$--ray sources in the GeV regime is represented by those radio-loud AGN whose jets are pointed at a small angle to the line of sight.  An intimate link is tied between radio VLBI and high-energy $\\gamma$--ray astronomy by the fact that the bright compact radio emission of these so-called blazars provides excellent targets for parsec-scale resolution VLBI observations of their jet structure.  Of special interest are the questions of where in the AGN jets the bright $\\gamma$--ray emission is produced, and how the emission is interacting with its immediate environment and with other parts of the jet.  This knowledge would enable us to put crucial constraints on the processes of jet formation, collimation, and acceleration.  GLAST is expected to yield densely sampled $\\gamma$--ray light curves of hundreds of extragalactic jets that are bright enough to be detected on time scales of days to weeks, and thousands on time scales of months to years.  Most of these objects will be blazars; in fact, all but two of the firmly identified extragalactic EGRET sources (LMC, Sreekumar et al. 1992; Centaurus\\,A, Sreekumar et al. 1999) are blazars.  In contrast to the blazars, most radio galaxies have larger inclination angles.  That allows better (deprojected) linear resolution with VLBI observations, and in the case of stratified jet structures, it allows observations of the slower jet layers, e.g. a sheath, whose emission may be swamped by the much brighter beamed emission from faster jet regions, e.g. a fast spine, in blazars.  \\citet{Ghi05} have presented such a spine-sheath stratified-jet model, predicting detectable $\\gamma$--ray emission from the nuclei of some radio galaxies.  In their model, each of the two emission regions sees photons coming from the other part relativistically enhanced because of the relative speed difference, giving rise to an additional inverse-Compton emission component.  In this letter we provide additional support for the identification of the broad-line radio galaxy 3C\\,111 as a $\\gamma$--ray source, responsible for a portion of the source 3EG\\,J0416+3650.  Such an association was suggested as possible in the third EGRET catalog \\citep[3EG]{Ha99}, but was considered unlikely because of the large positional offset of 3C\\,111 from 3EG\\,J0416+3650.  The present work was stimulated by the recent report of \\citet{Sg05}, who used multiwavelength data to strengthen the case for the association between 3C\\,111 and 3EG\\,J0416+3650.  We show here that 3EG\\,J0416+3650 is most likely composed of at least two, and more likely three, separate sources, one of which is in good positional agreement with 3C\\,111. As demonstrated recently, based on VLBA monitoring data at $\\lambda 2$\\,cm from the MOJAVE program \\citep{Kad08}, the parsec-scale jet of 3C\\,111 shows a variety of physically different regions in a relativistic extragalactic jet, such as a compact core, superluminal jet components, recollimation shocks, and regions of interaction between the jet and its surrounding medium, which are all possible sites of $\\gamma$--ray production. Its relatively large inclination angle of $\\sim$19$^\\circ$ makes 3C\\,111 a particularly well-suited target for tests of structured-jet models, such as the model of \\citet{Ghi05}.  We describe our re-analysis of the available EGRET data on 3C\\,111, as well as follow-up \\textit{Swift} OVOT and XRT observation, in Sect.~\\ref{sect:analysis}, and discuss our results and their implications in Sect.~\\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions} The $\\gamma$--ray source 3EG\\,J0416+3650 seems to be composed of at least three variable sources.  One of the sources, detected only above 1000\\,MeV, is close to the radio galaxy 3C\\,111, and plausibly associated with it; a new source, GRO\\,J0426+3747, is also seen only above 1000\\,MeV, and has no obvious identification at other wavelengths.  Most of the flux $>$100\\,MeV is from the third source, for which there is no evidence in the $>$1000\\,MeV data.  The position obtained here for the third source is very near the catalog position of 3EG\\,J0416+3650.  It  might be associated with the XRT source Swift\\,J041554.3+364926, but the EGRET position uncertainty makes a firm association impossible.  The region under study is $8-10^\\circ$ from the Galactic plane, so either 3EG\\,J0416+3650 or GRO\\,J0426+3747 (or both) could be local.  A recent stacking search found no evidence for $\\gamma$--ray emmision from radio galaxies as a class \\citep{Cil04}.  That paper specifically excluded 3C\\,111 from consideration because of its blazar-like superluminal motion, so there is no direct conflict with the results presented here.  That study made no systematic cut on the angle to the line-of-sight of the radio jet, so for most of the galaxies included, that angle was considerably larger than the $\\sim$19$^\\circ$ which has been estimated for 3C\\,111 \\citep{Kad08}.  Since a larger line-of-sight angle implies substantially smaller Doppler boosting and therefore lower $\\gamma$--ray output, there is no obvious conflict between the present results and those of \\citet{Cil04}.  A more significant challenge is posed by the lack of EGRET detection of $\\gamma$--ray emission from the (much closer) AGN M87.  Superluminal motion has been detected in the jet knot HST-1, about 0.86\\,arcsec downstream from the radio core, both in the optical and radio bands \\citep{Bir99,Che07}. Most estimates of our line-of-sight angle to its jet are in the range $30-45^\\circ$.  If we assume that M87 is similar to 3C\\,111, it is not clear whether its Doppler boosting is sufficiently low to compensate for its much smaller distance, and thereby account for its non-detection by EGRET.  Our re-analysis of the available EGRET data supports the previous association of the source 3EG\\,J0416+3650 with the broad line radio galaxy 3C\\,111, but with 3C\\,111 responsible for only a portion of the $\\gamma$--ray emission.  Furthermore, it explains the relatively large positional offset noted in \\citet{Ha99}. We have compiled an historical SED of 3C\\,111 which shows that the X--ray data may well extrapolate into the EGRET range, particularly during flares, which is in agreement with the intermittent nature of detections in the individual EGRET viewing periods. It is interesting to note that we detect 3C\\,111 at almost exactly the $\\gamma$--ray flux that is predicted by equation (12) in \\citet{Ghi05}: $(1.41-14.1) \\times 10^{-11}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$, scaling from the 5\\,GHz values in Table~\\ref{tab:3c111_sed}.  Note that 3C\\,111 (z=0.0485) would be the most distant radio galaxy detected in $\\gamma$ rays, about twice as far as NGC\\,6251 (Mukherjee et al. 2002; z=0.0247).  The detection of $\\gamma$--ray emission from 3C\\,111 further supports the hypothesis that radio galaxies may represent an important class of LAT sources.  GLAST was launched on 11 June, 2008; its LAT detector will continuously scan the entire sky over a 3-hr interval.  With a factor of 30 greater sensitivity than EGRET, and with a factor of $\\sim$3 better PSF above 1\\,GeV, it will be more efficient than EGRET for detecting transients over the whole sky. In the case of 3C\\,111, GRO\\,J0426+3747, and 3EG\\,J0416+3650, the smaller PSF and greater effective area of the LAT will clearly separate these three sources."
    },
    "0808/0808.1789_arXiv.txt": {
        "abstract": "We present U, B, V, R, I, H$\\alpha$ and NUV photometry of 14 galaxies in the very local Universe (within 10 Mpc). Most objects are dwarf irregular galaxies (dIrr) and are probably associated with the NGC 672/IC 1727 and NGC 784 galaxy groups. The galaxies are at low redshift (51$\\leq$v$_{\\odot} \\leq$610 km sec$^{-1}$) and most appear projected on the sky as a six degree long linear filament. We show that the galaxy positions along this filament correlate with their radial velocity, hinting to an interpretation as a single kinematic entity. Our CCD photometry indicates that all objects qualify as ``dwarf galaxies'' with M$_B\\geq-18$ mag. We examine the star formation (SF) properties of individual objects in the context of their immediate environment. The current SF rate (SFR) is derived directly from the H$\\alpha$ line flux. An approximate SF history is derived by comparing the multi-band photometry with results of galaxy evolution models from Bruzual \\& Charlot (2003a, 2003b), assuming short SF bursts separated by long quiescence periods.  Relations between the current SFR and the HI mass or the absolute B magnitude for the galaxies in these groups indicate that these objects behave like normal galaxies. A comparison of the photometric measurements with evolutionary synthesis model predictions indicates that most objects can be understood as containing at least one ``old'' stellar population ($\\geq$1-10 Gyr) and one ``young'' population ($\\leq$30 Myr). For both groups, the recent SF bursts appear to have occurred at similar times, a few to a few 10s of Myr ago, arguing for synchronicity in star formation in these objects.  In an attempt to evaluate the possible role of galaxy-galaxy interaction, we investigate the trend of the SFR with an object's projected distance from the brightest and most massive galaxies of each group. We do not find a steadily decreasing star formation as function of this distance; such a result could be expected if the star formation would have been triggered by interactions. We propose that one possible explanation of the $\\sim$synchronous star formation in all objects is accretion of cold gas from intergalactic space onto dark matter haloes arranged along a filament threading the void where these dwarf galaxies reside. We point out this galaxy sample as an ideal target to study hierarchical clustering and galaxy formation among very nearby objects. ",
        "introduction": "Star formation is one of the more important processes in the Universe. Studying both the SF history in galaxies and the influence of the environment on this process contributes to the understanding of cosmic evolution. In the past few decades increasing efforts have been made to better understand the dependence of SF on local and environmental properties. Tidal interactions were promoted as triggers and enhancers of SF (e.g., Larson \\& Tinsley 1978, Li et al.\\ 2008). In practice, the influence of tidal interactions is more complex, depends upon various factors, and varies among different environments. Hashimoto et al.\\ (1998) studied the variation and dependence of these influences in different environments for different types of galaxies and found different SFRs in the field and in clusters, even when the galaxy densities are similar. They also found that for a given galaxy concentration index, galaxies in lower-density environments show higher SFR levels than objects in higher galaxy density regions. At least two processes were found to influence the susceptibility of SF to the environment: gas-removal processes responsible for the variation with galaxy density of the SFRs of normal galaxies, and galaxy-galaxy interactions responsible for the prevalence of SF bursts in intermediate-density environments. At low redshifts, galaxies in groups may show a reverse relation with a somewhat higher SFR (but not extreme burst conditions) the more isolated a galaxy is (see Martig \\& Bournaud 2008). However, Brosch et al.\\ (2004) examined specifically the influence of interactions on SF in dwarf galaxies and concluded that galaxy interactions are probably not a primary trigger of present SF in such objects.  About half the galaxies in the Universe are probably in groups (Mamon 2007) and these constitute an interesting environment for the SF investigation. The evolution of galaxies within groups may depend upon the density of the group and the level of interaction between its galaxies. Many groups consist of normal galaxies surrounded by clouds of dwarf galaxies (Miller 1996, Cote et al.\\ 1997, Mateo 1998). The dwarf galaxies (DGs) may have formed due to tidal interaction of individual galaxies or of groups of galaxies, or as leftover material from the colliding and merging of more massive galaxies (e.g., Hunsberger et al.\\ 1996, Bournaud \\& Duc 2006). They may also represent early stages of hierarchical clustering, with the stars that provide the observed luminosity forming in a low-mass halo.  The DGs in a group contain usually only a few percents of the group mass and in many cases may be accreted by larger galaxies. Although dwarf galaxies constitute the majority of the galaxy population, large uncertainties surround their formation and evolutionary histories. The dwarf galaxy population follows a strong morphology-density relation, with passively evolving systems mostly found in close proximity to massive galaxies, in contrast to the more widespread gas-rich, star forming population. Studying galaxy groups, in particular the  dwarf-rich ones, contributes to the understanding of cosmic structure and evolution, since such groups constrain cosmological models, imply the shape of the dark halos around the massive galaxies, and play a role in the evolution of these galaxies (see Bournaud \\& Duc 2006 and references therein).  The present study complements similar ones on other groups of galaxies in the local Universe, such as for the NGC 628 and M81 groups (e.g., Sharina et al.\\ 2006, Karachentsev \\& Kaisin 2007). Hickson (1982) published a catalog of compact groups of galaxies in the local Universe, and investigated various characteristics of these groups (e.g., Hickson et al.\\ 1992, de Oliveira \\& Hickson 1991, 1994). Samples of galaxies in other environments have been examined, such as in the Virgo cluster (e.g., on dwarf galaxies, Almoznino \\& Brosch 1998a, b; Brosch, Heller \\& Almoznino 1998; Heller, Almoznino \\& Brosch 1999; Heller \\& Brosch 2001). The influences of galaxy interactions on star formation were reviewed by Larson \\& Tinsley (1978), Hashimoto et al.\\ (1998) and Li et al.\\ (2008). In most previous papers the galaxies were either normal (i.e., large spirals, lenticulars or ellipticals) or dwarf galaxies; the specific aspect the present paper adopts is to study these properties in a fairly isolated group or groups of DGs.  Tidal interactions between two strongly-interacting major galaxies tend to form sometimes rather concentrated clouds of DG satellites around them, or long tails of DGs as seen e.g., in the Hickson compact group 100 (de Mello et al. 2008). If one would observe dwarf galaxies in a group apparently aligned on the sky in a linear configuration, and with no major galaxies in the immediate neighborhood, this could imply the presence of a dark matter filament onto which the intergalactic matter collapses, forming the observed DGs.  This work studies a sample galaxies in the very local Universe ($\\sim$10 Mpc)  consisting of the NGC 672/IC 1727 group of galaxies, the NGC 784 group of galaxies, and several other galaxies in the same vicinity with similar redshifts. Our attention was drawn to this sample when inspecting the precursor observations of the ALFALFA survey (Giovanelli et al. 2005) because of its obvious linear structure that could qualify as a ``filament'' of galaxies, and because of its apparent isolation. The galaxies are located approximately in the anti-Virgo direction, in a region of the Universe that is a void in the galaxy distribution; this is the ``Local Mini-void'' (Karachentsev et al. 2004). As will be shown below, all galaxies are ``dwarfs'' (M$_B\\geq$-18 mag) and most are dIrr (e.g., Karachentseva \\& Karachentsev 1998, Huchtmeier et al. 2000). We study the SF properties and SF history of each galaxy using surface photometry, and examine general similarities and dependencies on the local, environmental and mutual properties.  Current $\\Lambda$CDM simulations predict that low-amplitude filamentary structures criss-cross the voids (Peebles 2007). The galaxies corresponding to the halos in those filaments are expected to be low-luminosity, star-forming galaxies (Hoyle et al. 2005). Saintonge et al. (2008) analyzed the ALFALFA catalog covering a portion of the nearby void in front of the Pisces-Perseus Supercluster at cz$\\sim$2000 km sec$^{-1}$ and, within a volume of 460 Mpc$^3$, did not detect a single galaxy. In contrast, one could expect to detect 38 HI sources in such a volume based on scaling the predictions of Gottl\\\"{o}ber et al. (2003) with a dark-to-HI mass ratio of 10:1.  Dekel \\& Birnboim (2006) discussed the bimodality observed in galaxy properties about a characteristic stellar mass of $3\\times10^{10}$ M$_{\\odot}$. The bimodality implies that less massive galaxies tend to be ungrouped, blue, star-forming disks and would be formed in low galaxy-density regions, while more massive galaxies are typically grouped, red, old-star spheroids and would reside in clusters and in denser groups. In haloes below a critical shock-heating mass of 10$^{12}$ M$_{\\odot}$ disks are built by cold-gas streams, not heated by a virial shock, yielding efficient and prompt star formation. The Dekel \\& Birnboim scenario explains naturally why one could expect to see dwarf galaxies even in regions that are not conducive to intense tidal interactions, provided dark matter haloes of the right mass are present and HI material for SF is available.  In what follows we argue that the objects studied here may be the first case in the local Universe where we witness the first stages of aggregation of matter that will, eventually, form a major galaxy. The plan of the paper is as follows: in Section~\\ref{sec.sample} we describe the galaxy sample, the observations, mostly collected at the Wise Observatory, and the data reduction, in \\S~\\ref{sec.results} we present the results and discuss individual objects, in \\S~\\ref{sec.discuss} we discuss the results and present arguments in favour of the interpretation proposed above, and in \\S~\\ref{sec.summary} we conclude and summarize our findings.  \\section {Observations and data reduction} \\label{sec.sample}  \\subsection {The sample}  The sample consists of the catalogued NGC 672/IC 1727 and NGC 784 galaxy groups and of several additional galaxies in the same sky area and at similar redshifts; these are listed in Table \\ref{gal}. The galaxies were chosen primarily from the ALFALFA survey (see Giovanelli et al.\\ 2005, Saintonge et al.\\ 2008). % ALFALFA is an unbiased HI survey of the extragalactic sky visible from Arecibo. The precursor ALFALFA observations (Giovanelli et al.\\ 2005) covered the sky area  (0h$\\leq$RA$\\leq$6h) and (26$^{\\circ} \\leq \\delta\\leq 28.5^{\\circ}$) and 10 galaxies in the vicinity of NGC 672/IC 1727 with similar redshifts (cz$\\leqslant$600 km sec$^{-1}$) were included among the 166 identified objects.  Figure \\ref{fig:galaxies} shows the projection of the different objects on the sky, with the redshift of each object indicated. It is clear that the galaxies form some kind of linear structure, which we call ``filament''; the projected distribution is approximately linear with the objects at the north-east side showing generally lower redshifts. The length of the filament is $\\sim$six degrees; a few galaxies diverge significantly from the linear distribution and we consider them to be ``isolated''. The distance to the galaxies is of order 5 Mpc; the $\\sim6^{\\circ}$ projected angular extent translates into a projected physical length for the linear structure of $\\sim$500 kpc. %  As already mentioned, most of the galaxies studied here were previously classified as dwarf irregulars (e.g., Karachentseva \\&  Karachentsev 1998, Huchtmeier et al.\\ 2000). In Table~\\ref{gal} we list not only the galaxy name, which in some cases is taken from the Arecibo Galaxy Catalogue (AGC: a private compilation of Haynes and Giovanelli maintained at Cornell), but also the heliocentric velocity in km sec$^{-1}$ and the 50\\% width of the HI line (also in km sec$^{-1}$) taken from Saintonge et al. (2008) or from the HyperLEDA data base (Paturel et al. 2003). Column 7 lists the position angle (PA) of the major axes of the objects, determined from a visual inspection of the blue galaxy images on-line at the Canadian Astronomy Data Centre (CADC). The PA is measured in degrees, clockwise from North through West. AGC  111945 was too faint and small to derive its PA. The last column of the table lists the galaxy group to which an object could belong.   \\begin{table}[!h]  \\caption{The sample galaxies arranged by group membership} \\vspace{0.5cm} \\label{gal} \\begin{footnotesize} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|c|} \\hline Galaxy & ALFALFA name & $\\alpha$[hh:mm:ss.s] & $\\delta$[$^\\circ$:':''] & v$_{\\odot}$ & w$_{50}$ & PA ($^\\circ$) & Group\\\\ \\hline NGC 672 & HI014753.9+272555 & 01:47:54.5 & +27:25:58 & 429 & 205 & 332 & NGC 672\\\\ IC 1727 & HI014729.9+271958 & 01:47:29.9 & +27:20:00 & 330 & 115 & 35 & NGC 672\\\\ AGC 110482 & HI014214.9+262202 & 01:42:17.3 & +26:22:00 & 357 & 30 & 53 & NGC 672\\\\ AGC 111945 & HI014441.4+271707 & 01:44:42.7 & +27:17:18  & 423 & 36 & ? & NGC 672\\\\ NGC 111946 & HI014640.9+264754 & 01:46:42.2 & +26:48:05 & 367 & 21 & 0 & NGC 672\\\\ AGC 112521 & HI014105.8+272007 & 01:41:08.0 & +27:19:20 & 274 & 26 & 291 & NGC 672\\\\ LEDA169957 & --- & 01:36:35.9 & +23:48:54 & 563 & 47 & 307 & Isolated\\\\ NGC 784 & HI020115.0+284953 & 02:01:16.9 & +28:50:14 & 193 & 88 & 0 & NGC 784\\\\ AGC 111977 & HI015519.2+275645 & 01:55:20.4 & +27:57:13 & 207 & 29 & 297 & NGC 784\\\\ AGC 111164 & HI020009.3+284954 & 02:00:10.2 & +28:49:53  & 164 & 27 & 24 & NGC 784\\\\ UGC 1281 & --- & 01:49:32.0 & +32:35:23 & 156 & 98 & 306 & NGC 784\\\\ AGC 122834 & HI020347.0+291153 & 02:03:47.0 & +29:11:53 & 51 & 10 & 297 & NGC 784?\\\\ AGC 122835 & HI020533.0+291358 & 02:05:33.0 & +29:13:58 & 50 & 23 & - & N784; no Opt ID \\\\ UGC 1561 & --- &02:04:05.1 & +24:12:30 & 610 & 47 & 90 & Isolated\\\\ NGC 855 & HI021404.3+275302 & 02:14:03.6  & +27:52:36 & 594 & 81 & 305 & Isolated\\\\ \\hline \\end{tabular} \\end{center} \\end{footnotesize} \\end{table}  \\begin{figure}[t] \\centering{ \\includegraphics[width=12cm]{Gal3DwithVel.eps} \\caption{The galaxies as projected on the sky, represented by circles of different sizes. The circle sizes are proportional to the galaxies' semi-major axes. Redshifts are marked next to each galaxy symbol. A linear structure is evident from the distribution of 11 of the 14 objects plotted here.} \\label{fig:galaxies}} \\end{figure}   Though we found no previous results covering the SF in the sample galaxies, some objects were investigated before: the NGC 672/IC 1727 pair and their interaction have been studied by de Vaucouleurs et al.\\ (1976), Combes et al.\\ (1980), Hodge \\& Kennicutt (1983), Sohn \\& Davidge (1996), Karachentseva \\& Karachentsev (1998);  NGC 784 and some members of its group have been studied by Karachentseva \\& Karachentsev (1998), Huchtmeier et al.\\ (2000), Tikhonov \\& Karachentsev (2006), Tully et al.\\ (2006); other galaxies of the sample were studied by e.g., Wallington et al.\\ (1988), Huchtmeier \\& Richter (1989), Moustakas \\& Kennicutt (2006). Most galaxies were measured in HI by the ALFALFA survey (e.g., Giovanelli et al.\\ 2005) or are included in Karachentsev's catalog of neighboring galaxies (Karachentsev et al. 2004). Saintonge et al.\\ (2008) studied the HI properties of the NGC 672/IC 1727 and NGC 784 galaxy groups among other objects.  Saintonge et al.\\ (2008) included the NGC 672/IC 1727 and the NGC 784 groups along with two new candidate members of the NGC  784 group among the objects they studied. Only one of those two candidates, AGC 122834, has an optical counterpart. The other, AGC 122835 at 02:05:33.0 +29:13:58, is also listed in Table~\\ref{gal}. In addition to the 12 objects retrieved from the ALFALFA data sets, NED was queried for objects up to 300 arcmin away and with similar velocities; this added three additional neighboring galaxies. All the objects fall within the ``Local Mini-void'' Karachentsev et al.\\ (2004, Fig. 4). Note that the ALFALFA survey is likely to find even more members of this filament once more survey data for other declination strips will be released and redshifts of low surface brightness objects will be measured.  The literature contains a number of distance estimates to objects in our sample. Sohn \\& Davidge (1996) assigned a distance of 7.9$^{+1.0}_{-0.9}$ Mpc to NGC 672 based on the brightness of the red supergiants. Karachentsev et al.\\ (2004) reported 7.2 Mpc to the NGC 672/IC 1727 pair based on the luminosity of the brightest stars. A distance of 5.0 Mpc was measured in a similar way for NGC 784 by Drozdovsky \\& Karachentsev (2000). The distance to AGC 111977 and AGC 111164 was reported as 4.7 Mpc by Karachentsev et al.\\ (2004) using the red giant branch method. The distances to UGC 1281 was estimated by Karachentsev et al.\\ (2004) to be 5.4 Mpc using the brightest stars. The reported distance for NGC 855 (using the surface brightness fluctuations method) was 9.73$\\pm$1.7 Mpc (Tonry et al. 2001).  The question of distances to these galaxies was discussed by Giovanelli et al. (2005) who used the peculiar velocity model of the local Universe of Tonry et al. (2000). The model provides the expected peculiar velocity at a given point in space, but can be inverted to yield the distance to a galaxy at certain celestial coordinates and with a measured redshift. Giovanelli  al. caution that the ``thermal'' rms peculiar velocity of galaxies in the Tonry et al. model translates into a distance uncertainty of about 2.7 Mpc and thus the distances to nearby objects, and in particular to the NGC 672 and NGC 784 groups, are highly uncertain.  Since the distance estimates were produced by a variety of methods, one wonders whether the distance differences between members of the sample is believable. In particular, when Karachentsev et al.\\ (2004) quote 4.7 Mpc to AGC 111977 and AGC 111164 and 9.73 Mpc to NGC  855, can we indeed assume that these objects are separated by $\\sim$5 Mpc in depth? Is it possible that the objects could be at about the same distance, despite their somewhat different redshifts?  Given the uncertainty in distance, we decided to disregard peculiar velocities and use distances assuming only Hubble expansion with H$_{0}$=73 km sec$^{-1}$/Mpc to evaluate intrinsic parameters for each galaxy. The HI masses and the SFRs were calculated assuming these distances. However, in the discussion of the 3D structure traced by the galaxies we retain the option of considering the objects to be aligned along a filament (that may be inclined to the line of sight) and being at approximately the same distance from us.   \\begin{figure}[t] \\centering{ \\includegraphics[width=12cm]{VelocityVsX.eps} \\caption{Position-velocity diagram for all the galaxies as projected on the sky. The horizontal axis marked X is the distance along a line fitted to the galaxy positions. The two symbols at the top-right part of the diagram are UGC 1561 and NGC 855.} \\label{fig:gals_X_vs_Vel}} \\end{figure}  We plot in Figure~\\ref{fig:gals_X_vs_Vel} the position of each galaxy along the line fitted to their projected location on the sky vs. the heliocentric radial velocity of each object. The plot shows that most galaxies line up nicely on a $\\sim$6$^{\\circ}$ linear feature, with those at higher Right Ascension having generally a lower radial velocity. This kinematic relation can be understood as additional evidence that the objects are probably part of the same large structure in the nearby Universe. The two points at the upper right in Figure~\\ref{fig:gals_X_vs_Vel} are UGC 1561 and NGC 855; it is possible that these might not be part of the same kinematic structure delineated by the other galaxies.  \\subsection {Observations}  All broad-band and H$\\alpha$ observations were performed with the Wise Observatory (WiseObs) 1-m telescope between September 2006 and February 2008. We obtained CCD images of all 14 galaxies of the sample using the PI VersArray camera with a $1340\\times1300$ pixel thinned and back-illuminated CCD and a scale of 0.58\"/pixel at the f/7 Ritchey-Chr\\'{e}tien focus. The CCD saturation level is 65535 counts per pixel, the gain is 2.1 $e^{-}$/ADU, and the readout noise is 2.87 $e^{-}$ per pixel. The seeing on most nights was 2.2\"-2.8\". Images of each galaxy were obtained through UBVRI and zero-redshift H$\\alpha$ filters. Landolt (1992) and spectrophotometric (e.g., Oke 1990) standard stars were imaged along with the galaxy fields on nine photometric nights for H$\\alpha$ and broad-band calibration. Most of the galaxies were calibrated on more than one night and yielded a typical photometric error of 0.06 mag for all broad bands.  Typical exposure times were 20 minutes for the U and H$\\alpha$ filters, and 10 minutes for the other broad-band filters. Each galaxy was imaged at least three times through each filter, but for most galaxies 4-5 images per filter were obtained. Thus, at least 30 minutes of integrated exposure time were collected through the B, V, R and I filters, and at least 60 minutes through the U and H$\\alpha$ filters.  The WiseObs set of filters was replaced between the two sessions - fall of 2006 and fall of 2007. Due to instrumental problems we discarded all U-band images taken in 2006. Three galaxies were imaged and calibrated through the new U-band filter in 2007/8. Note also that all long-exposure images in the I band showed interference fringes. These were removed by an interactive script that uses the IRAF \\textit{rmfringe} command and an image of the fringe pattern with no objects in it, created by median-combining many night-time I-band images of different sky fields.  The H$\\alpha$ filter used has $\\lambda_{c}=6559\\pm$2{\\AA}, FWHM=56$\\pm$2{\\AA}, peak transmittance 58.9$\\pm0.3\\%$, and out-of-band transmittance $\\leq$0.06\\% (Spector 2006). Note that the measured H$\\alpha$ line flux includes, in principle, the [NII] lines at 6548{\\AA} and 6584{\\AA}. Since the ratio of [NII] to H$\\alpha$ line fluxes in dwarf irregular galaxies is $\\leq$10\\% (e.g, Kennicutt 1983), we disregard the possible [NII] contribution through the $H\\alpha$ filter. This is specifically justified for the four objects in our sample with global spectrophotometry from Moustakas \\& Kennicutt (2006), these all have [NII]/H$\\alpha \\leq$0.08.  The galaxies of the sample have radial velocities $\\leq$610 km sec$^{-1}$; this translates into a minor redward shift of the H$\\alpha$ emission line ($\\leq12${\\AA}), for an effective filter transmission of at least $\\sim95\\%$ of the peak transmittance. Since the galaxies are at $\\sim$zero redshift, the R filter was used for continuum measurement. Assuming a pessimistic photometric error of 0.1 mag, and adding in quadrature the possible [NII] contribution and the filter transmission contribution, the resulting H$\\alpha$ flux measurement accuracy is 0.15 mag. This increases to 0.18 mag for the net-H$\\alpha$ calculation (see below) due to the photometric error in the R band. For some galaxies, where both R and H$\\alpha$ have low photometric errors, the real calculated photometric error is smaller than this ``typical'' value.  We collected redshifts and HI line fluxes, which were converted to mass using $M_{HI} = 2.356 \\times 10^{5}D^{2}F_{c}$ with $M_{HI}$ in solar units and $D$ in Mpc, from the ALFALFA survey (see Giovanelli et al.\\ 2005, Saintonge et al.\\ 2008) and other published data (e.g, Huchtmeier \\& Richter 1989, Karachentsev et al.\\ 2004, NED). NUV images for most of the sample galaxies were retrieved from the GALEX archive (\\textit{http://galex.stsci.edu/GR4/}). The galaxies were measured with the same aperture as the broad-band images, and the measurements were calibrated using the formal GALEX procedure (Morrissey et al.\\ 2007). The GALEX NUV band has an effective wavelength of 2315.7{\\AA}, a peak at 2200{\\AA}, and effective bandwidth is from 1771{\\AA} to 2831{\\AA} (Morrissey et al.\\ 2007). We assume similarity between this filter and the 2200{\\AA} filter calculated in the Bruzual \\& Charlot models (2003b) and use the derived NUV magnitudes to construct a (NUV-V) colour index. The typical photometric error of this colour index is similar to those of the other indices used here.  \\subsection {Galaxy photometry} The galaxies were measured using the IRAF ``Ellipse'' function at a surface brightness level of $\\mu_B \\sim25.5$ mag arcsec$^{-2}$. The sky background level was determined using a galaxy-dependent method. For each image, the sky background was measured in an elliptical annulus whose inner semi-major axis (SMA) was $\\sim$1.35 times the primary SMA and its outer SMA was $\\sim$1.75 times the primary SMA. The ``Primary SMA'' term refers to the SMA of the elliptical aperture within which the galaxy was measured (at a brightness level of $\\mu_B \\sim25.5$ mag arcsec$^{-2}$).  The final result for each band was corrected for extinction and for a possible colour term. The results were also corrected for Galactic extinction (GE) according to the prescriptions of Schlegel et al.\\ (1998). Since the galaxies are in a similar sky area, the typical Milky Way (MW) extinction values for this area are $A_{U}\\sim0.4$, $A_{B}\\sim0.34$, $A_{V}\\sim0.26$, $A_{R}\\sim A_{H_{\\alpha}}\\sim0.2$, and $A_{I}\\sim0.17$, but exact values were calculated  specifically for each galaxy. For the GALEX NUV measurement the typical MW extinction was $\\sim0.9$ mag. The error in the determination of the extinction values is $\\sim10\\%$.  The continuum flux contribution through the R filter was subtracted from the H$\\alpha$ measurement to provide the net-H$\\alpha$ (nH$\\alpha$) line flux. The derived nH$\\alpha$ was subtracted from the average $R$ flux, thus the final $R$-band magnitudes listed below are free of H$\\alpha$ line emission. The nH$\\alpha$ images were produced by subtracting the R image from the H$\\alpha$ image, after normalization by the averaged total flux of the reference stars. These nH$\\alpha$ images were used to find, count and measure individual HII regions.  The photometry results, H$\\alpha$ fluxes, derived SFR (see below) and HI mass data are presented in Table~\\ref{T3}. We compared our derived photometry with published values and found reasonable compatibility. Four of our objects were observed by Moustakas \\& Kennicutt (2006) in their scanning-slit spectrophotometric survey of spiral galaxies. For these objects we found also reasonable correspondence with the Moustakas \\& Kennicutt values for the H$\\alpha$ flux being, on average, 8\\% brighter than ours.  Note that the brightest objects of the sample, NGC 672 and NGC 784, have absolute magnitudes -16.9 and -16.4 if located at a distance of 5 Mpc, and could be considered ``dwarf galaxies'' even if at twice the distance. This demonstrates that the sample is indeed composed exclusively of low-luminosity galaxies.   \\subsection {Comparison to galaxy evolution models} The SFH of each galaxy was derived by comparing the observed global colours to predictions of galaxy evolution models  (GEMs). The GEMs are based on the GALAXEV library (Bruzual \\& Charlot 2003a, 2003b). Since the main SF characteristic in dwarf galaxies is that it is generally proceeding in bursts, we selected to compare our observational results to model predictions of stellar populations formed in bursts. We used the 26 default models pre-calculated in GALAXEV using the Padova 1994 evolutionary tracks: 13 models use the Salpeter (1955) IMF and 13 the Chabrier (2003) IMF. No significant differences were found between the results obtained with either of the two IMFs. The comparisons were done for the optical colours [(U-B), (B-V), (V-R) and (V-I)], as well as for the (H$\\alpha$-V) colour, which is produced in a procedure described below, and for the predicted (2200{\\AA}-V) colour assumed to be identical to the (GALEX NUV-V) colour. %  Using the H$\\alpha$-V colour enhances the fit reliability since the line emission is determined by the amount of young massive stars in a galaxy. The models predict the number of LyC photons and we used this and the V magnitude to derive a distance-independent colour for each galaxy. The H$\\alpha$ luminosity is related to the LyC photon flux (Osterbrock 1989): \\begin{equation} N_{c} = 7.43 \\times 10^{11}L(H\\alpha), \\end{equation} where $L(H\\alpha)$ is in erg/s, which leads to: \\begin{equation} H\\alpha-V = 129.8 -2.5 \\times log[N_{c}] - V_{abs}, \\end{equation} where both $log[N_{c}]$ and $V_{abs}$ [mag] are calculated in the model.  All models use standard lower and upper mass cutoffs of $m_{L}=0.1M_{\\odot}$ and $m_{U}=100M_{\\odot}$, and differ in metallicity, which ranges from 0.0001 to 0.05. The models calculate spectra and galactic colours at 220 time steps, from t=0 to t=20 Gyr. They were used to create a colour data base for each time step and for each metallicity. Since the models describe a single generation of stars formed in an instantaneous SF burst that only ages after being formed, a script was written to find the best $\\chi^{2}$ fit of any combination of colours and burst times from the data base. The minimal $\\chi^{2}$ criterion calculated by the script is given by: \\begin{equation} {\\rm Min}(\\chi^{2}) = {\\rm min} [\\sum (\\Delta colour)^{2}/err^{2}] / (N_{colours} - N_{bursts}). \\end{equation} where $\\Delta colour$ is the difference between each measured colour and the linearly-combined model-produced colour, $err$ is the error of the measured colour, and the sum refers to the measured colours (NUV-V, U-B, B-V, V-R, V-I and H$\\alpha$-V). $N_{colours}$ is the number of available colours and $N_{bursts}$ is the number of bursts in the allowed combination.  Here we used primarily a combination of two bursts and the variables are the percentage of light produced by each population, the metallicity of this population, and the time of the burst. Since we fit essentially five predicted variables (fraction of light from the first burst in the observed brightness, the times of each of the two bursts, and the metallicity of each burst) to five measured colours [(U-B), (B-V), (V-R), (R-I) and (H$\\alpha$-V)], the fit should be, in principle, fully determined. For cases where we have also (NUV-V), the fit to two bursts should be overdetermined. Instances where we lack U-band photometry cause the fit to be under-determined.  For some cases we were forced to assume more than two SF bursts; it is probable that these fit results are not unique. The script outputs not only the best fit, but also the next nine fits to determine the spread of the multiple solutions. The dispersion of the nine additional results around the best fit, in terms of the burst times, burst weight, and metallicities, was used as the error estimate.  The dispersion around the best fit was typically 1-3 Myr for the very recent bursts, whereas for the earlier bursts that occurred a few Gyr ago it was $\\lesssim$1 Gyr. For very ``old'' bursts, $\\sim$10 Gyr, the dispersion is high and the time resolution is low; we therefore indicate that these SF bursts took place at least 10 Gyr ago without specifying exactly when. In addition, there are sometimes two different possibilities described by the 10 best fits, which have similar $\\chi^{2}$ values and we list both options below. The dispersion of the relative weights of the bursts (fraction of light measured now) is $\\sim10\\%$ around the best fit values.  The least constrained variable in the fitting procedure is the metallicity, since the observed colours can be constructed from bursts at two specific times, but with different metallicities. For a wide range of observed colours, different metallicity models shows differences smaller than the observational error. For this reason, the best fit sometimes indicates a galaxy metallicity decreasing with time and such apparently non-physical solutions are retained, since they could hint at an influx of extragalactic low-metallicity gas. ",
        "conclusions": "The SFH of the galaxies was derived by comparing the measured galactic colours to GEMs, and the results are summarized in Table \\ref{T2}. Most sample galaxies were apparently formed more than 10 Gyr ago, with some possibly forming somewhat later (6-10 Gyr ago). All galaxies, with the exception of AGC 122834, experienced a recent SF burst at similar epochs, $\\sim$1 to a few tens of Myr ago. This could be interpreted as a result of interactions that were at least partially responsible for the SF bursts, or as the influence of an SF mechanism that might be at work synchronizing the SF in separate and non-interacting galaxies. AGC 122834 does not seem to have experienced a significant recent burst; this could imply that it is not a member of the NGC 784 group, or that it is sufficiently distant from the other galaxies not to be affected by whatever mechanism is synchronizing the SF in the other objects.  \\begin{figure*}[ht] \\centering{ \\includegraphics[width=12cm]{F_SFRVsMHI.eps} \\caption{The relation between SFR and HI mass found for the objects studied here is similar to that measured in other environments for late-type galaxies.} \\label{fig:SFRHI}} \\end{figure*}   \\begin{figure*}[ht] \\centering{ \\includegraphics[width=12cm]{F_SFRVsMB.eps} \\caption{Relation between SFR and absolute B magnitude. A similar relation was found for galaxies in the nearby Universe (e.g., Karachentsev \\& Kaisin 2007)} \\label{fig:SFRVsmB}} \\end{figure*}   \\begin{figure*}[ht] \\centering{ \\includegraphics[width=12cm]{F_SFRVsD.eps} \\caption{SFR vs. distance from the main galaxies. In this plot we consider the distances of each galaxy from the main galaxy in its group.} \\label{fig:SFRvsDIST}} \\end{figure*}  Dwarf galaxies are common in the immediate vicinity of larger galaxies (Miller 1996, Cote et al.\\ 1997). Some are believed to form as results of strong interactions or mergers between larger galaxies (e.g., Hunsberger et al.\\ 1996, Bournaud \\& Duc 2006) and can also be explained by hierarchical clustering scenarios that sometimes include supernovae winds (e.g., Dekel \\& Silk 1986, Nagashami \\& Yoshii 2004). Models can properly explain their observed colours only if the stars in these galaxies were formed in short bursts separated by long quiescence periods (e.g., Tosi et al.\\ 1991 and references therein). The recent bursts could possibly be caused by the accretion of gas retained in the dark halos of these dwarfs (Dekel \\& Silk 1986) or collected from intergalactic space (see below). %    We tested correlations of the individual SFRs against various parameters. Figure~\\ref{fig:SFRHI} shows the relation between the SFR and the HI mass. The galaxies, though belonging to the same groups, show a variety of SFRs; this was found for other galaxy groups (e.g., Karachentsev \\& Kaisin 2007). We found a linear correlation of $log[SFR]$ with $log[HI_{mass}]$, which implies that the more material is available for SF, the more intense the SF would be.  Kennicutt (1998), Taylor (2006) and others found that irregular and spiral galaxies follow the relation [SFR]$\\propto M_{HI}^{1.4}$ (see also Karachentsev \\& Kaisin 2007). The galaxies here, as Figure \\ref{fig:SFRHI} shows, follow a relation with a similar slope of 1.3$\\pm0.15$. %  We also found that $log[SFR]$ correlates with M$_B$, both derived assuming only Hubble expansion; see Figure~\\ref{fig:SFRVsmB}). This indicates that the more intense the SF is, the more significant is the existing, relatively young, stellar population. Karachentsev \\& Kaisin (2007) found a similar result for 150 galaxies in the local Universe where the relation was [SFR]$\\propto$L$_{B}$. The galaxies in our sample follow a similar linear relation (Figure \\ref{fig:SFRVsmB}) with a slope of $\\sim$-0.4$\\pm$0.04, which corresponds to [SFR]$\\propto L_{B}^{1\\pm0.1}$, similar to that found by Karachentsev \\& Kaisin.     The SFR behaviour vs. the projected distance from the main galaxies of each group is shown in Figure~\\ref{fig:SFRvsDIST}. The underlying assumption here is that the DGs might be formed in tidal interactions among the major galaxies in each group, thus the further away a DG would be from its ``main progenitor'' major galaxy, the weaker should its SF process be. It is difficult to evaluate the role of interactions in the triggering of SF in these groups, primarily since we do not have accurate 3D distances. The SFR is highest in the interacting pair NGC 672/IC 1727. The Sdm galaxies NGC 784 and UGC 1281, which might also be interacting, show relatively high SFR values as well. However, as Figure~\\ref{fig:SFRvsDIST} shows, the six other DGs do not show a steadily decreasing SFR with (projected) distance from their ``main progenitors'', but rather a $\\sim$uniform and low SF of $\\sim10^{-4}$ M$_{\\odot}$/yr. More isolated galaxies (that are also dIrr) show  independently high SFRs.   We interpret this behaviour, along with the similar and recent SF burst times, as possible evidence for synchronicity in star formation in galaxies formed along a dark matter (DM) filament threading a nearby void. Brosch et al.\\ (2004) found that galaxy interactions in dIrr galaxies are probably not a primary SF trigger and this seems to be the case here. The recent SF bursts in the dIrr galaxies of our sample, all taking place in the last few tens of Myrs, could be caused by accretion of extragalactic gas onto most of the galaxies. The accreted gas then forms stars upon collapse onto the disks. The lack of dependence of the SFR on the (projected)  distance from the brightest and most massive galaxies, and the timing similarity of the recent bursts, show that the potential main disturbers have minimal or nil influence on the dIrr galaxies in their vicinity. Another element, such as the postulated influx of extragalactic gas, could probably be the trigger of the synchronous SF.   At this point one remark is in order regarding our modeling the SF process as a collection of instantaneous bursts, separated by long quiescent intervals. The justification for this was that this SF mode is the one recognized to take place in dwarf galaxies, and all the objects here are dwarfs. However, if we examine the HI depletion time due to the star formation, by dividing the total HI mass of each object by its calculated SFR, we find that all objects could sustain SF at an almost constant level for long periods indeed, of order 10$^9$ yrs. Note that we have no information about the HI distribution; this could be very extended given the size of the ALFALFA beam, with only part of the HI taking part in the SF process. The long time needed to exhaust the HI reservoir could indicate that a quasi-continuous and low-level SF process might also explain the results. We emphasize, however, that the important issue here is that almost all objects show present-day SF.  The described behaviour and the proposed interpretation are, after all, not that surprising. It has been known that only strong interactions at short separations between galaxies of comparable masses enhance SF, and that the SFR increases as the strongly interacting galaxies come closer (e.g., Icke 1985). This could perhaps be the case in the NGC 672/IC 1727 (VV338) pair, but not in most of the galaxies studied here. In addition, it has been known that fairly isolated galaxies (in relatively denser environments) show generally higher SFR the more isolated they are. This could perhaps be the case for NGC 855 and UGC 1561. %  Kere\\v{s} et al. (2005) showed, using hydrodynamic simulations, that low-mass galaxies with baryonic mass M$_{gal}\\leq10^{10.3}$M$_{\\odot}$ (or halo mass M$_{halo}\\leq 10^{11.4}$M$_{\\odot}$) can accrete intergalactic gas in a ``cold mode'', with the accreted gas colder than 10$^5$K thus capable of directly forming stars. The cold accretion in the simulations is often directed along DM filaments, allowing galaxies to efficiently draw gas from large distances. Such DM filaments form in N-body simulations such as those presented by Hahn et al. (2007), where it was also shown that that haloes in filaments are more oblate than cluster halos at high masses and that haloes in filaments tend to have their spin vectors pointing along the filaments.  The prediction regarding the spin alignment result of Hahn et al. (2007) is that the major axes of disky galaxies in the filament discussed here should be approximately perpendicular to the line along which the galaxies appear to be arranged. We checked this by estimating the position angles of the major axes of the galaxies (listed in Table~\\ref{gal}) and comparing them to the general direction of the projected filament. The distribution of PA values shows two preferred values: one, with seven galaxies is centered at PA$\\simeq$305$^{\\circ}\\pm12^{\\circ}$. The other peak is broader and is centered at PA$\\simeq$34$^{\\circ}\\pm31^{\\circ}$. The general direction of the galaxy filament, using the same convention of measuring PAs, is $\\sim$303$^{\\circ}$. The conclusion is therefore that the Hahn et al. (2007) prediction regarding the correlated galaxy spins is fulfilled in a mixed-up way for this nearby galaxy alignment, with seven objects having their spin approximately aligned with the general direction of the filament and the other objects with their spin axes approximately perpendicular to the filament direction.  Dekel \\& Birnboim (2006) showed that low-mass galaxies are built by cold gas streams in haloes below a critical shock-heating mass M$_{shock}\\leq10^{12}$M$_{\\odot}$, not by intergalactic gas heated by a virial shock. This accretion mode yields efficient star formation in low-mass galaxies. The only ingredient then needed to understand the galaxies studied here is to assume that the phenomenon can take place not only at z$\\geq$2, but also at z$\\approx$0. We could, therefore, assume that most of the 14 galaxies studied here do line up along a DM filament in a nearby void, and that they all exhibit a present-day synchronous star formation burst triggered by cold gas accretion from intergalactic space that is being focused by the DM filament.     We presented U, B, V, R, I, H$\\alpha$ and NUV photometric measurements of 14 fairly isolated  galaxies in the local Universe. The galaxies are in the same sky area, most are arranged on a linear configuration, and have radial velocities $\\leq$610 km sec$^{-1}$. 11 of these galaxies, and one HI cloud with no optical counterpart, were previously assigned to two groups of galaxies. All visible galaxies are ``dwarfs'' with M$_B\\geq$--18 and can be morphologically classified as ``very late type'', including the dominating objects in each group. Three other galaxies were often considered to be field galaxies but are similar in absolute magnitude and classification to the other objects. The objects appear to form a single kinematically well-behaved ensemble that does not separate naturally into two groups. We derived the SFR and SFH of each object using photometry in the various bands and examined these properties in the context of the galaxy environment.  The galaxies show relations between the SFR and the HI mass, or m$_B$, known for field galaxies. The SFHs imply that similar low-intensity SF bursts took place in most galaxies a few Myr ago, despite their being spread along a $\\geq$0.5 Mpc long feature. The behaviour of the SFRs, which do not decline steadily with increasing distance from the brightest galaxies in each group,  indicates that interactions do not affect the SF in these galaxies to the extent one would expect from strongly interacting galaxies. In particular, no signs of strong gravitational interactions, such as tidal tails or global disturbances, are observed. The main galaxies do appear to be interacting but on a one-to-one basis and do indeed show enhanced SF. Other dIrr galaxies in the same vicinity show only low SFRs. More isolated galaxies show independently high SFRs.  We propose that the observational evidence argues in favor of interpreting the galaxies as located on a DM filament that is itself located in a low-galaxy-density region, and is accreting intergalactic cold gas focused by the filament. We are therefore witnessing in the classical NGC 672 and NGC 784 groups of relatively bright and nearby late-type galaxies the basic phenomenon of hierarchical clustering, the direct formation and growth of small galaxies out of intergalactic gas accreted on a dark matter ``backbone''. This nearby galaxy collection offers, therefore, an ideal opportunity to study the phenomenon of hierarchical clustering in significant detail.  One possible conclusion from our proposed interpretation is that the gas being accreted could be observable in the vicinity of the galaxies; a way to detect it would be through very low column density detection of Lyman $\\alpha$ absorption lines at the redshifts of the galaxies in spectra of background QSOs. Another would be to search for diffuse HI to much lower limits than achieved during the ALFALFA survey; it is unlikely that an extension of the ALFALFA scans will bring useful returns and perhaps this is a project that should wait for the entrance in operation of SKA."
    },
    "0808/0808.3573_arXiv.txt": {
        "abstract": "{Near-infrared (hereafter NIR) data may provide complementary information to the traditional optical population synthesis analysis of unresolved stellar populations because the spectral energy distribution of the galaxies in the 1-2.5\\,$\\mu$m range is dominated by different types of stars than at optical wavelengths. Furthermore, NIR data are subjected to less absorption and hence could constrain the stellar populations in dust-obscured galaxies.} {We want to develop observational constraints on the stellar populations of unresolved stellar systems in the NIR.} {To achieve this goal we need a benchmark sample of NIR spectra of ``simple'' early-type galaxies, to be used for testing and calibrating the outputs of population synthesis models. We obtained low-resolution (R$\\sim$1000) long-slit spectra between 1.5 and 2.4\\,$\\mu$m for 14 nearby early-type galaxies using SofI at the ESO 3.5-m New Technology Telescope and higher resolution (R$\\sim$3000) long-slit spectra, centered at the Mg{\\sc I} at $\\sim$1.51\\,$\\mu$m for a heterogeneous sample of 5 nearby galaxies observed with ISAAC at Antu, one of the 8.2-m ESO Very Large Telescope.} {We defined spectral indices for CO, Na{\\sc I}, Ca{\\sc I} and Mg{\\sc I} features and measured the strengths of these features in the sample galaxies. We defined a new global NIR metallicity index, suitable for abundance measurements in low-resolution spectra.  Finally, we present an average NIR spectrum of an early-type galaxy, built from a homogenized subset of our sample.} {The NIR spectra of the sample galaxies show great similarity and the strength of some features does correlate with the iron abundance [Fe/H] and optical metal features of the galaxies. The data suggest that the NIR metal features, in combination with a hydrogen absorption feature may be able to break the age-metallicity degeneracy just like the Mg and Fe features in the optical wavelength range.} ",
        "introduction": "State-of-the-art space-based instrumentation can only resolve galaxies from the Local Group well and in the more distant objects we usually can see only the tip of the red giant branch stars which is rarely sufficient for population analysis. This leaves us with the difficult task of trying to recover the stellar population of unresolved galaxies from their integrated properties. The complex mix of stellar populations found in most of them usually makes it possible to constrain only the most recent generation of stars. However, such properties as the stellar kinematics and the present-day metal content are dominated by the overall star formation history, and together with the well known age-metallicity degeneracy \\citep{fab73,oco86,wor94} they often lead to non-unique solutions for the stellar populations.  In this respect, the Lick/IDS system of indices pioneered by \\citet{bur84} and \\citet{fab85}, and developed further by \\citet{tra98, tra00, tra05}, has been particularly successful for interpreting the integrated optical light of galaxies.  However, new constraints are necessary to interpret more complicated systems and one possibility is to widen the spectral range towards the near-infrared (NIR) because the light in different wavebands is dominated by different populations of stars. The NIR passbands are dominated by light from older and redder star and therefore offer us the possibility to study other stellar populations than is possible with optical spectra alone. In addition, abundance determinations through optical spectroscopy are not possible for heavily reddened evolved stellar populations such as dusty spheroids or some bulge globular clusters hidden by dust \\citep[e.g.,][]{ste04}. NIR spectroscopic observations could overcome these problem because the extinction in the $K$-band is only one-tenth of that in the $V$-band.  Most of the previous work at NIR was focused on either active galactic nuclei (AGNs) or objects with very strong star formation, including recent surveys of ultra-luminous infrared galaxies \\citep{gol95,mur99,mur01,bur01}, luminous infrared galaxies \\citep{gol97,reu07}, starbursts \\citep{eng97,coz01}, Seyfert galaxies \\citep{iva00,sos01,reu02,boi02}, LINERs \\citep{lar98,alo00,sos01}, and interacting galaxies \\citep{van97,van98}.  Relatively few NIR spectroscopic observations exist for ``normal'' galaxies. Only \\citet{man01} provided low-resolution ($R\\sim400$) template spectra for galaxies of different Hubble types, including some giant ellipticals. Such data, together with the corresponding analysis is a necessary first step towards developing a system of spectral diagnostics in the NIR because well-understood galaxies with relatively simple star forming history will allow us to tune the NIR population synthesis models. Recently, \\citet[][hereafter S08]{sil08} studied the stellar populations of eleven early-type galaxies in the nearby Fornax cluster by means of $K$-band spectroscopy.  A few prominent NIR features were first studied by \\citet{ori93} who demonstrated that they represent a superb set of indicators for constraining the average spectral type and luminosity class of cool evolved stars. This conclusion was later confirmed by \\citet{for00} and \\citet{iva04}. Furthermore, the same NIR features appear to be promising abundance indicators \\citep{fro01,ste04}.  Here we describe two new data sets of high quality NIR spectra of ellipticals/spirals designed to provide a benchmark for future NIR studies of unresolved galaxies: {\\it (i)} low-resolution spectra covering the range from 1.5 to 2.4\\,$\\mu$m that include strong features such as CO, Na{\\sc I} and Ca{\\sc I} that are traditionally studied, and {\\it (ii)} moderate resolution spectra around the Mg{\\sc I} absorption feature at 1.51\\,$\\mu$m. This is the second strongest Mg feature in the $H$- and $K$-band atmospheric windows (after the feature at 1.71\\,$\\mu$m) and at zero redshift it is located in a region of poor atmospheric transmission, making it difficult to observe in stars. However, the redshift of external galaxies moves it into a more transparent region \\citep{iva01}, just the opposite of the 1.71\\,$\\mu$m Mg{\\sc I} which becomes affected by the red edge of the $H$-band atmospheric window.  We are only few years from the launch of the James Webb Space Telescope \\citep{gar06}, a space-based infrared telescope that will have unprecedented capabilities. Therefore, in the near future the application of NIR spectroscopy to the study of galaxy properties will be limited not by lack of data but by our understanding of spectral features at these wavelengths. Improving the characterization of the NIR indices is a timely step in this direction.  The paper is organized as follows. The sample selection is discussed in Sect.\\,\\ref{sec:Sample}. The NIR spectroscopic observations and data reduction are described in Sect.\\,\\ref{sec:Obs}.  The definition of the new NIR spectral indices and their measurements are given in Sect.\\,\\ref{sec:Indices}. Results are discussed in Sect.\\,\\ref{sec:Discus} and summarized in Sect.\\,\\ref{sec:Results}. ",
        "conclusions": "}  \\subsection{General appearance of the spectra}  The spectra of our sample galaxies appear qualitatively similar in most of the NIR features except for a large spread in the Na{\\sc I} values is present.  This similarity is not surprising because we selected mostly giant ellipticals and spheroids, which are all relatively metal rich, have no significant recent star formation, and only weak nuclear activity, if any. This conclusion agrees with \\citet{man01} who demonstrated that galaxies within the same Hubble type have nearly identical NIR spectra. On the other hand, the large spread of NaI, with respect to the observational errors, appears to be real, and suggests variety of star formation and enrichment histories among the galaxies in our sample. Two of the three galaxies with systematically weaker NaI exhibit strong H$\\beta$, indicating that the weak NaI might be related to the presence of younger stellar populations (the third galaxy lacks optical spectroscopy).  The strongest features are the CO absorption bands in both the $H$ and $K$-band atmospheric windows. They originate in K and M stars, as can be seen from the libraries of stellar spectra described in \\citet{lan92}, \\citet{ori93}, \\citet{dal96}, \\citet{mey98}, \\citet{wal97} and \\citet{for00}. A number of weaker metal absorption features are visible as well: Si{\\sc I} at 1.589\\,$\\mu$m, Mg{\\sc I} at 1.711\\,$\\mu$m, Na{\\sc I} at 2.206 and 2.209\\,$\\mu$m and Ca{\\sc I} at 2.261, 2.263, and 2.266\\,$\\mu$m. They are also present mostly in cool stars \\citep[i.e.][]{kle86,ori93,wal97,for00,mey98}. As mentioned above, we only consider the redder features with $\\lambda\\geq$1.65$\\mu$m.  Not surprisingly, our galaxies show no Fe{\\sc II} at 1.644\\,$\\mu$m or H$_2$ lines line emission, usually associated with supernova activity and supernova related shocks, respectively. There is also no Br$\\gamma$ emission.  \\subsection{Behavior of the NIR spectral indices}\\label{sec:Behaviour} Evolutionary synthesis models \\citep[i.e.][]{wor94} are the usual method to interpret the behavior of spectral features in galaxies, but the simple comparison of these features with the galaxy parameters can be informative too.  For example, the optical indices, Mg$_2$\\,, $<$Fe$>$ and \\Hb show tight correlation with the central velocity dispersion \\citep[e.g.][]{ter81,ber98,meh03,mor08}, suggesting that the chemical and the dynamical evolution of ellipticals are intertwined.  The EW of our NIR spectral indices, the iron abundance and metallicity are plotted versus the central velocity dispersion in Fig.\\,\\ref{fig:VelDips}. The loci of the S08 data (see their Fig. 13) are shown with dashed lines. Among the galaxies, NGC\\,5077 (\\#12) and NGC\\,6909 (\\#14) possess very weak Na{\\sc I}, Ca{\\sc I} lines, and their relative errors are significant. The correlations between the NIR indices and velocity dispersion for the sample galaxies are plotted in Fig\\,\\ref{fig:VelDips}; the fit do not consider NGC\\,4281 (\\#4), NGC\\,5077 (\\#12) and NGC\\,6909 (\\#14) because of the reasons explained further on.  The correlation with CO shows a similar slope as the one found by S08 but with a different zero-point. The different definition of the index and especially the position of the continuum bandpass can lead to a systematic variation of the EW values. This effect seems to be more evident in the CO band, whose equivalent width is measured with only extrapolated continuum blue-wards of the feature.  The agreement with S08 is closer in the Na{\\sc I} versus $\\sigma$ relation although our data show bigger scatter. Without pretending to trace a general conclusion we found relatively large NaI values compared to stars and something similar was also found by S08. The iron abundance [Fe/H] and the total abundance [Z/H] only show very weak correlations. Active galaxies are clustered at the high end of the velocity dispersion distribution which is understandable because their velocity dispersion measurement may be affected by the black hole and the active nucleus. Note that our highest velocity dispersion galaxy is a Type 2 Seyfert and the CO band may suffer some dilution, as discussed in \\citet{iva00}.   \\begin{figure} \\includegraphics[width=9truecm]{vd02a.eps} \\caption{The equivalent widths of the NIR spectral indices, iron abundance [Fe/H], and metal abundance [$Z$/H] in the sample galaxies as a function of the central velocity dispersion. The labels correspond to the row numbers in the Table \\ref{tab:GalProperties1}. The open and filled circles refer to quiescent and active galaxies, respectively. The small and large circles refer to galaxies with an age between 5 and 10\\,Gyr and $>$10\\,Gyr, respectively. Galaxies with no known ages were assumed to be old. The dashed lines corresponds to the relations find by S08. The solid line in each panel represents the linear regression (y = ax + b) through all the data points except for \\#4, \\#12 and \\#14. The Pearson correlation coefficient (r) and the results of the linear fit are given. Typical error bars are shown.} \\label{fig:VelDips} \\end{figure}   The EW of our NIR spectral indices are plotted against each other in Fig.\\,\\ref{fig:IR_IR} and compared to the relations for cluster stars and solar neighborhood giants by S08 and \\citet{dav08}, respectively. Three galaxies -- NGC\\,4281 (\\#4), NGC\\,5077 (\\#12) and NGC\\,6909 (\\#14) -- are above the relation of S08 for the cluster stars, which hints at different chemical enrichment history with respect to the rest of the sample. Unfortunately the literature lacks much data about these objects and although NGC\\,6909 has the lowest velocity dispersion in our sample, these galaxies do not stand out in any respect, including in the optical Mg$_2$ versus $\\sigma$ diagram. Further investigation of these galaxies is necessary. The rest of our objects populates a locus that follows similar trend as the galaxies of S08. In the Ca{\\sc I} versus CO plot we see a correlation with significant scatter, and with a different slope than in S08.  Finally, the Mg{\\sc I} versus CO plot is dominated by scatter.  Galaxies with traces of nuclear activity, evident in other wavelength ranges than the NIR, do not seem to separate from the rest of the sample, and they show no traces of emission lines. Therefore, the contribution of their AGNs is negligible with respect to the rest of the galaxy.   \\begin{figure} \\includegraphics[width=9truecm]{IR_IR02a.eps} \\caption{The equivalent widths of the NIR indices plotted against each other. The symbols are the same as in Fig.\\,\\ref{fig:VelDips}. The dashed lines correspond to the relations found by S08. The dotted line corresponds to the relation by \\citet{dav08} for giant stars in the solar neighborhoods. All these relations are plotted in their observed range. The solid line represents the linear regression (y = ax + b) through all the data points except for \\#4, \\#12 and \\#14. The Pearson correlation coefficient (r) and the results of the linear fit are given. Typical error bars are shown.} \\label{fig:IR_IR} \\end{figure}   Figure\\,\\ref{fig:IR_FeH} shows the EW of the NIR spectral indices versus the Mg$_2$ measurements from \\citet{ben93}. Na{\\sc I}, and to a lesser extent Ca{\\sc I} and CO, do show trends with respect to Mg$_2$, with NGC\\,4281 (\\#4), NGC\\,5077 (\\#12) and NGC\\,6909 (\\#14) standing out. The correlation of Na{\\sc I} with both $\\sigma$ and Mg$_2$ suggests that the Na{\\sc I} feature is dominated by the stellar photosphere rather than by the interstellar medium.   \\begin{figure} \\includegraphics[width=9truecm]{IR_FeH02a.eps} \\caption{The equivalent widths of the NIR indices plotted against the Mg$_2$ indices measured by \\citet{ben93}. The symbols are the same as in Fig.\\,\\ref{fig:VelDips}. The solid line represents the linear regression (y = ax + b) through all the data points except for \\#4, \\#12 and \\#14. The Pearson correlation coefficient (r) and the results of the linear fit are given. Typical error bars are shown.} \\label{fig:IR_FeH} \\end{figure}   Given that [Fe/H] measurements are not available for the entire sample and that Mg$_2$ is not representative of the total chemical abundance of a galaxy because of the varying abundance ratios, we attempted to create a combined Iron-and-$\\alpha$-element index similar to that defined by \\citet{gon93} and more recently by \\citet{tom03}: \\begin{equation} {\\rm [MgFe]'} = \\sqrt{{\\rm Mg}b\\times (0.72 \\times {\\rm Fe}5270+0.28 \\times {\\rm Fe}5335)} \\end{equation} Such a combined indicator is expected to decrease the effect from the varying $\\alpha$/Fe ratio. Since not all the components of this indicator are available we directly substitute the iron and magnesium indices by defining a new indicator: \\begin{equation}\\label{Eqn:FeMg2} {\\rm [MgFe]''} = \\sqrt{{\\rm Mg}_2\\times {\\rm Fe}5335} \\end{equation} The results are shown in Fig.\\,\\ref{fig:IR_MgFe}a-c.  The NaI index of NGC\\,5077 (\\#12) and the CO index of NGC\\,6909 (\\#14) deviate from the main loci of the other galaxies.  To address this issue we plotted the optical indices of these galaxies versus the combined [MgFe]$'$ index (Fig.\\,\\ref{fig:IR_MgFe}d). This is an analog of the typical plot \\citep[i.e.][]{wor94} that allows to disentangle the age-metallicity degeneracy: the inverse \\Hb index is roughly proportional to age while the [MgFe]$'$ is dominated by metal abundance. The two galaxies exhibit strong \\Hb which suggests that they are dominated by populations of 3\\,Gyr or younger (see for example Fig. 1 in S08). NGC\\,5077 and NGC\\,6909 are also well separated from the rest of the galaxies on the \\Hb versus Na{\\sc I} plot (Fig.\\,\\ref{fig:IR_MgFe}e) which leads us to the conclusion that the NIR indices can be used to create a similar diagnostic plot, but measuring the Br$\\gamma$ feature requires better quality data than the ones described here.\\\\ We investigated the behavior of the NIR indices in relation to the H and K-band magnitude and the H-K color but no clear correlations were found.   \\begin{figure} \\includegraphics[width=9truecm]{IR_MgFe02a.eps} \\caption{(a-d) The equivalent widths of the NaI, CaI, CO and H$\\beta$ indices plotted against the [MgFe]$''$ index. (e) The equivalent widths of the H$\\beta$ index as a function of the NaI index. The symbols are the same as in Fig.\\,\\ref{fig:VelDips}. The dashed lines correspond to the relations found by S08. The solid line represents the linear regression (y = ax + b) through all the data points except for \\#4, \\#12 and \\#14. The Pearson correlation coefficient (r) and the results of the linear fit are given. Typical error bars are shown.} \\label{fig:IR_MgFe} \\end{figure}   \\subsection{Combined NIR metal index}  Ground based NIR spectroscopy is much more time consuming than the corresponding optical observations -- higher and variable background and detectors with worse cosmetics often require one to sacrifice either resolution or signal-to-noise to obtain the data in a reasonable amount of time. To alleviate this problem, we defined a combined spectral index of all the major $K$-band metal features: \\begin{equation} \\langle {\\rm CONaCa} \\rangle = ( {\\rm CO} + \\ion{Na}{i} + \\ion{Ca}{i} ) / 3.0 \\end{equation} Various weighting schemes were tried to minimize the scatter of the basic relations (Fig.\\,\\ref{fig:CONaCa}). However, the simple average yielded the tightest relations. The carbon and oxygen are $\\alpha$ elements, while the sodium and calcium originate in both high and low mass stars. The average value of the CO index for the sample galaxies is $\\sim$15\\,\\AA, the Na{\\sc I} is $\\sim$4.3\\,\\AA\\, and the Ca{\\sc I} is $\\sim$2.7\\,\\AA\\, which means that at least 2/3 of the new index is dominated by metals produced mostly in high mass stars, i.e. early in the history of the elliptical galaxies. The other implication is that a relatively limited amount of recent star formation could affect the new index more than it would an iron peak dominated index. The lack of iron peak features in the NIR is well known as it was pointed out in \\citet{iva01}.   \\begin{figure} \\includegraphics[width=9truecm]{nir02a.eps} \\caption{The equivalent widths of the newly defined $<$CONaCa$>$ as a function of the central velocity dispersion (a), Mg$_2$ (b), H$\\beta$ (c), and [FeMg]$''$ (d) indices. The symbols are the same as in Fig.\\,\\ref{fig:VelDips}. The solid line represents the linear regression (y = ax + b) through all the data points except for \\#4, \\#12 and \\#14. The Pearson correlation coefficient (r) and the results of the linear fit are given. Typical error bars are shown.} \\label{fig:CONaCa} \\end{figure}   The new index improves the correlations. For example, the Pearson correlation coefficient is $\\sim$10\\% higher for the $\\langle {\\rm CONaCa} \\rangle$ vs. Mg2 and even $\\sim$20\\% higher for the $\\langle {\\rm CONaCa} \\rangle$ vs. [MgFe]'', with respect to the respective correlations where only one IR index is used. We derived these relations using only the galaxies with known parameters, i.e. excluding the poorly studied NGC\\,4281(\\#4) and, NGC\\,5077(\\#12) and NGC\\,6909(\\#14) which systematically show peculiarities with respect to the bulk of our sample.  \\subsection{Average NIR spectrum of the sample galaxies} Studies of composite stellar systems (i.e. galaxies hosting an AGN) often need to subtract the contribution of the underlying galaxy. This prompted us to create an average spectrum of the galaxies in our sample. We used a homogenized subset of eight galaxies, excluding the objects with young populations. We also excluded those with low signal-to-noise ratio. The average age of the galaxies used to create the composite spectrum is 9$\\pm$2\\,Gyr (the median is 9.5\\,Gyr) and the average [Fe/H] is 0.17$\\pm$0.13 (the median is 0.16).  The composite spectrum is shown in Fig.\\,\\ref{fig:AverSpec} and the values of the spectrum are listed in Table\\,\\ref{tab:AverSpec}, together with the r.m.s. values per Angstrom.  We included the $H$-band section because the artifacts caused by the spectral type mismatch of the standards are minimized by the averaging of galaxies observed at different redshifts.   \\begin{figure} \\includegraphics[width=9truecm]{comb01.eps} \\caption{Average $H$- (a) and $K-$band (b) spectrum of sample galaxies normalized to unity at 22000\\ \\AA\\, (thick line). The thin lines correspond to the $\\pm1\\sigma$ confidence levels. Some of the more prominent spectral features are marked.} \\label{fig:AverSpec} \\end{figure}   \\begin{table} \\caption{Average spectrum of the sample galaxies observed at NTT in arbitrary units and normalized to unity at 22000\\,\\AA. The regions with zero r.m.s. at the beginning and the at the and are covered only by one or two spectra. The full table is given only in the electronic version of the journal.} \\label{tab:AverSpec} \\begin{center} \\begin{tabular}{ccc} \\hline \\hline $\\lambda$ (\\AA) & F$_\\lambda$ & r.m.s. \\\\ \\hline 14944 & 2.217 & 0.000 \\\\ 14945 & 2.384 & 0.000 \\\\ 14946 & 2.384 & 0.000 \\\\ 14947 & 2.386 & 0.000 \\\\ 14948 & 2.130 & 0.366 \\\\ 14949 & 2.572 & 0.254 \\\\ 14950 & 2.575 & 0.251 \\\\ 14951 & 2.578 & 0.248 \\\\ 14952 & 2.583 & 0.245 \\\\ 14953 & 2.588 & 0.243 \\\\ 14954 & 2.595 & 0.241 \\\\ 14955 & 2.602 & 0.240 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}   \\subsection{The Mg{\\sc I} feature at 1.51\\,$\\mu$m}  \\citet{iva04} pointed out the possibility of using the Mg{\\sc I} feature at 1.51\\,$\\mu$m as an $\\alpha$-element abundance indicator in the NIR. With exploratory purposes we obtained spectra of a small and diverse set of galaxies (Table\\,\\ref{tab:log1} and Fig.\\,\\ref{fig:MgI_H_band}). To carry out quantitative analysis we defined an index (Table\\,\\ref{tab:bandDef}) and for the five galaxies we measured equivalent widths of 3.3, 5.0, 3.8, 3.9 and 4.5\\AA, respectively for Mrk\\,1055, NGC\\,1144, NGC\\,1362, NGC\\,4472 and NGC\\,7714. The typical uncertainty is $\\sim0.3$\\AA. No corrections for velocity dispersion were applied. Given the diverse nature of the objects, we refrain from drawing any conclusions but we note that the relation between the optical and the NIR Mg features is not straightforward because the only two Mg$_2$ values that we have from the literature are inversely proportional to our measurements for NGC\\,1362 and NGC\\,4472. Nevertheless, these observations prove that it is feasible to measure the NIR Mg{\\sc I} feature at 1.51\\,$\\mu$m and give a basis for future NIR synthetic spectral modeling.   \\begin{figure} \\includegraphics[width=9truecm]{spe04.eps} \\caption{Spectra of the galaxies observed with ISAAC in the region of the Mg{\\sc I} (1.51 $\\mu$m). The central wavelengths of the individual Mg{\\sc I} lines are shown with vertical solid lines.  The passbands for the feature and the adjacent continua are shaded. The spectra were normalized to unity and shifted vertically by 0.5 for display purposes.  } \\label{fig:MgI_H_band} \\end{figure}"
    },
    "0808/0808.2480_arXiv.txt": {
        "abstract": "We have determined Al, $\\alpha$, Fe--peak, and neutron capture elemental abundances for five red giant branch (RGB) stars in the Galactic globular cluster M10.  Abundances were determined using equivalent width analyses of moderate resolution (R$\\sim$15,000) spectra obtained with the Hydra multifiber positioner and bench spectrograph on the WIYN telescope.  The data sample the upper RGB from the luminosity level near the horizontal branch to about 0.5 mag below the RGB tip.  We find in agreement with previous studies that M10 is moderately metal--poor with [Fe/H]=--1.45 ($\\sigma$=0.04).  All stars appear enhanced in Al with $\\langle$[Al/Fe]$\\rangle$=+0.33 ($\\sigma$=0.19), but no stars have [Al/Fe]$\\ga$+0.55.  We find the $\\alpha$ elements to be enhanced by +0.20 to +0.40 dex and the Fe--peak elements to have [el/Fe]$\\sim$0, which are consistent with predictions from type II SNe ejecta. Additionally, the cluster appears to be r--process rich with $\\langle$[Eu/La]$\\rangle$=+0.41. ",
        "introduction": "Although few chemical analysis studies of M10 exist, the general consensus is that this cluster exhibits all of the classical characteristics observed in other Galactic globular clusters.  With a metallicity of [Fe/H]$\\approx$--1.5 (Kraft et al. 1995), M10 lies near the median metallicity distribution for halo globular clusters (Laird et al. 1988).  Small sample (N$\\la$15) analyses of red giant branch (RGB) stars in this cluster have revealed it to have [$\\alpha$/Fe]$\\sim$+0.30 and [el/Fe]$\\sim$0 for Fe--peak elements (Kraft et al. 1995; Mishenina et al. 2003).  These values are consistent with the current generation of M10 stars having been polluted by the ejecta of type II supernovae (SNe) without significant contributions from type Ia SNe.  While it has long been known that nearly all globular cluster giants show star--to--star variations of the light elements (A$\\la$27), the source of many of these anomalies has yet to be determined.  Numerous observations of globular cluster stars from the main sequence to above the RGB luminosity bump have revealed declining [C/Fe] and increasing [N/Fe] ratios as a function of increasing luminosity (e.g., see reviews by Kraft et al. 1994; Gratton et al. 2004; Carretta 2008).  These observations show clear evidence of CN--cycle products being brought to the surface and are a confirmation of first dredge--up predictions (Iben 1964).  Smith \\& Fulbright (1997) and Smith et al. (2005) have verified this trend in M10 as well as a CN band anticorrelation with [O/Fe] for stars at various RGB luminosities.  However, the large spread in [N/Fe] of about 1.0 dex found by Smith et al. (2005) in M10 stars may be evidence for primordial variations superimposed on in situ mixing.  The C and N abundance anomalies are known to exist in both globular cluster and field giants, but that likeness does not extend to the well documented O/Na, Mg/Al, and O/F anticorrelations and Na/Al correlation seen solely in globular cluster stars (e.g., Gratton et al. 2004).  These abundance relationships are clear signs of proton--capture nucleosynthesis, but where these processes are operating is still a mystery.  Kraft et al. (1995) examined the O/Na anticorrelations of 15 bright giants in M10 along with M3 and M13, which are all globular clusters of similar metallicity ([Fe/H]$\\approx$--1.5), because M10 and M13 have extremely blue horizontal branches (HB) but M3 has a uniform distribution of blue HB, RR Lyrae, and red HB stars.  The study showed that M10 appears to be an intermediate case in terms of O depletion and Na enhancement in that the average [O/Fe] is lower in M10 than in M3, but no M10 giants were super O--poor (i.e., [O/Fe]$<$--0.6), suggesting the process driving O depletion does not itself determine HB morphology.  In this paper we have examined five additional RGB stars in M10 that are located above the luminosity of the horizontal branch but below the RGB tip. We have derived Al, $\\alpha$, Fe--peak, and heavy element abundances to examine how M10 fits into context with other globular clusters of similar metallicity and HB morphology. ",
        "conclusions": "\\subsection{Al Abundances}  We have determined at least upper limits of [Al/Fe] for five giants with the cluster having $\\langle$[Al/Fe]$\\rangle$=+0.33 ($\\sigma$=0.19) and a full range of 0.50 dex.  Both the star--to--star dispersion and average [Al/Fe] ratios are in agreement with observations of other Galactic globular clusters of similar metallicity (e.g., Kraft et al. 1998; Sneden et al. 2004; Cohen \\& Mel{\\'e}ndez 2005; Johnson et al. 2005; Yong et al. 2005); however, the highest [Al/Fe] ratio found in our sample is about a factor of three smaller than the $>$+1.0 dex ratios observed in M3 and M13 (Pilachowski et al. 1996; Sneden et al. 2004; Johnson et al. 2005; Cohen \\& Mel{\\'e}ndez 2005), which possess similar metallicity and, in the case of M13, a similar HB morphology.  This may be due to our small sample size coupled with observations of stars well below the RGB tip, where additional Al enhancement due to extra in situ mixing may be operating (e.g., Denissenkov \\& VandenBerg 2003).  Kraft et al. (1995) found M10 to be an intermediate case between M3 and M13 with regard to the amount of O depletion and Na enhancement and therefore given the likely Na--Al correlation present in this cluster one would not expect [Al/Fe] values much greater than about +0.80 dex.  A complete list of our determined abundances for Al and all other elements is provided in Table 3.  It has been shown that [Fe/H] determinations based on Fe I lines in metal--poor stars suffer from larger LTE departure effects than their metal--rich counterparts because of overionization due to decreased UV line blocking (e.g., see review by Asplund 2005).  Correcting for this effect would drive the [Fe/H] abundance up, perhaps by as much as $\\sim$+0.30 dex at [Fe/H]=--3 (Th{\\'e}venin \\& Idiart 1999, but see also Gratton et al. 1999; Kraft \\& Ivans 2003), and thus decrease the derived [Al/Fe] ratio found here.  While a few NLTE studies for Al exist (e.g., Gehren et al. 2004; Andrievsky et al. 2008) finding offsets of order a few tenths of a dex, the actual Al NLTE correction for stars in the metallicity and luminosity regime studied here are mostly unknown.  Fortunately, our sample does not vary widely in either metallicity or luminosity and any NLTE corrections are likely to be very similar, suggesting at least the relative star--to--star dispersion is a real effect.  In Figure \\ref{f2} we compare abundances of various elements in M10 versus those in the similar cluster M12.  The [Al/Fe] abundances for both clusters are comparable and each displays a modest star--to--star dispersion.  Given that the scatter is about a factor of two larger than those observed in the Fe--peak and $\\alpha$ elements, it is likely that the Al distribution is real and not an artifact of observational uncertainty.  To see how M10 fits into the context of other Galactic globular clusters, we have plotted [Al/Fe] as a function of both horizontal branch ratio (HBR) and galactocentric distance (R$_{\\rm GC}$) for M10 and seven other clusters in Figure \\ref{f3}. The top panel suggests there is no significant relation between HBR and either the average [Al/Fe] ratio or the star--to--star dispersion.  However, it should be noted that Carretta et al. (2007) do find a relationship between the extent of O/Mg depletions and Na/Al enhancements and the maximum temperature of stars located on the zero--age HB.  The bottom panel may indicate a trend of increasing cluster average [Al/Fe] with increasing galactocentric distance; however, the sample size for each cluster varies between less than 10 to nearly 100 stars.  Consequently, M10 does not appear to exhibit anomalous [Al/Fe] ratios compared to other globular clusters.  \\subsection{$\\alpha$, Fe--Peak, and Heavy Elements}  Nearly all globular clusters with [Fe/H]$<$--1 have [$\\alpha$/Fe]$\\sim$+0.30 to +0.50, solar Fe--peak to Fe ratios, and are r--process rich (e.g., Gratton et al. 2004).  The star--to--star scatter present is usually $\\la$0.10 dex for the $\\alpha$ and Fe--peak elements and $\\sim$0.30--0.50 dex for the neutron capture elements, which is still significantly less than the 0.50--1.00 dex variations seen in light elements such as O, Na, and Al.  In M10 we find the expected enhancement and small star--to--star dispersion of the two $\\alpha$ elements Ca and Ti with $\\langle$[Ca/Fe]$\\rangle$=+0.42 ($\\sigma$=0.12) and $\\langle$[Ti/Fe]$\\rangle$=+0.24 ($\\sigma$=0.06).  These values are consistent with the results from Kraft et al. (1995) that found $\\langle$[Ca/Fe]$\\rangle$=+0.29 ($\\sigma$=0.07) and $\\langle$[Ti/Fe]$\\rangle$=+0.21 ($\\sigma$=0.12) for a set of 10 other upper RGB stars in this cluster.  The proxy Fe--peak elements Sc and Ni exhibit near solar abundance ratios in all stars with cluster average values of $\\langle$[Sc/Fe]$\\rangle$=+0.03 ($\\sigma$=0.19) and $\\langle$[Ni/Fe]$\\rangle$=+0.09 ($\\sigma$=0.06), which are roughly consistent with Kraft et al. (1995).  These abundances patterns are mirrored in M12 (see Figure \\ref{f2}), but with M10 showing a smaller range of [Cr/Fe] and [Co/Fe] abundances.  The combination of $\\alpha$ enhancement and near solar Fe--peak ratios is consistent with this cluster being primarily polluted by the ejecta of type II SNe (e.g., Woosley \\& Weaver 1995).  For stars near M10's metallicity, La is produced primarily via the s--process in $\\sim$1--3 M$_{\\odot}$ stars and Eu from the r--process in $\\sim$8--10 M$_{\\odot}$ stars (e.g., Busso et al. 1999; Truran et al. 2002).  Our derived La and Eu abundances are consistent with the picture of massive stars producing most of the heavy elements in this cluster with $\\langle$[La/Fe]$\\rangle$=+0.08 ($\\sigma$=0.29) and $\\langle$[Eu/Fe]$\\rangle$=+0.54 ($\\sigma$=0.10).  Comparing the ratio of r-- to s--process elements gives [Eu/La]=+0.41 and implies M10 is slightly more r--process rich than the average globular cluster.  However, this value is within the 1$\\sigma$ range of $\\langle$[Eu/Ba,La]$\\rangle$=+0.23 ($\\sigma$=0.21) found by Gratton et al. (2004) after combining data from the literature on 28 globular clusters.  A larger sample size of M10 stars is likely to decrease the star--to--star scatter observed in our La and Eu sample but will probably not change the result that the cluster is r--process rich."
    },
    "0808/0808.3353_arXiv.txt": {
        "abstract": "{The peculiar hot star \\tc\\ in the open cluster IC\\,2602 is a blue straggler as well as a single-line binary of short period (2.2d).} {Its high-energy properties are not well known, though X-rays can provide useful constraints on the energetic processes at work in binaries as well as in peculiar, single objects. } {We present the analysis of a 50\\,ks exposure taken with the \\xmm\\ observatory. It provides medium as well as high-resolution spectroscopy. } {Our high-resolution spectroscopy analysis reveals a very soft spectrum with multiple temperature components (1--6\\,MK) and an X-ray flux slightly below the `canonical' value ($\\log[L_X(0.1-10.)/L_{BOL}]\\sim-7$). The X-ray lines appear surprisingly narrow and unshifted, reminiscent of those of $\\beta$\\,Cru and $\\tau$\\,Sco. Their relative intensities confirm the anomalous abundances detected in the optical domain (C strongly depleted, N strongly enriched, O slightly depleted). In addition, the X-ray data favor a slight depletion in neon and iron, but they are less conclusive for the magnesium abundance (solar-like?). While no significant changes occur during the \\xmm\\ observation, variability in the X-ray domain is detected on the long-term range. The formation radius of the X-ray emission is loosely constrained to $<$5\\,R$_{\\odot}$, which allows for a range of models (wind-shock, corona, magnetic confinement,...) though not all of them can be reconciled with the softness of the spectrum and the narrowness of the lines.} {} ",
        "introduction": "\\tc\\ (=HD\\,93030) is a luminous star of type B0.2V belonging to the open cluster IC\\,2602, situated at 152\\,pc \\citep{rob99}. The cluster is 30\\,Myr old, and the massive, hot \\tc\\ therefore appears to be a rare example of blue straggler. Its optical spectrum display hints of enhanced nitrogen (a three-fold enrichment with respect to solar) and depleted carbon (at least by an order of magnitude) and oxygen (a factor of about 1/4, see \\citealt{hub08}). In addition, \\tc\\ is also a binary system, with short period (2.2d, \\citealt{llo95}) and small eccentricity ($e$=0.13, \\citealt{hub08}). The companion remains undetected in the visible spectrum (its flux contributes to $<$0.1\\% of the total flux) and should therefore be of a much later spectral type ($M\\sim1$\\,M$_{\\odot}$, \\citealt{hub08}). The binarity and anomalous abundances might suggest that the blue straggler character of \\tc\\ results from a past episode of mass-transfer between these two stars.\\\\  Because of its peculiar properties, the spectrum of \\tc\\ was investigated several times, notably to search for the presence of a magnetic field \\citep{bor79,hub08}. The results are rather inconclusive, but an intriguing period of about 9 minutes was detected in the most recent spectropolarimetric results. \\\\  We decided to further study the star in the high-energy domain. As it was detected by Einstein and Rosat, \\tc\\ is known as the brightest X-ray source of IC\\,2602, but only approximate X-ray properties have up to now been derived. The current generation of X-ray facilities (Chandra, \\xmm) provides a detailed insight on the X-ray emission, especially thanks to their grating instruments. However, these observatories have observed only a handful of B stars at high spectral resolution ($\\tau$\\,Sco, $\\beta$\\,Cru, $\\epsilon$\\,Ori, Spica) and no blue straggler. The analysis of \\tc\\ thus nicely fills a gap, and could help better understand the high-energy characteristics of hot stars.\\\\  The paper is organized as follows: the data and reduction processes are presented in Sect. 2, the general properties of \\tc\\ in the X-ray domain are examined in Sect. 3.1, the detailed characteristics of the X-ray lines are derived in Sect. 3.2, and we finally conclude in Sect. 4. ",
        "conclusions": "\\xmm\\ observations revealed the atypical character of \\tc. Overall, its X-ray emission appears very soft as well as rather weak. Almost all the flux is found below 1\\,keV; indeed, spectral fits indicate a dominant temperature of about 0.2\\,keV. The total unabsorbed luminosity in the 0.1--10\\,keV range amounts to 1.0$\\times10^{31}$~erg\\,s$^{-1}$, which yields a $L_X/L_{BOL}$ ratio slightly lower than the `canonical' value for OB stars. Though no significant variability is detected during the 50\\,ks \\xmm\\ observation, an increase (resp. decrease) of the flux is clearly observed when comparing with previous Einstein (resp. ROSAT) observations. \\\\  The high-resolution spectrum of \\tc, revealed by the RGS, does not show any significant continuum emission but is solely composed of lines of H- and He-like ions of N, O, and Ne as well as some iron lines. To the resolution limits of the RGS, these lines appear narrow and unshifted: the observed width of these lines is mainly instrumental and the intrinsic width is limited to $<350$\\,\\kms. It must be noted that the nitrogen lines are anomalously strong: the spectral fits reveal a large enrichment in nitrogen (abundance 3 times solar), together with a strong deficit in carbon, which both agree well with the values found in the optical spectrum \\citep{hub08}. Neon, oxygen and iron also appear slightly depleted in the X-ray spectrum. For the {\\it fir} triplets, the forbidden component is fully suppressed while the intercombination lines are slightly stronger than the resonance component. The {\\it f/i} ratios suggests a formation radius rather close to the star, below 5 stellar radii.\\\\  With its soft X-ray spectrum and its narrow X-ray lines, \\tc\\ clearly appears different from O-type stars. Comparing \\tc\\ to other B-type objects with high-resolution spectra, we find that the temperature distribution is quite typical of `normal' B stars: the DEM is similar to that of $\\epsilon$\\,Ori and only slightly broader than for $\\beta$\\,Cru and Spica \\citep{zhe07}. Similar temperatures, fluxes and overall line properties were also found in a detailed analysis of $\\beta$\\,Cru \\citep{coh08}, though in the latter case, more precise constraints could be found on the formation radius. On the other hand, \\tc\\ seems very different from $\\tau$ Sco as far as softness, abundances, and flux level are concerned, but both objects present unshifted, narrow X-ray lines \\citep{mew,coh03}. Such lines are at odds with the wind-shock model, though it is still unclear how this model could apply in the case of low mass-loss rates, as those of B-type stars. Actually, narrow lines are often attributed to magnetic confinement. However, this process is still uncertain for \\tc, as magnetic field searches were inconclusive up to now \\citep{hub08}. As for $\\beta$\\,Cru, additional observations, especially polarimetric ones, are requested before \\tc\\ can be fully understood and the origin of its X-ray emission pinpointed."
    },
    "0808/0808.3769_arXiv.txt": {
        "abstract": "The uniformity of the helium-to-hydrogen abundance ratio in X-ray emitting intracluster medium (ICM) is one of the commonly adopted assumptions in X-ray analyses of galaxy clusters and cosmological constraints derived from these measurements.  In this work, we investigate the effect of He sedimentation on X-ray measurements of galaxy clusters in order to assess this assumption and associated systematic uncertainties. By solving a set of flow equations for a H-He plasma, we show that the helium-to-hydrogen mass ratio is significantly enhanced in the inner regions of clusters.  The effect of He sedimentation, if not accounted for, introduces systematic biases in observable properties of clusters derived using X-ray observations. We show that these biases also introduce an apparent evolution in the observed gas mass fractions of X-ray luminous, dynamically relaxed clusters and hence biases in observational constraints on the dark energy equation of state parameter, $w$, derived from the cluster distance-redshift relation.  The Hubble parameter derived from the combination of X-ray and Sunyaev-Zel'dovich effect (SZE) measurements is affected by the He sedimentation process as well. Future measurements aiming to constrain $w$ or $H_0$ to better than 10\\% may need to take into account the effect of He sedimentation.  We propose that the evolution of gas mass fraction in the inner regions of clusters should provide unique observational diagnostics of the He sedimentation process. ",
        "introduction": "Clusters of galaxies are powerful cosmological probes and have the potential to constrain properties of dark energy and dark matter.  Recent development in X-ray observations of galaxy clusters have produced a large statistical sample of clusters and start to deliver powerful cosmological constraints \\citep{Allen2004Constraints_on_,Allen2008Improved,Mantz2007New_constraints, Vikhlinin2008} that are complimentary to and competitive with other techniques (e.g., supernova, baryon acoustic oscillation, and weak lensing). This has motivated construction of the next-generation of X-ray satellite missions (e.g., \\emph{eROSITA}) to push the precision cosmological measurements based on large X-ray cluster surveys.  However, in the era of precision cosmology, the use of clusters as sensitive cosmological probes require solid understanding of cluster gas physics, testing of simplifying assumptions, and assessing associated systematic uncertainties.  One of the commonly adopted assumptions in X-ray cluster analyses include the uniformity of the helium-to-hydrogen abundance ratio with nearly primordial composition in X-ray emitting intracluster medium (ICM).  At present, there is no observational test of this assumption, since both H and He in the ICM are fully ionized, which makes it difficult to measure their abundances using traditional spectroscopic techniques.  Theoretically, on the other hand, it has long been suggested that heavier He nuclei slowly settle in the potential well of galaxy clusters and cause a concentration of He toward their center \\citep{Abramopoulos1981On_the_equilibr,Gilfanov1984Intracluster,Qin2000BARYON-DISTRIBU,Chuzhoy2003Gravitational,Chuzhoy2004Element,Ettori2006Effects}. In the era of precision cosmology, this could be a source of significant systematic uncertainties in X-ray measurements of galaxy clusters and cosmological parameter derived from these measurements \\citep{markevitch2007Helium_abundance}.  Thus, the primary goal of the present work is to assess the validity of this assumption and associated systematic uncertainties in X-ray measurements of key cluster properties as well as cosmological parameters derived from these observations.  In this work, we investigate the effects of He sedimentation on X-ray measurements of galaxy clusters by solving a set of diffusion equations for a H-He plasma in the ICM.  By taking into account observed temperature profiles obtained by recent X-ray observations \\citep{Vikhlinin2006Chandra_Sample,2007Pratt, 2008Leccardi,2008George}, we show that the observed temperature drop in the cluster outskirts lead to a significant suppression of He sedimentation, compared to the results based on the isothermal cluster model \\citep{Chuzhoy2004Element}.  Our analysis indicates that the He sedimentation has negligible effect on X-ray measurements in the outer regions of clusters (e.g., $r_{500}$), and it does not affect cluster mass measurements obtained at the sufficiently large cluster radius.  The effect of He sedimentation, on the other hand, introduces increasingly larger biases in X-ray measurements in the inner regions and could affect cosmological constraints, including the dark energy equation of state parameter $w$ derived from distance-redshift relation as well as $H_0$ derived from the combination of X-ray and Sunyaev-Zel'dovich effect (SZE).  The paper is organized as follows.  In \\S~\\ref{sec:xray}, we describe the dependence of X-ray clusters measurements on He abundance in the ICM. The physics of He sedimentation and cluster models are discussed in \\S~\\ref{sec:diffusion}.  In \\S~\\ref{sec:results}, we present results of our He sedimentation calculations and investigate their effects on cluster properties and cosmological constraints derived from X-ray cluster observations.  Main conclusions are summarized in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this work, we investigate effects of He sedimentation on X-ray measurements of galaxy clusters and their implication for cosmological constraints derived from these observations.  By solving a set of flow equations for a H-He plasma and using observationally motivated cluster models, we show that the efficiency of He sedimentation is significantly suppressed in the cluster outskirts due to the observed temperature drop, while it is dramatically enhanced in the cluster core regions. Our sedimentation model based on the observed temperature profile suggest that the effect of helium sedimentation is negligible at $r_{500}$, and it does not affect cluster mass measurements obtained at the sufficiently large cluster radius. However, the effect of sedimentation increases toward the inner regions of clusters and introduces biases in X-ray measurements of galaxy clusters. For example, at $r_{2500}$, biases in X-ray measurements of gas mass, total mass, and gas mass fractions, are at the level of $5-10\\%$. The effect of He sedimentation could also introduce biases in the estimate of the Hubble parameter derived from the combination of X-ray and SZE measurements, which could explain the observed offset in the X-ray+SZE derived $H_0$ and independent measurement from the Hubble Key project. We emphasize, however, that the magnitude of these biases depends sensitively on the cluster age, temperature, and magnetic and/or turbulent suppression in the ICM.  We show that the process of He sedimentation introduces the apparent evolution in the observed gas mass fractions of X-ray luminous, dynamically relaxed clusters.  The effect of He sedimentation could lead to biases in observational constraints of dark energy equation of state $w$ at a level of  $\\lesssim$10\\%. These biases tend to make the value of $w$ more negative. Current measurements based on $f\\gas$ evolution \\citep{Allen2004Constraints_on_,Allen2008Improved} should not be significantly affected by these biases.  However, future measurements aiming to constrain $w$ to better than 10\\% may need to take into account the effect of He sedimentation.  For cosmological measurements, one way to minimize these biases is to extend the X-ray measurements to a radius well beyond $r_{2500}$.  At the same time, the evolution of cluster gas mass fraction in the inner regions of clusters should provide unique observational diagnostics of the He sedimentation process in clusters."
    },
    "0808/0808.1841_arXiv.txt": {
        "abstract": "The morphology of the outer rings of early-type spiral galaxies is compared to integrations of massless collisionless particles initially in nearly circular orbits.  Particles are perturbed by a quadrupolar gravitational potential corresponding to a growing and secularly evolving bar. We find that outer rings with R1R2 morphology and pseudorings are exhibited by the simulations even though they lack gaseous dissipation. Simulations with stronger bars form pseudorings earlier and more quickly than those with weaker bars. We find that the R1 ring, perpendicular to the bar, is fragile and dissolves after a few bar rotation periods if the bar pattern speed increases by more than $\\sim 8\\%$, bar strength increases (by $\\gtrsim 140\\%$) after bar growth, or the bar is too strong ($Q_T>0.3$). If the bar slows down after formation, pseudoring morphology persists and the R2 ring perpendicular to the bar is populated due to resonance capture. The R2 ring remains misaligned with the bar and increases in ellipticity as the bar slows down. The R2 ring becomes scalloped and does not resemble any ringed galaxies if the bar slows down more than 3.5\\% suggesting that bars decrease in strength before they slow down this much. We compare the morphology of our simulations to B-band images of 9 ringed galaxies from the Ohio State University Bright Spiral Galaxy Survey, and we find a reasonable match in morphologies to R1R2' pseudorings seen within a few bar rotation periods of bar formation. Some of the features previously interpreted in terms of dissipative models may be due to transient structure associated with recent bar growth and evolution. ",
        "introduction": "Rings in barred galaxies can exist interior to the bar, encircling the bar or exterior to the bar. For a review on classification and properties of ringed galaxies see \\citet{buta96}. The outer rings of barred galaxies are classified as R1 or R2 depending upon whether the ring is oriented with major axis perpendicular to the bar (R1) or parallel to it (R2) (e.g., \\citealt{romero06}). If the ring is broken, partial or is a tightly wrapped spiral it is called a pseudoring and denoted R1' or R2'.  Some galaxies contain both types of rings and are denoted R1R2' or R1R2. R1' and R2' morphologies were predicted as morphological patterns that would be expected near the outer Lindblad resonance (OLR) with the bar \\citep{schwarz81,schwarz84}. Rings are often the site of active star formation and so are prominent in blue visible band images, H$\\alpha$ narrow band images, and HI emission \\citep{buta96}.  Orbital resonances, denoted Lindblad Resonances, occur at locations in the disk where \\begin{equation} \\Omega_b = \\Omega \\pm \\kappa/m \\end{equation} where $\\Omega_b$ is the angular rotation rate of the bar pattern and $m$ is an integer. Here $\\Omega(r)$ is the angular rotation rate of a star in a circular orbit at radius $r$ and $\\kappa(r)$ is the epicylic frequency.  The $m=2$ OLR is that with $\\Omega_b = \\Omega + \\kappa/2$. Orbits of stars are often classified in terms of nearby periodic orbits that are closed in the frame rotating with the bar. Near resonances orbits become more elongated and have higher epicyclic amplitudes. Exterior to the OLR periodic orbits parallel to the bar are present whereas interior to the OLR both perpendicular and parallel periodic orbits are present. For a steady pattern, closed orbits interior to the OLR are expected to be aligned with major axis perpendicular to the bar whereas those exterior to the OLR are aligned parallel to it (e.g., \\citealt{cont89,kalnajs91}).  A common assumption is that rings form because gas accumulates at resonances.  This follows as gas clouds cannot follow self-intersecting orbits without colliding.  Because of dissipation in the gas, the bar can exert a net torque on the gas leading to a transfer of angular momentum.  The torque is expected to change sign at resonances so gas can move away from them or accumulate at them. The CR region is expected to be depopulated leading to gas concentrations at the OLR and ILR resonances. Gaseous rings form when gas collects into the largest periodic orbit near a resonance that does not cross another periodic orbit \\citep{schwarz84}.  \\citet{schwarz81,schwarz84} first demonstrated the efficiency of this process. Other papers have confirmed and extended this work (e.g., \\citealt{combes85,byrd94,salo99,rau00,rau04}). Because dissipation is thought to be important, spiral like features and ovals that are not perfectly aligned with the bar, similar to those observed, are predicted. In some cases galaxy morphology and kinematics have not been successfully modeled with a single steady state bar component. Improvements in the models have been made with the addition of an additional exterior oval or spiral component (e.g., \\citealt{hunter88,lindblad96}).  Previous work accounting for ring galaxy morphology has primarily simulated the gas dynamics using sticky particle simulations that incorporate dissipative or inelastic collisions. \\citet{rau00} ran N-body stellar simulations  coupled with sticky gas particles. These simulations have self-consistent bars so that the orbits of the stars in the bars are consistent with the bar's gravitational potential. The disadvantage of using N-body simulations is that the properties of the bar such as its pattern speed and strength cannot be set. They can only be changed indirectly by varying the initial conditions of the simulations. An alternative approach is to set the bar perturbation strength, shape and pattern speed and search for likely bar parameters consistent with the properties of observed galaxies (e.g., \\citealt{salo99,rau04,rau08}).  Previous work has explored the affect of bar strength and pattern speed on ring morphology (e.g., \\citealt{salo99,rau04,rau08}) and length of time since the bar grew (e.g., \\citealt{rau00,ann00}). Here we explore the role of bar evolution on ring galaxy morphology.  By bar evolution we mean changes in bar pattern speed and strength during and after bar growth. N-body simulations lacking live halos predict long lived bars with nearly constant pattern speeds (e.g., \\citealt{voglis07}). However angular momentum transfer between a bar and the gas disk either interior or exterior to the bar or between a bar and a live halo can cause the pattern speed to vary (e.g., \\citealt{debattista98,bournaud02,das03,ath03,sellwood06,martinez06}). Thus constraints on the secular evolution of bars could tell us about the coupling between bars, gas and dark halos.  Gas and stars exterior to a bar are sufficiently distant and moving sufficiently slowly compared to the bar that they are unlikely to cause strong perturbations on the orbits of stars in the bar. Because a calculation of the gravitational potential involves a convolution with an inverse square law function, high order Fourier components are felt only extremely weakly exterior to the bar. The dominant potential term exterior to the bar is the quadrupolar term which decreases with radius to the third power, $\\Phi \\propto r^{-3}$.  Here we explore the role of a changing quadrupolar potential field on the morphology of stars exterior to a bar. In this work we focus on collisionless stellar orbits and leave investigating the study of dissipative effects for future study.  In Section 2 we describe our simulations and present the results obtained by varying the parameters.  In Section 3 we compare the results of our simulations with 9 galaxies from the Ohio State University Bright Spiral Galaxy Survey (\\citet{eskridge02}, hereafter OSUBSGS). Finally in Section 4 we summarize and discuss our results. ",
        "conclusions": "We have presented integrations of collisionless  massless particles perturbed by growing and secularly evolving bar perturbations. We find that collisionless simulations can exhibit double ringed R1 and R2 outer ring morphology with rings both perpendicular (R1) and parallel (R2) to the bar. In the last period of bar growth, strong open spiral structure is exhibited resembling an R1' pseudoring.   For 2-3 periods following bar growth R1 and R2 rings are seen with the R2 ring changing in orientation and azimuthal density contrast. Thus R1R2' pseudoring morphology is displayed within a few bar periods following bar growth. Our simulations start with particles in nearly circular orbits with velocity dispersions equivalent to 7 km/s for a 200 km/s rotation curve. This suggests sticky particle simulations have been successful in exhibiting R1R2 ring morphology because the velocity dispersion of orbits is damped and so particles are in initially nearly circular orbits.  In our collisionless simulations we find that the outer rings with major axis perpendicular (R1) to the bar are fragile. If the bar pattern speed increases more than 8\\% after bar growth, or if the bar strength is higher than or increases past $|\\epsilon| \\gtrsim 0.16$ or $Q_T \\gtrsim 0.32$ the R1 outer ring will dissolve after $\\sim 20$ twenty bar periods. The simulations are then nearly mirror symmetric and do not display asymmetries typical of pseudorings.  Stronger bars can form R1' pseudorings earlier. However if the bar strength $|\\epsilon| \\gtrsim 0.16$ or $Q_T \\gtrsim 0.32$ the R1 ring will dissolve after $\\sim 20$ bar rotation periods. If the bar strength increases to this value subsequent to formation, the R1 ring also dissolves.  We find that a decrease in the bar pattern speed after bar growth causes particles to be captured in orbits parallel to the bar which are increased in epicyclic amplitude as the bar slows down. Strong R1 and elongated R2 rings persist in these simulations. Misalignments between the R2 ring and the bar also persist so the galaxy can exhibit R1R2' pseudoring morphology for a longer period of time.  If the bar pattern speed slows down more than $\\sim 3.5\\%$ the R2 ring develops a scallop above and below the bar. As these are not observed in galaxies, bars probably do not slow down more than $\\sim 3.5\\%$ without also varying in strength.  \\citet{sandage94} find that early type barred galaxies often have semi-detached outer rings (e.g, NGC 1543, \\citealt{buta96} and NGC 4457). These galaxies may contain bars that have increased in pattern speed or were once strong and so destroyed their R1 ring. If the bar weakens the R1 and R2 rings can be left behind as two nearly circular rings, similar to those observed in the unusual double outer ringed galaxy NGC 2273.  We find that the morphology of our simulations resembles that of R1' ringed galaxies if the simulation time is chosen during or just after bar formation.  We find we can match pseudoring morphology with simulations that have bar strengths estimated from the bar shapes. Stronger and longer spiral arms are seen later in the simulation and in more strongly barred systems. The constraint on simulation timescale suggests that R1' ring morphology is a signpost of recent bar formation. We note that sticky particle and SPH simulations  exhibit R1 pseudoring morphology a few bar rotation periods longer than ours suggesting that the dissipationless simulations explored here underestimate the longevity of these features.  We find that galaxies with R1R2' morphology are well matched by simulations a few bar rotation periods following bar growth. As R1 rings are fragile, we infer that these galaxies have had stable bars that have not experienced large changes in either pattern speed or strength.  The exploration of parameter space in the collisionless dissipationless limit done here can be used by future work to differentiate between phenomena that would be exhibited by collisionless models and that that is a result of dissipation. A better understanding of the role of dissipation in affecting outer ring morphology should allow observationally based constraints on the secular evolution of bars.  Only 10-20\\% of early type galaxies exhibit outer rings with pseudorings being more prevalent in later type galaxies \\citep{buta96}. Not all but most galaxies classified with outer rings are barred suggesting that only 15-40\\% of barred galaxies exhibit outer rings. Here we have found that R1' and R1R2' galaxies are likely to represent different times since bar formation with R1' galaxies representing an earlier timescale during or just after bar formation and R1R2 morphology representing galaxies with stable bars a few bar rotation periods following bar formation.  Galaxies in these two transient categories probably comprise a significant fraction of all outer ring galaxies. This suggests that most outer ring galaxies represent morphology that is only present for a few bar rotation periods. It is interesting to ask what timescales these morphologies correspond to. Bar rotation periods for the ringed galaxies in our sample range from $\\sim 100-200$ Myr (see \\ref{tab:tab3}). The R1' classification, may only last a few bar rotations or 1/2 Gyr and the R1R2' classification only $\\sim$ 1 Gyr. Both of these timescales are short compared to the lifetime of a galaxy. Ringed galaxies lacking R1 rings may be longer lived but may provide evidence for bar evolution. It is likely that only a low fraction of barred galaxies might be considered systems that are not evolving secularly or have not formed in the last Gyr."
    },
    "0808/0808.1593_arXiv.txt": {
        "abstract": "The nearby radio galaxy Centaurus A is poorly studied at high frequencies with conventional radio telescopes because of its very large angular size, but is one of a very few extragalactic objects to be detected and resolved by the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}). We have used the five-year {\\it WMAP} data for Cen~A to constrain the high-frequency radio spectra of the 10-degree giant lobes and to search for spectral changes as a function of position along the lobes. We show that the high-frequency radio spectra of the northern and southern giant lobes are significantly different: the spectrum of the southern lobe steepens monotonically (and is steeper further from the active nucleus) whereas the spectrum of the northern lobe remains consistent with a power law. The inferred differences in the northern and southern giant lobes may be the result of real differences in their high-energy particle acceleration histories, perhaps due to the influence of the northern middle lobe, an intermediate-scale feature which has no detectable southern counterpart. In light of these results, we discuss the prospects for {\\it Fermi Gamma-ray Space Telescope} detections of inverse-Compton emission from the giant lobes and the lobes' possible role in the production of the ultra-high energy cosmic rays (UHECR) detected by the Pierre Auger Observatory. We show that the possibility of a {\\it Fermi} detection depends sensitively on the physical conditions in the giant lobes, with the northern lobe more likely to be detected, and that any emission observed by {\\it Fermi} is likely to be dominated by photons at the soft end of the {\\it Fermi} energy band. On the other hand we argue that the estimated conditions in the giant lobes imply that UHECRs can be accelerated there, with a potentially detectable $\\gamma$-ray signature at TeV energies. ",
        "introduction": "\\label{intro}  Centaurus A is the closest radio galaxy to us (we adopt $D = 3.7$ Mpc, the average of 5 distance indicators in Ferrarese \\etal\\ 2007). Its proximity makes it one of the brightest extragalactic radio sources in the sky at low frequencies (only exceeded by Cygnus A: Baars \\etal\\ 1977), but also means that the outer `giant' double lobes (throughout the paper we use the nomenclature adopted by Alvarez et al 2000) subtend an angle of $\\sim 10^\\circ$ on the sky, although their total physical size ($\\sim 600$ kpc in projection) is not unusually large for an radio galaxy of Cen A's luminosity. The large angular size of the lobes has prevented the type of {\\it spatially resolved}, multifrequency study of their spectral structure that is commonplace for more distant radio galaxies (e.g. Alexander \\& Leahy 1987). While detailed low-frequency maps of the giant lobes have been available for many years (e.g. Cooper, Price \\& Cole 1965), published radio data from ground-based observations only exist up to 5 GHz (Junkes \\etal\\ 1993) and the spectral study of Alvarez \\etal\\ (2000), involving (at many frequencies) painstaking graphical integration of contour maps, was only able to determine overall spectra for the giant lobes, finding no evidence for deviation from a single power law between 408 MHz and 5 GHz. The low-frequency two-point spectral index maps of Combi \\& Romero (1997), however, do show some evidence for position-dependent spectral steepening, particularly towards the end of the southern giant lobe.  Cen~A is widely believed to be a restarting radio galaxy, in the sense that the inner lobes are the result of the current nuclear activity, while the giant outer lobes are the result of a previous outburst (Morganti \\etal\\ 1999). This picture is supported by the observation of hot thermal X-ray emission, apparently the result of strong shocks, surrounding both inner lobes (Kraft \\etal\\ 2003, 2007; Croston \\etal , in prep.) which implies that they are propagating supersonically into the intergalactic medium (IGM) of Cen~A and are disconnected from the giant lobes. However, the nature of the intermediate-scale northern middle lobe (NML: Morganti \\etal\\ 1999) is not completely clear in this model. In standard spectral ageing models (e.g. Jaffe \\& Perola 1973), we might expect to see spectral steepening at high frequencies in the giant lobes; a measurement of spectral ageing gives a model-dependent constraint on the time since the last injection of high-energy electrons into these lobes. Such a constraint on spectral age could be compared with other estimates of the dynamical age of the radio source, and would therefore be of considerable interest, but the work of Alvarez \\etal\\ (2000) only sets upper limits on this quantity, since they do not see any deviation from a power-law spectrum.  \\begin{figure*} \\epsfxsize 17.5cm \\epsfbox{montage.eps} \\caption{Large-scale structure of Centaurus A with a resolution of $\\sim 0.83^\\circ$. Contours are at $1,2,4\\dots$ times the base level specified for each map, and take no account of background, except for the 408-MHz map, from which a constant background of 43 Jy beam$^{-1}$ has been subtracted. Top row, from left to right: 408 MHz (3 Jy beam$^{-1}$), 1.4 GHz (1 Jy beam$^{-1}$), 5 GHz (0.25 Jy beam$^{-1}$), 20 GHz (1.0 Jy beam$^{-1}$). Bottom row, from left to right: 30 GHz (0.5 Jy beam$^{-1}$), 40 GHz (1.0 Jy beam$^{-1}$), 60 GHz (1.0 Jy beam$^{-1}$), 90 GHz (2.0 Jy beam$^{-1}$). Circles in the bottom right-hand corner of each image indicate the beam size (diameter shows FWHM). The 20-GHz data have different noise characteristics from the other {\\it WMAP} images because they have not been convolved with a Gaussian; instrumental noise on scales smaller than the effective beam is therefore visible. See the text (Section \\ref{wmap-pro}) for discussion of the convolution, effective resolution and beam area of these images.} \\label{convolved-data} \\end{figure*}  The giant lobes of Cen~A are also interesting because they are predicted to be strong sources of inverse-Compton emission as the relativistic electrons in the lobes scatter cosmic microwave background (CMB) photons to high energies; a detection of inverse-Compton emission from Cen~A would constrain the magnetic field strength in the lobes of Fanaroff \\& Riley (1974) class I (hereafter FRI) radio sources in general: we have little information on the magnetic field strengths in these low-power radio galaxies at present. However, X-ray emission from this process would be distributed on similar scales to the giant lobes, making it hard to detect. Cooke, Lawrence \\& Perola (1978) claimed an early detection of the giant lobes using {\\it Ariel V}, but Marshall \\& Clark (1981) argued that this was the result of point source contamination, placing a much lower upper limit on the flux from {\\it SAS 3} observations. At soft X-ray energies (e.g. Arp 1994) the situation is seriously confused by the presence of known thermal X-ray emission from the interstellar medium of the host galaxy, which more recently has been extensively studied with {\\it Chandra} and {\\it XMM-Newton} (e.g.\\ Kraft \\etal\\ 2003), and is also hard to study because the X-ray emission fills the field of view of modern soft X-ray imaging instruments such as {\\it ROSAT} (Arp 1994), {\\it XMM}, and {\\it ASCA} (Isobe \\etal\\ 2001), presenting almost insuperable problems of background modelling and subtraction. However, the spectrum of the inverse-Compton emission should be hard up to high energies (exactly how high depends on the model adopted for the electron energy spectrum, as we will discuss below). There are thus also interesting constraints from observations at MeV to GeV energies made with the {\\it Compton Gamma-Ray Observatory} (e.g. Steinle \\etal\\ 1998; Sreekumar \\etal\\ 1999), which do not detect the giant lobes but again set upper limits on their high-energy flux densities. More sensitive hard X-ray/$\\gamma$-ray observations exist with wide-field instruments like {\\it INTEGRAL} (e.g. Rothschild \\etal\\ 2006) and the {\\it Swift} Burst Alert Telescope (e.g. Markwardt \\etal\\ 2005) but as these are coded-aperture instruments they have limited sensitivity to extended emission (see e.g. Renaud \\etal\\ 2007). At present, therefore, there is no unambiguous detection of X-ray or $\\gamma$-ray inverse-Compton emission from the giant lobes of Cen~A. One of us has shown (Cheung 2007) that {\\it Fermi}\\footnote{Formerly known as the {\\it Gamma-ray Large Area Space Telescope}, {\\it GLAST}.} may have the sensitivity to detect inverse-Compton emission off the CMB from the lobes of Cen~A at energies from $\\sim$100 MeV to 10 GeV. However, the details of this depend on modelling of the electron energy spectrum at high energies, which in turn depends on high-frequency radio data.  Finally, Cen~A's giant lobes are possible sources of ultra-high energy cosmic rays (UHECRs). Cen~A's proximity means that all aspects of the active galaxy -- central AGN, inner jets and lobes, and giant lobes -- have long been considered as possible UHECR accelerators (see e.g. Cavallo 1978; Romero \\etal\\ 1996, and, more recently, Gureev \\& Troitsky 2008 and references therein). Interest in Cen~A has been spurred by the remarkable discovery that 2 of the 27 UHECR events detected so far by the Pierre Auger Observatory (hereafter `PAO'; Abraham \\etal\\ 2007) appear to be arriving from the direction of the centre of Cen~A, while at least 2 additional events may be associated with it (e.g. Gorbunov \\etal\\ 2008a; Wibig \\& Wolfendale 2007; Fargion 2008) due to the large angular extent of the giant radio lobes (Gorbunov \\etal\\ 2008b; Moskalenko \\etal\\ 2008). Most scenarios discussed in the literature to date assume that UHECRs are produced near the supermassive black hole (SMBH) or in the inner jets (e.g. Cuoco \\& Hannestad 2008; Kachelriess, Ostapchenko \\& Tomas 2008), but an explanation in terms of the giant lobes has the advantage that it can easily account for the PAO events seen on larger scales. In order to investigate this quantitatively we need information about the magnetic field strengths and the {\\it leptonic} particle acceleration in these lobes, which can be provided by a combination of high-frequency radio observations and inverse-Compton constraints or measurements.  The {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}) has observed the whole sky at frequencies around 20, 30, 40, 60 and 90 GHz (known as K, Ka, Q, V and W bands respectively) with the aim of measuring structure in the CMB (e.g.\\ Hinshaw \\etal\\ 2008). The currently available {\\it WMAP} data represent 5 years of observations. Cen~A is clearly detected, and spatially resolved, in the {\\it WMAP} observations at all frequencies (e.g. Page \\etal\\ 2007) and Israel \\etal\\ (2008) have recently presented {\\it WMAP}-derived measurements of the flux density of the whole source, showing that there is clear steepening in the integrated spectrum at high frequencies. Thus the data are available to carry out a study of the variation of the radio spectrum as a function of position, to fit spectral ageing models to the large-scale lobes and investigate whether we can learn anything about the source dynamics, and to make predictions of the expected inverse-Compton emission from the giant lobes.  In this paper we present the results of such a study. We first combine the 5-year {\\it WMAP} data on Cen~A with single-dish radio images at lower frequencies to make spatially resolved measurements of the radio spectra from 408 MHz to 90 GHz. We then discuss the implications of the high-frequency detections for the dynamics of Cen~A, for possible inverse-Compton detections of the giant lobes at high energies, and for acceleration of UHECRs and their possible $\\gamma$-ray emission signatures. ",
        "conclusions": ""
    },
    "0808/0808.3305_arXiv.txt": {
        "abstract": "In this paper, we review some of the properties of dense molecular cloud cores. The results presented here rely on three-dimensional numerical simulations of isothermal, magnetized, turbulent, and self-gravitating molecular clouds (MCs) in which dense core form as a consequence of the gravo-turbulent fragmentation of the clouds. In particular we discuss issues related to the mass spectrum of the cores, their lifetimes and their virial balance. ",
        "introduction": "We performed 3D numerical simulations of magnetized, self-gravitating, and turbulent isothermal MCs using the TVD code (Kim et al. 1999) on grids with $256^{3}$ and $512^{3}$ cells (V\\'{a}zquez-Semadeni et al. 2005; Dib et al. 2007a; Dib et al. 2008a). The basic features of these simulations are: Turbulence is driven until it is fully developed (at least for 2 crossing timescales) before gravity is turned on. The Poisson equation is solved to account for the self-gravity of the gas using a standard Fourier algorithm. Turbulence is constantly driven in the simulation box following the algorithm of Stone et al. (1998). The kinetic energy input rate is adjusted such as to maintain a constant rms sonic Mach number $M_{s}=10$. Kinetic energy is injected at large scales, in the wave number range $k=1-2$. Periodic boundary conditions are used in the three directions. In physical units, the simulations have a linear size of 4 pc, an average number density of $\\bar{n}=500$ cm$^{-3}$, a temperature of 11.4 K, a sound speed of $0.2$ km s$^{-1}$, and an initial {\\it rms} velocity of 2 km s$^{-1}$. The Jeans number of the box is $J_{box}=4$ (number of Jeans masses in the box is $M_{box}/M_{J,box}=J_{box}^{3}=64$, where $M_{box}=1887$ M$_{\\odot}$). The simulations vary by the strength of the magnetic field in the box with $B_{0}= 0, 4.6, 14.5$, and 45.8 $\\mu$G for the non-magnetized, the strongly supercritical, the mildly supercritical, and the subcritical cloud models, respectively. Correspondingly, the plasma beta and mass-to-magnetic flux (normalized for the critical value for collapse $M/\\phi =(4 \\pi^{2} G)^{-1/2}$; Nakano \\& Nakamura 1978) values of the box are $\\beta_{p,box}=\\infty, 1, 0.1$, and $0.01$, and $\\mu_{box}=\\infty, 8.8, 2.8$, and $0.9$, respectively. Cores are identified using a clump-finding algorithm that is based on a density threshold criterion and a friend-of-friend approach as described in Dib et al. (2007a). We restrict our selection of cores to epochs where the Truelove criterion (Truelove et al. 1997) is not violated in any of them. Thus, the derived properties of our numerical cores can be best compared to those of starless prestellar cores. ",
        "conclusions": ""
    },
    "0808/0808.3902_arXiv.txt": {
        "abstract": "We discuss the bounds on the mass of Dark Matter (DM) particles, coming from the analysis of DM phase-space distribution in dwarf spheroidal galaxies (dSphs). After reviewing the existing approaches, we choose two methods to derive such a bound. The first one depends on the information about the current phase space distribution of DM particles only, while the second one uses both the initial and final distributions. We discuss the recent data on dSphs as well as astronomical uncertainties in relevant parameters. As an application, we present lower bounds on the mass of DM particles, coming from various dSphs, using both methods. The model-independent bound holds for any type of fermionic DM. Stronger, model-dependent bounds are quoted for several DM models (thermal relics, non-resonantly and resonantly produced sterile neutrinos, etc.). The latter bounds rely on the assumption that baryonic feedback cannot significantly increase the maximum of a distribution function of DM particles. For the scenario in which all the DM is made of sterile neutrinos produced via non-resonant mixing with the active neutrinos (NRP) this gives $m_{\\dw}>1.7$ keV.  Combining these results in their most conservative form with the X-ray bounds of DM decay lines, we conclude that the NRP scenario remains allowed in a very narrow parameter window only. This conclusion is independent of the results of the Lyman-alpha analysis. The DM model in which sterile neutrinos are resonantly produced in the presence of lepton asymmetry remains viable. Within the minimal neutrino extension of the Standard Model (the $\\nu$MSM), both mass and the mixing angle of the DM sterile neutrino are bounded from above and below, which suggests the possibility for its experimental search. ",
        "introduction": "\\label{sec:introduction}  The nature of Dark Matter is one of the most intriguing questions of particle astrophysics. Its resolution would have a profound impact on the development of particle physics beyond the Standard Model.  Although the possibility of having massive compact halo objects (MACHOs) as a dominant form of DM is still under debate (see recent discussion in~\\cite{Calchi:07} and references therein), it is widely believed that Dark Matter is composed of non-baryonic particles.  However, the Standard Model of elementary particles does not contain a viable Dark Matter particle candidate -- a massive, neutral and long-lived particle. Active neutrinos, which are both neutral and stable, form structures in a top-down fashion~\\cite{Zeldovich:70,Bisnovatyi:80,Bond:80,Doroshkevich:81,Bond:83}, and thus cannot produce the observed quantity of early-type galaxies~\\cite[see e.g.][]{White:83,Peebles:84a}.  Therefore, the DM particle hypothesis implies the extension of the Standard Model (SM).  The DM particle candidates may have very different masses (for reviews of DM candidates see e.g.~\\cite{Bergstrom:00,Bertone:05,Carr:06,Taoso:07}): massive gravitons with the mass $\\sim 10^{-19}\\ev$~\\cite{Dubovsky:04}, axions with the mass $\\sim 10^{-6}\\ev$~\\cite{Holman:83}, sterile neutrinos having mass in the keV range~\\cite{Dodelson:93}, sypersymmetric (SUSY) particles (gravitinos~\\cite{Pagels:82}, neutralinos~\\cite{Haber:85}, axinos~\\cite{Covi:99} with their masses ranging from eV to hundreds GeV, supersymmetric Q-balls~\\cite{Kusenko:97b}, WIMPZILLAs with the mass $\\sim 10^{13}\\gev$~\\cite{Kuzmin:98,Chung:99}, and many others). Thus, the mass of DM particles becomes an important characteristic which may help to distinguish between various DM candidates and, more importantly, may help to differentiate among different models beyond the SM.   It was suggested in~\\cite{Tremaine:79} that quite a robust and model-independent \\emph{lower bound} on the mass of DM particles can be obtained by considering phase space density evolution of compact astrophysical objects, most notably dwarf spheroidal satellites (dSphs) of the Milky Ways. The idea was developed further in a number of works (see e.g.~\\cite{Madsen:84,Madsen:91,Madsen:90,Dalcanton:00,Hogan:00,Madsen:00}).%  Another way to distinguish between various DM models, and particularly to put a bound on the DM mass, is the analysis of the Lyman-$\\alpha$ (Ly-$\\alpha$) forest data~\\cite{Hui:97,Gnedin:01,Weinberg:03}.\\footnote{Absorption feature by neutral hydrogen at $\\lambda = 1216$~\\AA~at different redshifts in the spectra of distant quasars.} % This method essentially constrains the possible shape of the power spectrum of density fluctuations at comoving scales $\\sim$Mpc. Assuming \\emph{a DM model}, (i.e. a particular primordial velocity distribution of DM particles), one can obtain a relationship between the DM particle mass \\emph{in this model} and the shape of the power spectrum, probed by Ly-$\\alpha$.  Although very promising, the Ly-$\\alpha$ method is very complicated and indirect. First of all, under the assumption that the distribution of the neutral hydrogen traces that of the DM, one can reconstruct the power spectrum of density fluctuations at redshifts $z \\sim 2-5$ from the statistics of Lyman-$\\alpha$ absorption lines.  One can then perform a fit of the Lyman-$\\alpha$ data (often together with the measurements of anisotropy of temperature of cosmic microwave background and the data of large-scale structure surveys), to extract the information about cosmological models. This is usually done by using the Monte-Carlo Markov chain technique~\\cite{Lewis:02}.  At redshifts probed by Ly-$\\alpha$, the evolution of structure has already entered the (mildly) non-linear stage. Therefore, to properly relate the measured power spectrum with the parameters of a given cosmological model one would have to perform a prohibitively large number of hydrodynamic numerical simulations. Therefore, various simplifying approximations have to be realized~\\cite{Theuns:98,Gnedin:01,McDonald:05,Viel:04,Viel:05,Viel:05b,Viel:05c,Regan:06a}.  Apart from these computational difficulties, the physics entering the \\lya analysis is complicated, and not yet fully understood~(see e.g.~\\cite{Kim:07a,Bolton:07a,Viel:2002ui,Viel:2001hd,Viel:2003fx}). Moreover, the DM particles can significantly influence the background physics, further complicating the Ly-$\\alpha$ analysis~\\cite{Biermann:06,Gao:07,Stasielak:07}. For a recent overview of the \\lya method see e.g.~\\cite{Boyarsky:08c}.    The systematic uncertainties associated with both computational difficulties and complicated physics of \\lya systems are not fully explored.  Therefore, \\emph{it is very important to have a lower mass bound on DM particles from more direct and simple considerations}.  In this paper we discuss the Tremaine-Gunn and related DM mass bounds based on the phase-space density considerations as well as possible ways to strengthen them for several DM models.  The obtained phase-space density bounds are weaker yet comparable with Ly-$\\alpha$ bounds and therefore provide an interesting alternative. We consider a class of the so-called ``generic'' DM models, where DM particles are produced thermally and decouple while being relativistic, thus having the (relativistic) Fermi-Dirac momentum spectrum.  We also consider models of non-thermal DM production. In this case the primordial velocity spectrum of DM particles depends on the details of the production mechanism.  We analyze the case when the velocity spectrum can be approximated by the \\emph{rescaled} Fermi-Dirac spectrum, or has two such components (a colder and a warmer one).  A very important example of such a DM particle is the \\emph{sterile} (right-handed) neutrino.  Although known as a DM candidate for some 15 years~\\cite{Dodelson:93}, recently sterile neutrinos have attracted a lot of attention. It was shown~\\cite{Asaka:05a} that if one adds three right-handed (sterile) neutrinos to the Standard Model, it is possible to explain simultaneously the data on neutrino oscillations (see e.g.~\\cite{Fogli:05,Strumia:06,Giunti:06} for a review) and the Dark Matter in the Universe, without introducing any new physics \\emph{above electro-weak scale} $M_W \\sim 100\\gev$.  Moreover, if the masses of two of these particles are between $\\sim 100 \\MeV$ and electro-weak scale and are almost degenerate, it is also possible~\\cite{Asaka:2005pn} to generate the correct baryon asymmetry of the Universe (see e.g.~\\cite{Dolgov:97,Riotto:98}). The third (lightest) sterile neutrino can have mass in keV-MeV range\\footnote{There are several interesting astrophysical applications of $\\keV$ sterile neutrinos (see e.g. \\cite{Sommer:99,Kusenko:06a,Biermann:06,Hidaka:06,Hidaka:07,Stasielak:06} and references therein).} % and be coupled to the rest of the matter weakly enough to provide a viable (\\emph{cold} or \\emph{warm)} DM candidate.  This theory, explaining the three observed phenomena ``beyond the SM'' within one consistent framework, is called the \\emph{$\\nu$MSM}~\\cite{Asaka:2005pn,Asaka:05a} (see also~\\cite{Shaposhnikov:07b}).  Although weakly coupled, the DM sterile neutrino in the $\\nu$MSM can be produced in the correct quanities to account for all of the DM. There are several mechanisms of production: non-resonant active-sterile neutrino oscillations (\\textbf{\\emph{non-resonant production}} mechanism, \\textbf{NRP}) ~\\cite{Dodelson:93,Dolgov:00,Abazajian:01a,Asaka:06b,Asaka:06c}, resonant active-sterile neutrino oscillations in the presence of lepton asymmetry (\\textbf{\\emph{resonant production}} mechanism, \\textbf{RP})~\\cite{Shi:98,Abazajian:01a,Shaposhnikov:08a,Laine:08a}, decay of the gauge-singlet scalar field~\\cite{Shaposhnikov:06} (see also \\cite{Kusenko:06a,Petraki:07,Petraki:08}).  The \\lya analysis of the sterile neutrino DM, produced via NRP scenario was performed in a number of works~\\cite{Viel:06,Seljak:06,Viel:08}. These bounds were recently revisited in~\\cite{Boyarsky:08c}, using the SDSS \\lya dataset together with WMAP5~\\cite{Dunkley:2008ie}.  The lower bound on the DM mass was found to be in this case $8\\kev$ (at $99.7\\%$~CL). Ref.~\\cite{Boyarsky:08c} also analyzed a more general case CWDM case : a mixture of NRP sterile neutrino with cold DM (see also~\\cite{Palazzo:07}). These results were applied to the RP produced sterile neutrino in~\\cite{Boyarsky:08d}. It was shown that the mass as low as $2\\kev$ is compatible with \\lya data.  In this paper we will analyze in detail restrictions on the sterile neutrinos, produced via first two production mechanisms and briefly comment on the third one in the Discussion.  In the case of non-resonant production the primordial velocity spectrum is approximately \\emph{proportional} to the Fermi-Dirac distribution~\\cite{Dolgov:00,Hansen:01} (the exact spectrum was calculated in ~\\cite{Asaka:06b,Asaka:06c}).  The paper is organized as follows. In Section~\\ref{sec:overview} we review DM mass bounds, based on the phase-space density arguments. In Section~\\ref{sec:maxim-coarse-grain} we introduce the concept of maximal coarse-graining and propose a conservative modification of the original Tremaine-Gunn bound. In Section~\\ref{sec:analys-meas-valu} we analyze new observational data on recently discovered dSphs (see~\\cite{Gilmore:07,Simon:07} and references therein), and use it to determine the phase-space density of these objects.  Special attention is paid to determine various systematic uncertainties of measured values. Our results are summarized in Section~\\ref{sec:bounds}. We conclude with the discussion of the results, analysis of possible uncertainties and outlook for the further improvement of the mass bounds in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In this paper we suggested that a conservative way to put the bound on a DM particle mass may be based on the requirements that the maximum of the observed coarse-grained phase space density should not exceed the maximum of the initial distribution function of the DM particles. The maximum of the coarse-graned distribution function in the final state may be conservatively estimated from the observed quantities. This bound relies on the assumption that the maximum of the distribution function was not significantly increased by the interaction with baryons.  Although DM consists of the non-interacting particles, the remaining part of the galaxy -- the baryons -- interact with one another and dissipate their energy, finally concentrating towards the center. The baryons, which are condensed in the center, influence the shape of DM halo gravitationally, increasing the central DM density~\\cite{Blumenthal:86,Gnedin:04}. The opposite effect is the energy feedback from SNae, galactic winds and reionization, which creates the strong outflow, significantly decreasing the mass of the gas and thereby affecting the DM halo shape. Such a feedback is thought to be responsible to the formation of dwarf spheroidals from gas-rich dwarf spiral/irregular galaxies~\\cite{Lin:83,Moore:94,Mastropietro:05,Mayer:06,Mayer:07}. Clearly both gas condensation and feedback strongly influence the central PSD of DM~\\cite{Read:06}, and in principle can lead to the violation of the inequality (\\ref{ff}).  Numerical studies of galaxy mergers show that baryons can lead to the increase of the phase-space density during the merger (see e.g.~\\cite{Naab:07a}).  However, the method used in this work -- coarse-graining of the PSD over a large phase-space region -- reduces the influence of baryons. Indeed, we take the spatial averaging over the radius $R \\sim r_h$, which includes external part of the system, where the amount of baryons is small.  Additional studies are necessary to estimate effects of baryons and make our bounds more robust. We plan to address these issues elsewhere.  We would also like to stress that the initial velocities of DM particles in our approach are \\emph{thermal} velocities and they should not be confused with the so-called \\emph{Zeldovich} velocities~\\cite{Zeldovich:70}.  Numerical simulations of galaxy formation do not start at the time, when the DM phase-space distribution is spatially uniform (redshifts $z\\gtrsim 10^3$). Instead, the initial (linear) stage of the structure formation is computed analytically in the framework of the so-called \\emph{Zeldovich approximation}~\\cite{Zeldovich:70}.  This approximation is commonly used to set up initial conditions for the numerical simulations of non-linear stage of structure formation~\\cite{Bertschinger:95,Klypin:97,Klypin:00}, which start at redshifts $z\\sim 10$. The peculiar (\\emph{Zeldovich}) velocities acquired by DM particles at this stage due to structure formation and included into the initial conditions are normally $\\sigma \\sim 10\\km/\\sec$. Apart from Zeldovich velocities, DM particles also possess thermal velocities, which are discussed in this paper.  For cold enough Dark Matter these thermal velocities are much smaller than Zeldovich ones and, thus, are often neglected and not included into initial conditions.  Therefore, the numerical studies of PSD evolution\\footnote{Most of these studies use the quantity $Q(r) = \\rho(r)/\\sigma^3(r)$ as a PSD estimator} (see e.g.~\\cite{Taylor:01,Peirani:06,Peirani:07,Romano-Diaz:07,Hoffman:07,Romano-Diaz:06}) essentially investigate the change of PSD from Zeldovich to final stage.  It was found in some of these works that the PSD changes by $10^2-10^3$ in the process of collapse~\\cite{Peirani:06}.  This change of PSD can be understood as being simply an evolution from initial Zeldovich velocities $\\sigma_i \\sim 10\\km/\\sec$ to the final (virial) ones $\\sigma_f \\sim 10^2\\km/\\sec$ (with $Q_i/Q_f \\sim (\\sigma_f/\\sigma_i)^3 \\sim 10^3$).  Because initial thermal velocities may be much smaller than Zeldovich ones, initial PSD may differ from the final (observed) PSD not by 2--3, but by many orders of magnitude. This fact does not contradict to the results of simulations, described in e.g.~\\cite{Peirani:06} and, therefore, cannot be used to obtain an \\emph{upper} bound on the mass of DM particles (c.f~\\cite{Boyanovsky:08,Gorbunov:08a}).    This work was mostly concentrated on restrictions on the mass of the sterile neutrino DM, produced in through the non-resonant oscillations with active neutrino (NRP scenario). We see that our results (Section~\\ref{sec:bounds}) strongly disfavor such sterile neutrinos as the single DM component. This conclusion is not based on the \\lya method and therefore is not subject to its uncertainties (discussed in the Introduction).  However, several uncertainties can affect this conclusion, the major being baryonic feedback. To make this result really robust, apart from further modeling of the baryonic influence, one needs to strengthen the tension between upper and lower mass bounds discussed in this paper.  This is plausible and may be done either by improving the X-ray bounds with new observations or by strengthening the PSD consideration, which is in the first place related to better measurements of kinematics of dSphs.  In the presence of lepton asymmetry, the resonant production (RP) of sterile neutrino DM takes place~\\cite{Shi:98}. This mechanism is more efficient~\\cite{Shi:98,Laine:08a,Shaposhnikov:08a} than the NRP scenario and allows to achieve required DM abundance for weaker mixings (c.f. Fig.~4 in~\\cite{Laine:08a}).  This lifts the upper bound on the DM particle mass in this scenario up to $\\sim 50$~keV.  At the same time, for the same mass the primordial velocity distribution of RP sterile neutrino DM is colder than in NRP one.  This $f_{max}$ is as much as the order of magnitude bigger than~(\\ref{eq:16}) (c.f.~\\cite{Laine:08a}).  This brings down by a factor $\\sim 2$ the analog of the mass bound~(\\ref{eq:29}).  Analyzing available spectra for a range of lepton asymmetries, we see that models with $m_\\sf \\gtrsim 1 \\kev$ are allowed. Thus, there is a large open ``window'' of allowed DM masses (c.f.  Fig.\\ref{fig:sf-window}).  However, as the dependence of the velocity spectrum on the lepton assymetry is not monotonic, to obtain the exact shape of the lower bound on the mass at given mixing angle more work is needed. Nevertheless, our results show that the sterile neutrinos, produced in the presence of lepton asymmetry, are viable DM candidates, allowed by all current bounds.  {Finally, we would like to comment on the mechanism of production of sterile neutrinos from decay of massive scalar field, for example the inflaton~\\cite{Shaposhnikov:06} (for other models see~\\cite{Kusenko:06a,Petraki:07,Petraki:08,Gorbunov:08a}).  The primordial phase-space distribution function for this case was computed e.g. in~\\cite{Shaposhnikov:06,Petraki:07,Boyanovsky:08b,Gorbunov:08a}.  Maximal value of phase-space density for this distribution is that of degenerate Fermi gas. Notice that the distribution functions in~\\cite{Shaposhnikov:06,Petraki:07,Gorbunov:08a} $f(p)$ is formally unbounded for small momenta: $f(p) \\sim p^{-1/2}$.  From this one can easily find that the fraction of particles, having maximal phase-space density, is $\\sim 10^{-8}$.  As only this small fraction of all particles has maximal phase-space density, we expect the mass bound in this case to be stronger than~(\\ref{eq:28}). The detailed analysis will be presented elsewhere.  }      \\bigskip  After this work has been completed, we received a draft of the paper~\\cite{Gorbunov:08b}, where similar issues have been considered. Our results are consistent with those of discussed in~\\cite{Gorbunov:08b} wherever they overlap."
    },
    "0808/0808.3419_arXiv.txt": {
        "abstract": "The equation of state (EoS) of dark energy $w$ remains elusive despite enormous experimental efforts to pin down its value and its time variation. Yet it is the single most important handle we have in our understanding of one of the most mysterious puzzle in nature, dark energy. This letter proposes a new method for measuring the EoS of dark energy by using the gravitational waves (GW) of black hole binaries. The method described here offers an alternative to the standard way of large scale surveys.  It is well known that the mass of a black hole changes due to the accretion of dark energy but at an extremely slow rate. However, a binary of supermassive black holes (SBH) radiates gravitational waves with a power proportional to the masses of these accreting stars and thereby carries information on dark energy. These waves can propagate through the vastness of structure in the universe unimpeded. The orbital changes of the binary, induced by the energy loss from gravitational radiation, receive a large contribution from dark energy accretion. This contribution is directly proportional to $(1+w)$ and is dominant for SBH binaries with separation $R \\ge 1000$ parsec, thereby accelerating the merging process for $w > -1$ or ripping the stars apart for phantom dark energy with $w < -1$. Such orbital changes, therefore $w$, can be detected with LIGO and LISA near merging time, or with X-ray and radio measurements of Chandra and VLBA experiments. ",
        "introduction": "\\label{sec:intro}  One of our most crucial questions about nature at present is: what is dark energy? The fact that our universe is accelerating \\cite{wmap,lauramelchiorri} and that dark energy constitutes about $70 percent$ of the total energy density of the universe, are well established by now. Many theoretical models have been put forth which cast dark energy in the form of a cosmic fluid \\cite{quintessence,kessence,transplanck} with time variations in its equation of state $w(z)$. Yet the simplest explanation for dark energy remains to be a pure cosmological constant (cc) $\\Lambda$.  The trouble we face in understanding dark energy does not stem from a shortage of dark energy models, with $w(z_0) \\simeq -1$, that mimick at present the behavour of $\\Lambda$ and give rise to the observed acceleration of the universe. The puzzle rather lies on identifying which one of these possible candidates is the correct one. The best way to discriminate among the various possibilites and a pure $cc$, $\\Lambda$, is to experimentally measure the time variations of the dark energy equation of state $w(z)$. So far a popular parametrization for $w(z)$ is the linear one $w(z)=w_0 + w_{1} z +...$ with $w_1 = 0$ for a pure {\\it cc} \\cite{eos}. A knowledge of $w(z)$ is crucial for not only understanding the present accelerated expansion of the universe but also for making predictions for its future evolution and destiny. If $w(z) \\ll -1$ then dark energy is a phantom \\cite{phantom} which leads the universe to a Big Rip in the future. If $w = -1$ then we are probably \\cite{laurads} facing an eternal DeSitter state \\cite{eternalds} which at least at the classical level implies constant temperature and entropy therefore a cosmic heat death \\cite{fredlaura, fred}. Other forms of $w(z)$ can also allow for a Big Crunch \\cite{crunch} or bounces \\cite{bounce}. At present we can not infer which destiny our universe will meet without a better knowldege of $w(z)$.  Major experimental efforts for pinning down the value and time-variations of $w(z)$ are under way through large scale surveys from CMB \\cite{wmap}, large scale structure \\cite{SDSS}, and SN1a observations. The endeavor of measuring, to confident precision, such small time variations in $w(z)$ has proven extremely difficult, partly due to the inherited errors in the experiments that are not instrumental but which originate from noise accumulated from the background and foreground effects through which the signal we receive has propagated. In order to minimize such errors associated with the propagation of the signal through the vastness of structure in the universe, we would like to propose in this letter a complimentary method for measuring $w(z)$. This method uses compact and localized objects, such as Black Holes, for acquiring information about $w(z)$,  by exploiting the gravitational waves these objects emit when they are in binaries. The advantage of this method is twofold: first, gravitational waves propagate undisturbed through structure; and second, we have existing experiments which are either already operational or will be in the near future, such as LIGO and LISA missions designed to detect these binaries gravitational waves, or Chandra and VLBA experiments designed for X-ray and radio measurements.  The accretion of dark energy by Black Holes is reviewed in Sec.2., including a review of the main parameters of binaries and gravitational waves, useful for our purposed. Sec.3presents the method we propose, along with an investigation, discussion and some illustrations, on how gravitational waves from SBH's binaries can be used for extrapolating the equation of state $w(z)$ of dark energy. ",
        "conclusions": ""
    },
    "0808/0808.3713_arXiv.txt": {
        "abstract": "{Characterising the circumstellar dust around nearby main sequence stars is a necessary step in understanding the planetary formation process and is crucial for future life-finding space missions such as ESA's \\darwin{} or NASA's Terrestrial Planet Finder (TPF). Besides paving the technological way to \\darwin/TPF, the space-based infrared interferometers \\peg{} and FKSI (Fourier-Kelvin Stellar Interferometer) will be valuable scientific precursors.} {We investigate the performance of \\peg{} and FKSI for exozodiacal disc detection and compare the results with ground-based nulling interferometers.} {We used the GENIEsim software (Absil et al. 2006) which was designed and validated to study the performance of ground-based nulling interferometers. The software has been adapted to simulate the performance of space-based nulling interferometers by disabling all atmospheric effects and by thoroughly implementing the perturbations induced by payload vibrations in the ambient space environment.} {Despite using relatively small telescopes ($\\leq$ 0.5\\,m), \\peg{} and FKSI are very efficient for exozodiacal disc detection. They are capable of detecting exozodiacal discs respectively 5 and 1 time as dense as the solar zodiacal cloud and they outperform any ground-based instrument. Unlike \\peg, FKSI can achieve this sensitivity for most targets of the \\darwin/TPF catalogue thanks to an appropriate combination of baseline length and observing wavelength. The sensitivity of \\peg{} could, however, be significantly boosted by considering a shorter interferometric baseline length.} {Besides their main scientific goal (characterising hot giant extrasolar planets), the space-based nulling interferometers \\peg{} and FKSI will be very efficient in assessing within a few minutes the level of circumstellar dust in the habitable zone around nearby main sequence stars down to the density of the solar zodiacal cloud. These space-based interferometers would be complementary to Antarctica-based instruments in terms of sky coverage and would be ideal instruments for preparing future life-finding space missions.} ",
        "introduction": "\\begin{figure*}[t] \\centering \\includegraphics[width=7cm]{pegase_CV3.eps}\\hspace{1.5 cm} \\includegraphics[width=6cm]{FKSI.eps} \\caption{Left: overview of the \\peg{} space-based interferometer. Two 0.4-m siderostats are flying in a linear configuration with the beam combiner spacecraft located in the middle of the formation. Right: representation of FKSI, showing the two 0.5-m siderostats located on a 12.5-m boom.} \\label{fig:pegase_FKSI} \\end{figure*}  Nulling interferometry is the core technique of future life-finding space missions such as ESA's \\darwin{} \\citep{Fridlund:2006} and NASA's Terrestrial Planet Finder Interferometer \\citep[TPF-I,][]{Beichman:2006}. Observing in the mid-infrared (6-20\\,$\\mu$m), these missions would enable the spectroscopic characterisation of the atmosphere of habitable extrasolar planets orbiting nearby main sequence stars. This ability to study habitable distant planets strongly depends on the density of exozodiacal dust in the inner part of circumstellar discs, where the planets are supposed to be located. In particular, the detection of habitable terrestrial planets would be seriously hampered for stars presenting warm ($\\sim$300\\,K) exozodiacal dust more than 10 to 100 times as dense as our solar zodiacal disc, depending on stellar type, stellar distance and telescope diameter \\citep{beichman:2006b,defrere:2008}. Assessing the level of circumstellar dust around nearby main sequence stars is therefore a necessary pre-requisite for preparing the observing programme of \\darwin/TPF by reducing the risk of wasting time on sources for which exozodiacal light prevents Earth-like planet detection. In addition, the existence of planets is intrinsically linked to circumstellar discs and observing them provides an efficient way to study the formation, evolution and dynamics of planetary systems. At young ages, essentially all stars are surrounded by protoplanetary discs in which the planetary systems are believed to form \\citep{Meyer:2008}. In particular, the detection of gaps in these protoplanetary discs is very important for understanding the early dynamics of planets, including migration and orbital interaction. At older ages, photometric surveys primarily with IRAS, ISO, and Spitzer have revealed the presence of micron-sized grains around a large number of main sequence stars \\citep[see e.g.,][]{Trilling:2008,Hillebrand:2008}. This is interpreted as the sign of planetary activity, as the production of grains is believed (by analogy with the zodiacal cloud in our solar system) to be sustained by asteroid collisions and outgassing of comets in the first tens of astronomical units (AU). However, the presence of warm dust can generally not be unequivocally determined because the typical accuracy on both near-infrared photometric measurements and photospheric flux estimations is a few percent at best, limiting the sensitivity to typically 1000 times the density of our solar zodiacal cloud \\citep{Beichman:2006c}. Photometric measurements are therefore generally not sufficient to probe the innermost regions of the discs and interferometry is required to separate the starlight from the disc emission. Good examples are given by the detection of hot dust ($\\sim$1500\\,K) around Vega and $\\tau$ Cet with near-infrared interferometry at the CHARA array \\citep{Absil:2006b,Difolco:2007}. Nulling interferometry is a quite new technique even though it was initially proposed in 1978 \\citep{Bracewell:1978}. Several scientific observations using this technique have recently been carried out with the Bracewell Infrared Nulling Cryostat (BLINC, \\citealt{Hinz:2000}) instrument at the Multi-Mirror Telescope (MMT, Mont Hopkins, Arizona), with the Keck Interferometer Nuller (KIN, Hawaii, \\citealt{Serabyn:2006,Barry:2008,Serabyn:2008}), and are foreseen to begin in 2010 at the Large Binocular Telescope (Mount Graham, Arizona,\\citealt{Hinz:2008}). In Europe, ESA has initiated the study of a ground-based demonstrator for \\darwin, the Ground-based European Nulling Interferometer Experiment \\citep[GENIE,][]{gondoin:2004}. GENIE is a nulling interferometer conceived as a focal instrument for the VLTI which has been studied by ESA at the phase A level. Another European project is ALADDIN (Antarctic L-band Astrophysics Discovery Demonstrator for Interferometric Nulling, \\citealt{Foresto:2006}), a nulling interferometer project for Dome C, on the high Antarctic plateau. The performance of GENIE has been studied in detail (\\citealt{Absil:2006_paper1},\\citealt{Wallner:2006}) and recently compared to that of ALADDIN \\citep{Absil:2007_paper2}. Using 1-m collectors, ALADDIN would have an improved sensitivity with respect to GENIE working on 8-m telescopes, provided that it is placed above the turbulence boundary layer (about 30\\,m at Dome C). Circumstellar discs 30 times as dense as our local zodiacal cloud could be detected by ALADDIN around typical \\darwin/TPF targets in an integration time of few hours.  The low atmospheric turbulence on the high Antarctic plateau is a significant advantage with respect to other astronomical sites and one of the main reasons for the very good sensitivity of ALADDIN. However, as for any other ground-based site, the atmosphere effects (turbulence and thermal background) are still major limitations to the performance and active compensation by real-time control systems are mandatory. Observing from space would provide an efficient solution to improve the sensitivity by getting rid of the harmful effect of the atmosphere. Two infrared nulling interferometers could achieve the detection of circumstellar dust discs from space (see Fig.~\\ref{fig:pegase_FKSI}): \\peg, a two-telescope interferometer based on three free-flying spacecraft \\citep{Leduigou:2006} and the Fourier-Kelvin Stellar interferometer (FKSI), a structurally-connected interferometer also composed of two telescopes \\citep{Danchi:2006}. These two missions have been initially designed to study hot extrasolar giant planets at high angular resolution in the near- to mid-infrared regime (respectively 1.5-6.0 $\\mu$m and 3.0-8.0 $\\mu$m). Besides their main scientific goal, they could also be particularly well suited for the detection of warm circumstellar dust in the habitable zone around nearby main sequence stars. The objective would be to provide a statistically significant survey of the amount of exozodiacal light in the habitable zone around the \\darwin/TPF targets, and its prevalence as a function of other stellar characteristics (age, spectral type, metallicity, presence of a cold debris disc, etc). Following our performance studies of ground-based instruments such as GENIE at Cerro Paranal (\\citealt{Absil:2006_paper1}, hereafter Paper\\,I) or ALADDIN on the high Antarctic plateau (\\citealt{Absil:2007_paper2}, hereafter Paper\\,II), the present study addresses the performance of space-based nulling instruments for exozodiacal disc detection. We have limited our comparison to instruments working at similar wavelengths (ranging from 2 to 8 $\\mu$m), and purposely discarded ground-based instruments working in the N-band such as the KIN and the LBTI. The ultimate performance of these two mid-infrared instruments essentially depends on the spatial and temporal fluctuations of the sky and instrumental thermal backgrounds, which are very difficult to model with a sufficient accuracy for our comparative study. ",
        "conclusions": "Nulling interferometry is a promising technique to assess the level of circumstellar dust in the habitable zone around nearby main sequence stars. From the ground, instruments like GENIE (VLTI nuller, using two 8-m telescopes) and ALADDIN (Antarctic nuller, using two dedicated 1-m telescopes) could achieve the detection of exozodiacal discs with a density of several tens of zodis. The high Antarctic plateau is a particularly well suited site in that context, so that ALADDIN is expected to achieve the best sensitivity (down to 30 zodis in few hours of integration time). Observing from space provides the solution to go beyond this sensitivity by getting rid of the high thermal background constraining ground-based observations. In this paper, we have investigated the performance of two space-based nulling interferometers which have been intensively studied during the past few years (namely \\peg{} and FKSI). Even though they have been initially designed for the characterisation of hot extrasolar giant planets, \\peg{} and FKSI would be very efficient to probe the inner region of circumstellar discs where terrestrial habitable planets are supposed to be located. Within a few minutes, \\peg{} (resp.\\,FKSI) could detect exozodiacal discs around nearby main sequence stars down to a density level of 5 (resp.\\,1) times our solar zodiacal cloud and thereby outperform any ground-based instrument. FKSI can achieve this sensitivity for most targets of the \\darwin/TPF catalogue while \\peg{} becomes less sensitive for the closest targets with detectable density levels of about 40 times the solar zodiacal cloud. This outstanding and uniform sensitivity of FKSI over the \\darwin/TPF catalogue is a direct consequence of the short baseline length (12.5\\,m) used in combination with an appropriate observing wavelength of about 8\\,$\\mu$m, which is ideal for exozodiacal disc detection. Another advantage of FKSI is to be relatively insensitive to the uncertainty on stellar angular diameters, which is a crucial parameter driving the performance of other nulling interferometers. In terms of sky coverage, we show that these space-based instruments are able to survey about 50\\% of the \\darwin/TPF target stars. The sky coverage reaches 80\\% if they are used in combination with ALADDIN, which provides a complementary sky coverage. Beyond the technical demonstration of nulling interferometry in space, the present study indicates that \\peg{} and FKSI would be ideal instruments to prepare future life-finding space missions such as \\darwin/TPF.  \\begin{figure}[t] \\centering \\includegraphics[width=9.2cm]{sky_coverage.eps} \\caption{Sky coverage after 1 year of observation of GENIE (dark frame), ALADDIN (light frame) and \\peg{} (shaded area) shown with the \\darwin/TPF all sky target catalogue. The blue-shaded area shows the sky coverage of a space-based instrument with an ecliptic latitude in the [-30\\degre,30\\degre] range (such as \\peg). The sky coverage of FKSI is similar to that of \\peg{} with an extension of 40\\degre\\,instead of 60\\degre.}\\label{Fig:skycov} \\end{figure}"
    },
    "0808/0808.1716_arXiv.txt": {
        "abstract": "In the solar convection zone, rotation couples with intensely turbulent convection to drive a strong differential rotation and achieve complex magnetic dynamo action.  Our sun must have rotated more rapidly in its past, as is suggested by observations of many rapidly rotating young solar-type stars. Here we explore the effects of more rapid rotation on the global-scale patterns of convection in such stars and the flows of differential rotation and meridional circulation which are self-consistently established.  The convection in these systems is richly time dependent and in our most rapidly rotating suns a striking pattern of localized convection emerges.  Convection near the equator in these systems is dominated by one or two nests in longitude of locally enhanced convection, with quiescent streaming flow in between at the highest rotation rates.  These active nests of convection maintain a strong differential rotation despite their small size.  The structure of differential rotation is similar in all of our more rapidly rotating suns, with fast equators and slower poles.  We find that the total shear in differential rotation $\\Delta \\Omega$ grows with more rapid rotation while the relative shear $\\Delta \\Omega/ \\Omega_0$ decreases.  In contrast, at more rapid rotation the meridional circulations decrease in energy and peak velocities and break into multiple cells of circulation in both radius and latitude. ",
        "introduction": "Our sun is a magnetic star whose cycles of magnetic activity must arise from organized dynamo action in its interior. This dynamo action is achieved by turbulent plasma motions in the solar convection zone, which spans the outer 29\\% of the sun in radius.  Here vigorous convective motions and rotation couple to drive the differential rotation and meridional circulation. These global-scale flows are important ingredients in stellar dynamo theory, providing shear which may build and organize fields on global scales. When our sun was younger, it must have rotated much more rapidly, as is suggested both by the solar wind which continually removes angular momentum from the sun and by many observations of rapidly rotating solar-like stars.  In more rapidly rotating suns the coupling between rotation and convection is strong and must continue to drive global scales of flow. Understanding the nature of convection, differential rotation and meridional circulation in more rapidly rotating stars is a crucial step towards understanding stellar dynamos.  The manner in which the sun achieves global-scale dynamo action is gradually being sorted out. Helioseismology, which uses acoustic oscillations to probe the radial structure of the star as well as the convective flows beneath the surface, has revealed that the solar differential rotation profile observed at the surface prints throughout the convection zone with two prominent regions of radial shear.  The near-surface shear layer occupies the outer 5\\% of the sun, whereas a tachocline of shear at the base of the convection zone separates the strong differential rotation of that zone from the uniform rotation of the deeper radiative interior \\citep[e.g.,][]{Thompson_et_al_2003}.  The solar global magnetic dynamo, responsible for the 22-year cycles of activity, is now believed to be seated in that tachocline.  In such interface dynamo models, magnetic fields generated in the bulk of the convection zone are pumped into the tachocline where the radial shear builds strong toroidal magnetic fields, with magnetic buoyancy leading to loops that rise upward and erupt through the solar surface \\citep[e.g.,] []{Charbonneau_2005,Miesch_2005}.  The differential rotation plays an important role in building and organizing the global-scale fields while the meridional circulations may be important for returning flux to the base of the convection zone and advecting it equatorward, enabling cycles of magnetic activity.  By studying the coupling of rotation and convection over a range of conditions, we are likely to learn about the operation of the current solar dynamo and about the nature of dynamos operating in our sun's past and in other solar-like stars.   Observations of young solar-like stars indicate that they rotate as much as 50 times faster than the current solar rate.  Many of these more rapidly rotating suns possess strong magnetic fields. A correlation between rotation and magnetic activity is observed over a range of stellar types and populations, indicating that more rapidly rotating stars may have stronger stellar dynamos \\citep[e.g.,][]{Noyes_et_al_1984a,Patten&Simon_1996}.  Probing the nature of these dynamos and the impact of faster rotation on the internal stellar dynamics requires both accurate observations and detailed dynamical models of the stellar interiors.  The faster flows of differential rotation are much easier to detect than the relatively slow motions associated with meridional circulations; observations across the HR diagram indicate that differential rotation is a common feature in many stars. Asteroseismic observations with the Kepler and Corot missions may soon begin to constrain the internal rotation structure.  At present only measurements of surface differential rotation are available, as assessed with a variety of techniques including photometric variability \\citep{Donahue_et_al_1996, Walker_et_al_2007}, Doppler imaging \\citep{Donati_et_al_2003} and Fourier transform methods \\citep{Reiners&Schmitt_2003}.  Advances in supercomputing have enabled 3-D simulations that are beginning to capture many of the dynamical elements of the solar convection zone.  Early global-scale simulations of solar convection by \\citet{Gilman_1975,Gilman_1977,Gilman_1979} under the Boussinesq approximation were extended by the pioneering work of \\citet{Gilman&Glatzmaier_1981}.  Such global-scale simulations of solar convection conducted in full spherical shells sought to capture the largest scales of convective flows and began to study how they can establish differential rotation and meridional circulations.  However, the range of spatial and temporal scales present in solar convection are vast and thus the computational resources required by the modeling are daunting. Through recent advances in massively parallel computer architectures, solar convection simulations are now beginning to make detailed contact with the observational constraints provided by helioseismology \\citep[e.g.,][]{ Brun&Toomre_2002, Miesch_et_al_2006, Miesch_et_al_2008}. Other efforts have focused on the vigorous turbulence and the dynamo action achieved in the bulk of the solar convection zone \\citep{Brun_et_al_2004}, with recent studies beginning to include the tachocline as a region of penetrative overshoot, shear, and magnetic field amplification \\citep{Browning_et_al_2006}. Facilitated by these computational advances, models of convection and dynamo action within the cores of A-type stars have also begun to be investigated \\citep{Browning_et_al_2004, Brun_et_al_2005, Featherstone_et_al_2007}.  To date, most models of stellar differential rotation in stars like our sun that rotate more rapidly have been carried out in 2-D under the simplifying assumptions of mean field theory \\citep[e.g.,][]{Rudiger_et_al_1998, Kuker&Stix_2001, Kuker&Rudiger_2005_A&A, Kuker&Rudiger_2005_AN}.  The time is ripe to pursue the question with fully 3-D simulations of global-scale stellar convection.  In order to study solar-like stars that rotate faster, our previous work focused on a series of 3-D compressible simulations within a spherical shell using the anelastic spherical harmonic (ASH) code for stars rotating from one to five times the current solar rate \\citep{Brown_et_al_2004}.  These preliminary hydrodynamic simulations explored how stellar convection changes with more rapid rotation, including the differential rotation and meridional circulation that is achieved. Comparable studies have been carried out by \\citet{Ballot_et_al_2007} studying younger stars with deeper convection zones.  In our simulations we found that remarkable nests of vigorous convection emerge in the equatorial regions.  Namely, convective structures at low latitudes about the equator can exhibit strong spatial modulation with longitude, and at high rotation rates the convection is confined to narrow intervals (or nests) in longitude.  In the more turbulent simulations presented in this paper, which span a larger range of rotation rates from one to ten times the current solar rate, the phenomena of modulated convection has persisted and the active nests of convection are prominent features at the higher rotation rates. These nests of localized convection persist for long intervals of time and despite their small filling factor maintain a strong differential rotation.  The emergence of spatially localized convective states has been observed in other systems, particularly in theoretical studies of doubly-diffusive systems such as thermosolutal convection \\citep[e.g.,][]{Spina_et_al_1998, Batiste_et_al_2006}, in laboratory studies of convection in binary fluids \\citep[e.g.,][]{Surko_et_al_1991}, and in simulations of magnetoconvection where isolated ``convectons'' have been observed \\citep{Blanchflower_1999}. In shells of rapidly rotating fluid, temporally intermittent patches of localized convection emerged in Boussinesq simulations of the geodynamo \\citep{Grote&Busse_2000}. In many of these systems, spatial modulation occurs in the weakly nonlinear regime close to the onset of convection.  In contrast, our simulations of solar convection are in a regime of fully developed turbulent convection.  We describe briefly in \\S\\ref{sec:ASH} the 3-D anelastic spherical shell model and the parameter space explored by our simulations. In \\S\\ref{sec:convection} we discuss the nature of convection realized in more rapidly rotating stars and the emergence of spatially-localized patterns of convection.   In \\S\\S\\ref{sec:global scale flows}-\\ref{sec:energies} we examine the global-scale flows realized in our simulations, including differential rotation and meridional circulation, and their scaling with more rapid rotation. A detailed exploration of the active nests of convection is presented in \\S\\ref{sec:patches}.  We reflect in \\S\\ref{sec:conclusion} on the significance of our findings. ",
        "conclusions": "\\label{sec:conclusion}  When stars like our sun are young they rotate much more rapidly than the present sun.  In these stars rotation must strongly influence the convective motions and may lead to stronger global-scale dynamo action. We have explored here the effects of more rapid rotation on global-scale convection in simulations of stars like our sun.  The mean zonal flows of differential rotation become much stronger with more rapid rotation, scaling as $\\Delta \\Omega \\propto \\Omega_0^{0.3}$ or as $\\Delta \\Omega/\\Omega_\\mathrm{eq} \\propto \\Omega_0^{-0.6}$.  In striking contrast, the meridional circulations become much weaker with more rapid rotation, and the energy contained in them drops approximately as $\\Omega_0^{-0.9}$.  Accompanying the growing differential rotation is a significant latitudinal temperature contrast, with amplitudes of $100~\\mathrm{K}$ or higher in the most rapidly rotating cases.  The maximum temperature contrast near the surface occurs between the hot poles and the cool mid latitudes at about $\\pm 40^\\circ$.  If this latitudinal temperature contrast prints through the vigorous convection at the stellar surface, it may appear as an observable latitudinal variation in intensity.  The thermal contrast would presumably persist for long periods compared to stellar activity, offering a way to disentangle this intensity signature from that caused by spots of magnetism at the stellar poles.  These simulations are entirely hydrodynamic and this provides a major caveat to our findings on the scaling of differential rotation and latitudinal temperature contrast with rotation rate $\\Omega_0$. Prior MHD simulations of stellar convection have demonstrated that in some parameter regimes strong dynamo-generated magnetic fields can react back strongly on the differential rotation, acting to lessen angular velocity contrasts or largely eliminate them \\citep[e.g.,][]{Brun_et_al_2005, Featherstone_et_al_2007,Browning_2008}. It is unclear whether the scaling trends identified here for differential rotation as a function of $\\Omega_0$ will survive in the presence of dynamo action and magnetic fields.  Likewise, magnetic fields may lessen the strong temperature contrasts realized here, doing so through their feedback on the convective flows and energy transport. We expect that the weaker meridional circulations may be less affected by magnetic feedbacks, and thus the prediction that meridional circulations lessen in energy and amplitude with more rapid rotation may be of greater significance though harder to confirm observationally.  Weaker meridional circulations in more rapidly rotating stars will have a strong impact on many theories of stellar dynamo action, including the Babcock-Leighton flux-transport model favored for solar-type stars. We have begun simulations to explore the dynamo action realized at various rotation rates, and its impact upon the flows described here. Preliminary results appear in \\cite{Brown_et_al_2007c} and more detailed results will be forthcoming shortly.  A striking feature of these simulations is the emergence of a pattern of strongly modulated convection in the equatorial regions.  These nests of active convection are regions of enhanced convective vigor and transport which propagate at rates distinct from either the mean zonal flows of differential rotation or the individual convective cells.  In the most rapidly rotating systems, such as case~G10, convection at the equator is entirely dominated by motions inside the nest with only very weak radial motions present in the regions outside the nest.  Though their impact on the convection is most obvious in the rapidly rotating limit, we find some evidence for weak modulation even in our more slowly rotating cases.   All of our simulations stop short of the turbulent stellar surface, and it is thus difficult to estimate how these nests of active convection may affect stellar observations in detail.  The convective velocities associated with the nests are small compared to the nearly supersonic flows in stellar granulation, and in the sun such global-scale convective structures have evaded direct detection despite intensive searches throughout the near-surface layers by helioseismology.  The extremely localized nests found in our most rapidly rotating cases may however influence the thermal stratification and thus convective vigor in the near surface regions, as most of the flux at the equator is transported through a narrow range of longitudes.  These nests may act as traveling hot spots with enhanced convection even in the surface layers where the higher emerging flux escapes the system.  These spatially localized states of convection may also have some bearing on the active longitudes of magnetic activity observed in the sun, if the enhanced pummeling of the tachocline by the convection within the nest preferentially destabilizes magnetic structures within the tachocline that then rise to the surface.  Initial dynamo simulations indicate that nests of convection can coexist with magnetism in portions of parameter space.  Thus their strongest signature is likely to emerge in magnetic stars, where magnetic fields threading the bulk of the convection zone may be concentrated in the nests and mimic giant, propagating starspots which survive for very long epochs in time.  If the nests lead to active longitudes of enhanced magnetic activity in rapidly rotating stars, we might expect these long-lived magnetic structures to propagate at a rate different from the stellar rotation rate as measured at the surface or from the stellar differential rotation.  We recognize that our simulations remain separated by many orders of magnitude from the parameter space of real stellar convection.  As such, we must be cautious with our interpretations of the overall dynamics.  However, we have found that these nests of convection are a robust feature over a range of parameters and that they are able to persist as entities for as long as we could pursue their modelling.  Thus one should be prepared to consider the possibility of their presence also in real stellar convection zones, where they may appear as long-lived propagating features."
    },
    "0808/0808.2838_arXiv.txt": {
        "abstract": "Here we present a detailed analysis of solar acoustic mode frequencies and their rotational splittings for modes with degree up to 900. They were obtained by applying spherical harmonic decomposition to full-disk solar images observed by the Michelson Doppler Imager onboard the {\\em Solar and Heliospheric Observatory} spacecraft. Global helioseismology analysis of high-degree modes is complicated by the fact that the individual modes cannot be isolated, which has limited so far the use of high-degree data for structure inversion of the near-surface layers ($r > 0.97 R_{\\odot}$). In this work, we took great care to recover the actual mode characteristics using a physically motivated model which included a complete leakage matrix. We included in our analysis the following instrumental characteristics: the correct instantaneous image scale, the radial and non-radial image distortions, the effective position angle of the solar rotation axis and a correction to the Carrington elements. We also present variations of the mode frequencies caused by the solar activity cycle. We have analyzed seven observational periods from 1999 to 2005 and correlated their frequency shift with four different solar indices. The frequency shift scaled by the relative mode inertia is a function of frequency alone and follows a simple power law, where the exponent obtained for the $p$ modes is twice the value obtained for the $f$ modes. The different solar indices present the same result. ",
        "introduction": "\\label{S-Introduction}   The central frequencies of solar acoustic modes, which are obtained using spherical harmonic decomposition, have been successfully used to determine the solar interior structure to as close as 21 Mm to the solar surface ($r < 0.97 R_{\\odot}$) using modes with angular degrees $\\ell \\le 300$ ({\\it e.g.}, \\opencite{Gough96}). The inclusion of high-degree modes ({\\it i.e.}, up to $\\ell = 1000$) has the potential to improve dramatically the inference of the sound speed and the adiabatic exponent ($\\Gamma_1$) in the outermost 2 to 3\\% of the solar radius, allowing to construct localized kernels as close to the solar surface as 1.75 Mm \\cite{Rabello-Soares00}. The effects of the equation of state, through the ionization of hydrogen and helium, are felt most strongly in the outer layers of the Sun, making this shallow region of particular interest. Furthermore, dynamical effects of convection, and the processes that excite and damp the solar oscillations, are predominantly concentrated in this region. Although the spatial resolution of the modern helioseismic instruments allows us to observe oscillation modes up to $\\ell$ = 1000 and higher, only a small fraction of them are currently used ($\\ell \\le 300$). Unfortunately, analysis of high-degree data is complicated by the fact that the individual modes cannot be isolated ({\\it e.g.}, \\opencite{RKS01}).   The solar structure is not static, but changes over the solar cycle. It is well known that the mode frequencies change with solar activity. It seems that the {\\color{black} responsible} % mechanism is restricted to the outer layers of the Sun (\\opencite{LW90}), where the high-degree modes are confined. However, at the moment, there is no general agreement in the precise physical mechanism that gives rise to the frequency variation. It is likely a product of the change in the subphotospheric small-scale magnetic field strength with the solar activity cycle ({\\it e.g.}, \\opencite{Goldreich91}). Accordingly to \\inlinecite{DG04}, the frequency shift is easily explained in terms of a variation in the turbulent velocities associated with the magnetic field variation rather than the sole direct effect of the magnetic field itself. \\inlinecite{Li03}, using models of the structure and evolution of the Sun, found that turbulence near the surface of the Sun plays a major role in solar variability, and only a model that includes a magnetically modulated turbulent mechanism can agree with the observed correlation between the frequency shift and the solar cycle. In such a dynamic model, the evolution of the subsurface layers of the Sun {\\color{black} through} % the activity cycle plays an important role.   The frequencies of the global modes give the radius and latitude $(r,\\theta)$ part of the structure. While, local helioseismic techniques such as ring-diagram analysis \\cite{Hill88} allow the determination of the three-dimensional structure of the Sun, allowing the study of localized areas in the solar surface, such as those in active regions. Large variations of the mode frequencies observed in and near sunspots in comparison to magnetically quiet regions are well known to be correlated with variations in the average surface magnetic field between the corresponding regions. ({\\it e.g.}, \\opencite{Rajaguru01} and \\opencite{RBB07}). Whether the frequencies are changed directly by the magnetic field or indirectly through an associated change in the solar structure ({\\color{black} such as} % a pressure change) is still a matter of debate. A detailed analysis of the frequency-shift characteristics will hopefully help understand their physical origin.  \\inlinecite{Basu04}, using ring-diagram analysis, found that the sound speed is lower in the immediate subsurface layers of an active region than of  a magnetically quiet region, while the opposite is true for depths below about 7 Mm. However, \\inlinecite{Basu02}, using global analysis, have found no observable structural changes in the inner layers of the Sun below a depth of 21 Mm associated with the magnetic-activity induced frequency shifts. They were, however, unable to get closer to the solar surface due to the lack of high-degree modes in their mode set. High-degree global analysis is important to complete the picture of the near-surface layers. Besides, the determination of high-degree frequencies using different methods allows us to check the results against each other giving confidence in the results and avoiding systematic errors. We should point out that, although the high-degree modes have short lifetimes (one\\,--\\,ten % hours for $100 \\le \\ell \\le 600$ accordingly to \\opencite{Olga07}) propagating only locally, they are averaged over most of the solar surface using spherical harmonic decomposition (over a relatively long time series) and thus can still be called global analysis.  In the traditional global helioseismology data-analysis methodology, a time series of full-disk Doppler solar images is decomposed into spherical harmonic coefficients, characterized by its degree ($\\ell$) and its azimuthal order ($m$).  Each coefficient time series is Fourier transformed, and the order of the radial wave function ($n$) gets separated in the frequency domain. However, a spherical harmonic decomposition is not orthonormal over less than the full sphere -- {\\it i.e.}, the solar surface that can be observed from a single view point-- resulting in what is referred as spatial leakage.  At low and intermediate degrees, most of these leaks are separated in the frequency domain from the target mode (except for some $m$ leaks) and individual modes can be identified and fitted. However, at high degrees, the spatial leaks lie closer in frequency (due to a smaller mode separation) and, at high frequency, the modes become wider (as the mode lifetimes get smaller), resulting in the overlap of the target mode with the spatial leaks that merges individual peaks into ridges (see Figure~1 in \\opencite{RKS01}). The characteristics of the resulting ridge (central frequency, amplitude, {\\it etc}\\ldots) do not correspond to those of the underlying target mode. This has so far hindered the estimation of unbiased mode parameters at high degrees.   To recover the actual mode characteristics, we need a very good estimation of the relative amplitude of the spatial leaks present in a given $(\\ell, m)$ power spectrum, also known as the leakage matrix, which in turn requires a very good knowledge of the instrumental properties \\cite{RKS01}. In our previous papers (\\opencite{RKS01} and \\opencite{KRS04}, hereafter KRS), we described in detail the large influence of the instrumental properties on the amplitude of the leaks and as a consequence in the determination of unbiased high-degree mode parameters.   In the following, we will first describe the data used in this analysis and the ridge-to-mode correction applied to them (Sections~\\ref{useddata} and \\ref{ridgemodel}). In Section~\\ref{instreff}, we will discuss the influence on the mode parameters of each of the instrumental properties that were included in the spherical harmonic decomposition of the solar images. We then analyze in Section~\\ref{highdegree} the characteristics of the high-degree mode frequencies and their rotational splittings obtained in this work. Finally, in Section~\\ref{solarcycle}, we analyze the frequency variation induced by the solar cycle. ",
        "conclusions": "\\label{S-Conclusion}  In the determination of unbiased high-degree mode parameters, the instrumental characteristics must be taken into account in the image spatial decomposition or in the leakage matrix calculation itself to obtain a correct estimation of the relative amplitude of the spatial leaks. Among the instrumental characteristics analyzed here, the image scale is the one that affects the parameter determination the most. The image scale is the ratio of the image dimensions observed on the CCD detector and the dimensions in the actual Sun. An error in the image scale introduces an error in the estimated central frequency which increases with the mode frequency. A 0.27\\% error in the image scale would shift the estimated central frequency by as much as 11 $\\mu$Hz at 5 mHz. The radial distortion also has an important effect which is expected since it is similar to an image scale error. An instrumental property not taken into account here (due to a lack of a good estimation) that could have an important effect on the measured parameters is an azimuthally varying PSF.    The applied ridge-to-mode correction recovers frequencies at moderate degree that differ from the assumed corrected values by 1 $\\mu$Hz or less depending on the mode frequency. The fitting uncertainty of the recovered frequencies is in the range 0.07\\,--\\,0.18 $\\mu$Hz. The agreement for the $a_1$, $a_2$, and $a_3$ coefficients is very good, except maybe for $f$- and $p_1$-mode $a_1$ coefficients, their mean difference with respect to the assumed correct values is, respectively, three and two normalized by the ridge fitting uncertainties.  At high degree, the differences between our frequency determination and theoretical frequencies for the $p$ modes {\\color{black}shows} the same general variation with degree as the results obtained with ring analysis. For $n \\leq 5$, the global and ring analysis agree within 6 $\\mu$Hz. The high-degree $f$-mode frequencies obtained using ring analysis, like previous observations (\\opencite{Duvall98} and references within), are substantially lower than the theoretical frequencies. Surprisingly, the $f$-mode high-$\\ell$ set frequencies agree well with the model frequencies (within 3 $\\mu$Hz) whereas the ring-analysis frequency differences can be as large as 13 $\\mu$Hz for $\\ell > 700$ modes. The implications of the high-$\\ell$ frequencies and splitting coefficients on the solar structure and rotation will be addressed in a future paper.   {\\color{black}As noted} % by other authors for low- and moderate-degree modes ({\\it e.g.}, \\opencite{LW90}), the frequency shift induced by the solar cycle scales well with the mode inertia. We extended this analysis to high-degree modes and found out that scaling with the mode inertia normalized by the inertia of a radial mode of the same frequency follows a simple power law (given by Equation~\\ref{eq:plaw}) with one exponent at all frequency ranges, where the $f$ and $p$ modes are fitted independently of one another. The exponents obtained using four different solar indices agree within their fitting uncertainty, where: $\\bar{\\gamma_f} = 1.29 \\pm 0.07$ and {\\color{black}$\\bar{\\gamma_p} = 3.60 \\pm 0.01$}.  % The $f$-mode exponent is less than half of the $p$-mode value. The fundamental mode of solar oscillations has essentially the character of surface gravity waves and, contrary to the $p$ modes, it is essentially incompressible and independent of the hydrostatic structure of the Sun. Hence, it is not a surprise that these different types of modes have different exponents. Due to their different properties, it is also very likely that different physical effects are responsible for their frequency variation. Accondingly to \\inlinecite{DG05}, for the $f$ modes, the dominant cause of frequency shift is the variation of the subphotospheric magnetic field. For the $p$ modes, it is the decrease in the radial component of the turbulent velocity in the outer layers during the increase in solar activity, which is accompanied by a decrease in temperature (due to a decrease in the efficiency of convective transport). At low frequency ($\\nu < 2.3$ mHz), the $p$-mode frequency shifts have a different behavior than at high-$\\nu$: a step (with the same exponent $\\gamma_p$) or, as found by other authors, an exponent twice as large as the one at high-$\\nu$. Low-frequency modes have a large probability of being uncorrelated with the solar cycle, which was not taken into account in the case where a large exponent was estimated. {\\color{black} The $f$-mode frequency shifts also have a different behavior around 1.7 mHz: they increase abruptly by an order of magnitude. }    Modes with frequency around 3 mHz have the smallest probability that their variation is linearly uncorrelated with the solar index, while modes with $\\nu < 2.5$ mHz or $\\nu > 4.5$ mHz have the largest probability of being uncorrelated. A large  probability ($P_u$) of being uncorrelated does not necessarily means that a given mode is not physically correlated with the solar cycle, instead it could be due to uncertainties in the measurements, a low signal-to-noise ratio. The logarithm of $P_u$ is well correlated with the logarithm of the relative uncertainty of the estimated frequency shift $\\delta\\nu^e$, the Pearson correlation coefficient is 0.71 for medium-$\\ell$ modes and 0.54 for high-$\\ell$ modes. If a given mode has a large probability of being linearly uncorrelated, it is expected that the linear fitting of its frequency shifts will have a large uncertainty, hence the high correlation coefficient between $P_u$ and the estimated frequency shift uncertainty. However, it raises the question of what could be the physical process that would make those modes less sensitive to solar activity. For a given $\\ell$, the upper reflection point for lower-frequency modes is deeper in the Sun than for high-frequency modes. If the perturbation layer causing the frequency shift is above the upper turning point of the mode, it would not be affected by the solar cycle. Accordingly to model S of Christensen-Dalsgaard, the depth of the upper turning point increases sharply with decreasing frequency below 2.3 mHz. The upper turning point for a radial mode with $\\nu = 2$ mHz is 0.5 Mm deeper in the Sun than a three-millihertz mode (from Figure~2 in \\opencite{Cha01}). At high-frequency, the observed frequency shift seems to suddendly drop to zero. However, there is no agreement on the exact frequency that this happens, the observed values range from 3.7 to 5 mHz. The frequency-shift falloff is explained by an increase of chromospheric temperature and magnetic field at solar maximum (\\opencite{Jain96} and references within). In the presence of an inclined magnetic field, high-frequency modes tunnel through the temperature minimum and are particularly sensitive to changes in the chromosphere, which are expected to be well correlated with solar activity.   \\begin{acks} We are grateful to Tim Larson of Stanford University for discussing with us the results of his improved analysis of MDI medium-$\\ell$ data. The Solar Oscillations Investigation (SOI) involving MDI is supported by NASA grant NNG05GH14G at Stanford University.  SOHO is a mission of international cooperation between ESA and NASA. SGK is supported by NASA grant NNG05GD58G. NOAA Mg {\\sc ii} Core-to-wing ratio data are provided by Dr. R. Viereck, NOAA Space Environment Center. The solar radio 10.7 cm daily flux (2800 MHz) have been made by the National Research Council of Canada at the Dominion Radio Astrophysical Observatory, British Columbia. The International Sunspot Number was provided by SIDC, RWC Belgium, World Data Center for the Sunspot Index, Royal Observatory of Belgium. This study includes data from the synoptic program at the 150-Foot Solar Tower of the Mt. Wilson Observatory.  The Mt. Wilson 150-Foot Solar Tower is operated by UCLA, with funding from NASA, ONR, and NSF, under agreement with the Mt. Wilson Institute. This work utilizes data obtained by the Global Oscillation Network Group (GONG) program, managed by the National Solar Observatory, which is operated by AURA, Inc.  under a cooperative agreement with the National Science Foundation.  The data were acquired by instruments operated by the Big Bear Solar Observatory, High Altitude Observatory, Learmonth Solar Observatory, Udaipur Solar Observatory, Instituto de Astrof\\'isica de Canarias, and Cerro Tololo Interamerican Observatory.     \\end{acks}"
    },
    "0808/0808.0526_arXiv.txt": {
        "abstract": "{ Progress in understanding the formation and evolution of planetary nebulae (PN) has been restricted by a paucity of well-determined central star masses. To address this deficiency we aim to (i) significantly increase the number of known eclipsing binary central stars of PN (CSPN), and subsequently (ii) directly obtain their masses and absolute dimensions by combining their light-curve parameters with planned radial velocity data. Using photometric data from the third phase of the Optical Gravitational Lensing Experiment (OGLE) we have searched for periodic variability in a large sample of PN towards the Galactic Bulge using Fourier and phase-dispersion minimisation techniques. Among some dozen periodically variable CSPN found, we report here on three new eclipsing binaries: M 3-16, H 2-29 and M 2-19. We present images, confirmatory spectroscopy and light-curves of the systems. } ",
        "introduction": "The key parameter for any star is its mass. The least model-dependent method to obtain masses is via Keplerian orbits in binary systems. Eclipsing binaries provide critical information necessary to secure the masses and other fundamental stellar parameters. The mass is obtained from radial velocity (RV) orbits of both stars combined with the orbital inclination from photometric eclipses.  Yet directly derived masses have not been fully exploited in the case of central stars of planetary nebulae (CSPN), partly because so few eclipsing binary systems are known. Indirect methods, such as nebula modelling (e.g. Gesicki et al. 2006) or model atmosphere fitting (e.g. Napiwotzki 1999), offer a poor alternative to relatively model-independent techniques such as Keplerian orbits.  Obtaining direct masses for a sufficiently large number of CSPN is key to understanding the link between the AGB phase of stellar evolution and the final white dwarf stage. Such a study will also greatly improve our understanding of PN shaping mechanisms (e.g. Zijlstra 2007) and can provide compelling candidates for SN Ia progenitors (Tovmassian et al. 2008). Because most stars of low-intermediate mass experience the PN stage, PN are also major contributors to ISM enrichment and the chemical evolution of galaxies.    De Marco et al. (2008, and references therein) summarised the status of photometrically variable CSPN, including 12 close binaries with known periods. Three of these are eclipsing: UU Sge (Pollaco \\& Bell 1993), V477 Lyr (Pollaco \\& Bell 1994), and BE UMa (Ferguson et al. 1999). All three have viable mass determinations using both RV and light-curves, with 0.5--0.7 M$_\\odot$ for the hot 60--120 kK primary and 0.15--0.36 M$_\\odot$ for the irradiated 5--7 kK secondary. Their light-curves display strong irradiation effects, in addition to typical emission lines from the heated hemisphere of the secondary. All these systems are believed to be post-CE (common-envelope) binaries, where the secondary spirals in through the primary's AGB envelope. The remaining stars of De Marco et al. 2008 are non-eclipsing and therefore any mass estimates for them are subject to considerable uncertainty.      Either RV or photometric monitoring surveys may be used to \\emph{find} CSPN binaries. However, if the ultimate goal (as in this work) is to get the masses from binaries, then it is easier and more efficient to start with photometric surveys to preselect eclipsing binaries, for which RV orbits combined with photometrically-derived inclinations will secure the masses. Indeed, RV monitoring surveys have proven to be difficult with many candidates showing RV variability without definitive periodicity (e.g. De Marco et al. 2004).  This paper introduces a novel use of extant online photometric data from the OGLE microlensing survey. We describe in Section 2 a search for periodic variability in over 300 PN towards the Galactic Bulge. In Section 3 we present the discovery of three new eclipsing binary CSPN. We discuss the nature of these new discoveries in Section 4 and conclude in Section 5. ",
        "conclusions": "We have doubled the known population of eclipsing binary CSPN after searching OGLE-III photometry in a large PN catalogue towards the Galactic Bulge. The post-CE CSPN show bipolar morphologies consistent with the current hypothesis that close binaries lead to non-spherical nebulae. A following paper will present more new close binary CSPN, discuss the selection effects of the search in detail and provide a new, independent estimate the binary fraction of PN. Planned high-resolution spectroscopy will refine system parameters and derive masses of our new eclipsing binaries."
    },
    "0808/0808.2465_arXiv.txt": {
        "abstract": "We present Spitzer infrared (IR) photometry and spectroscopy of the lensed Lyman break galaxy (LBG), MS1512-cB58 at $z=2.73$.  The large (factor $\\sim 30$) magnification allows for the most detailed infrared study of an $L^*_{UV}(z=3)$ LBG to date.  Broadband photometry with IRAC (3-10 $\\mu$m), IRS (16 $\\mu$m), and MIPS (24, 70 \\& 160 $\\mu$m) was obtained as well as IRS spectroscopy spanning 5.5-35 $\\mu$m.    A fit of stellar population models to the optical/near-IR/IRAC photometry gives a young age ($\\sim9$ Myr), forrming stars at $\\sim 98$ M$_{\\odot}$ yr$^{-1}$, with a total stellar mass of $\\sim 10^9$ M$_{\\odot}$ formed thus far.  The existence of an old stellar population with twice the stellar mass can not be ruled out.  IR spectral energy distribution fits to the 24 and 70 $\\mu$m photometry, as well as previously obtained submm/mm, data give an intrinsic IR luminosity $L_{IR} = 1-2 \\times 10^{11}$ L$_{\\odot}$ and a star formation rate, SFR $\\sim20-40$ M$_{\\odot}$ yr$^{-1}$.  The UV derived star formation rate (SFR) is $\\sim3-5$ times higher than the SFR determined using $L_{IR}$ or $L_{H\\alpha}$ because the red UV spectral slope is significantly over predicting the level of dust extinction.  This suggests that the assumed Calzetti starburst obscuration law may not be valid for young LBGs.  We detect strong line emission from Polycyclic Aromatic Hydrocarbons (PAHs) at 6.2, 7.7, and 8.6 $\\mu$m.  The line ratios are consistent with ratios observed in both local and high redshift starbursts.  Both the PAH and rest-frame 8 $\\mu$m luminosities predict the total $L_{IR}$ based on previously measured relations in starbursts.  Finally, we do not detect the 3.3 $\\mu$m PAH feature.  This is marginally inconsistent with some PAH emission models, but still consistent with PAH ratios measured in many local star-forming galaxies. ",
        "introduction": "Much of the global star formation at $z>2$ occurs in ultraviolet-luminous Lyman Break Galaxies \\citep[LBGs,][]{steidel96,reddy05}.  For most LBGs, much of what we know about their star-formation (star formation rates, dust reddening) is derived from rest-frame UV properties, where considerable degeneracies exist (ie. dust reddening and starburst age).  In most $L^*_{UV}$ LBGs, the majority of the UV photons are absorbed by dust \\citep{adelberger00,reddy08}, which then emits the energy in the infrared.  Therefore, an accurate measurement of the total infrared luminosities would give a complete census of the reprocessed UV photons and thus, a better determination of the star formation rates in LBGs.  LBGs can be detected in the infrared at 24 $\\mu$m with the {\\it Spitzer Space Telescope}.  Observations in the $Spitzer$ 24 $\\mu$m bandpass are the most sensitive towards detecting dust obscured star formation at $z<3$.  However, even the deepest $Spitzer$ surveys ($f_{24}>20$ $\\mu$Jy, or $L_{IR}\\geqsim3\\times10^{11}$ L$_{\\odot}$ at $z=3$), can not detect the majority of LBGs \\citep[see IR luminosity function of][]{reddy08}.  Between $1<z<3$, the 24 $\\mu$m flux is dominated by the emission features of Polycyclic Aromatic Hydrocarbons (PAHs), resulting in uncertain bolometric corrections due to variations in the PAH emission and possible silicate absorption at 9.7 $\\mu$m.  As the IR spectra of typical LBGs have not been measured, templates of local starbursts have been used to extrapolate the rest-frame 5-12 $\\mu$m luminosities to total IR luminosities.  It is not clear, however, that the IR SEDs, and in particular the PAH emission, is the same in local and high redshift starbursts.  Much progress has been made in measuring the mid-IR spectral properties and broadband SEDs of dusty, ultraluminous star-forming galaxies at $z>1$ \\citep{yan05,menendez-delmestre07,sajina07,desai07,pope08}, and even an extremely luminous and dusty LBG \\citep{huang07}.  However, there have been no mid-IR spectra obtained of LBGs with typical luminosities ($M_{1500} \\sim -21$, $L_{IR}\\sim10^{11}$ L$_{\\odot}$) and dust extinction ($0.0<E(B-V)<0.3$).  There are a few LBGs that are sufficiently magnified through lensing for their mid-IR spectral properties and far-IR photometry to be measured \\citep[e.g.,][]{yee96,smail07,allam07}.  MS1512-cB58 (hereafter, cB58) is the first lensed LBG \\citep[$z=2.73$,][]{yee96} to be found with a large magnification \\citep[$\\sim30$,][]{williams96_cb58,seitz98}, allowing extensive follow-up at various wavelengths.  Fits of stellar population models to the optical and near-IR photometry (rest-frame UV and optical) indicate a very young ($t_{age} \\sim 10-20$ Myr), relatively dusty, $E(B-V) \\sim 0.3$, starburst \\citep{ellingson96}.  \\citet{pettini00} use the rest-frame UV continuum to determine the dust extinction and derive a reddening (and magnification) corrected star formation rate of 93 $M_{\\odot}$ yr$^{-1}$.  A Keck near-IR spectrum of nebular emission lines yields an intrinsic star formation rate of $\\sim 20$ M$_{\\odot}$ yr$^{-1}$ \\citep{teplitz00}.  Sub-millimeter \\citep{van_der_werf01,sawicki01} and millimeter \\citep{baker01} photometry also suggest, though with large uncertainty, a star formation rate significantly lower than the UV-derived SFR, calling into question the applicability of locally derived relations (reddening laws, IR SEDs) to this LBG, and perhaps the population as a whole.  However, this discrepancy can be reconciled if the IR SED is very warm, as predicted by the LBG dust emission models of \\citet{takeuchi04}.  Here we present Spitzer mid- and far-IR photometry, and mid-IR spectroscopy of cB58 to more accurately determine its IR properties and whether local starburst relations can be used for analysis of LBGs.  We use a $\\Lambda$CDM cosmology with $\\Omega_m = 0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$.  All intrinsic luminosities and star-formation rates are corrected assuming a lensing magnification, $\\mu = 30$ \\citep{seitz98}.  \\citet{seitz98} estimate an error on the magnification 30\\%. ",
        "conclusions": "We have obtained Spitzer photometry and mid-IR spectroscopy of the lensed LBG, MS1512-cB58.  This is the (intrinsically) least luminous galaxy to be studied in detail in the infrared at $z>1$.  Although cB58 is a ``typical'' LBG in terms of its UV and IR luminosities, the lack of a Balmer break in the SED indicates that it is much younger ($t_{age} \\sim 10 Myr$) than a large majority of LBGs.  As seen in previous studies \\citep{shapley05}, stellar population fits to the combined optical/near-IR/IRAC photometry give similar parameters to fits to the optical/near-IR photometry alone, but with smaller error bars.  The mid-IR photometry do not significantly constrain the mass of an older stellar population, as the rest-frame near IR flux from the current burst ($M_{burst} \\sim 10^9$ M$_{\\odot}$) can dwarf the flux from an older stellar population with twice the mass.  The far-IR photometry is reasonably well fit by local starburst templates and gives an $L_{IR} = 1-2\\times10^{11}$ L$_{\\odot}$.  In addition, the PAH luminosities and rest-frame 8 $\\mu$m luminosities agree with the $L_{PAH}$-to-$L_{IR}$ and $L_{8}$-to-$L_{IR}$ relations measured in local starbursts of comparable luminosity and high redshift, higher luminosity starbursts.  The inferred SFR from the infrared luminosity is consistent with the SFR derived from the extinction corrected $H\\alpha$ luminosity, but the UV-derived (extinction corrected) SFR is larger by a factor of 3-5.  This phenomenon has been noted by \\citet{reddy06}, where LBGs with starburst ages less than 100 Myr consistently have lower $L_{IR}$ than the UV spectral slope and luminosity suggest.  With cB58, we can be confident that this discrepancy is {\\it not} due to a poor determination of the infrared luminosity.  This suggests that the Calzetti obscuration law for starbursts may not be valid in very young, $<100$ Myr old, high redshift starbursts.  The high HI column density and large covering fraction of neutral gas suggests a dust geometry approximated by a uniform foreground sheet.  This geometry would result in a steeper reddening law like the LMC or SMC curves and would explain the overstimate of the UV obscuration using the Calzetti reddening law.  In addition, because many young ($t_{age}<100$ Myr) LBGs also exhibit spectral properites indicative of large covering fractions of neutral gas (eg. high equivalent width low ionization metal absorption lines), a steeper reddening law may also be appropriate for explaining the relatively red UV spectral slopes observed in these systems as well.  It is of course impossible to apply the relations measured in cB58 to the LBG population as a whole.  For example, there is significant scatter (beyond measurement errors) in the predicted opacities versus UV spectral slope in the starburst sample used to derive the Calzetti law \\citep{calzetti94}.  Therefore, cB58 may simply lie to one end of the natural dispersion already observed in low redshift starbursts.  However, the fact that the phenomenon of the UV overpredicting $L_{IR}$ has been observed in a large sample of young LBGs \\citep{reddy06} suggests that cB58 may be typical of these young sytems, and that a systematic bias in UV-derived SFR derivations may exist for the youngest LBGs.  SED fits to rest-frame UV and optical photometry of LBGs suggest that at least $\\sim 30$\\% have starburst ages less than 100 Myr \\citep{shapley01,papovich01} and, at least for LBGs with $L_{UV}\\sim L^*$, there is no correlation between age and observed UV luminosity \\citep{shapley01}.  Therefore, if the SFR in all of these young LBGs is overestimated by a factor of $\\sim 4-5$, then the global star formation rate density of all LBGs will be overestimated by a factor of $\\sim 2$.  $Spitzer$ observations of recently discovered lensed LBGs \\citep{smail07,coppin07,allam07} will allow us to better determine if these discrepancies persist in the rest of the young LBG population."
    },
    "0808/0808.2471_arXiv.txt": {
        "abstract": "In an earlier paper we quantified the mean merger rate of dark matter haloes as a function of redshift $z$, descendant halo mass $M_0$, and progenitor halo mass ratio $\\xi$ using the Millennium simulation of the $\\Lambda$CDM cosmology. Here we broaden that study and investigate the dependence of the merger rate of haloes on their surrounding environment. A number of local mass overdensity variables, both including and excluding the halo mass itself, are tested as measures of a halo's environment. The simple functional dependence on $z$, $M_0$, and $\\xi$ of the merger rate found in our earlier work is largely preserved in different environments, but we find that the overall amplitude of the merger rate has a strong positive correlation with the environmental densities. For galaxy-mass haloes, we find mergers to occur $\\sim 2.5$ times more frequently in the densest regions than in voids at both $z=0$ and higher redshifts. Higher-mass haloes show similar trends. We present a fitting form for this environmental dependence that is a function of both mass and local density and is valid out to $z=2$. The amplitude of the progenitor (or conditional) mass function shows a similar correlation with local overdensity, suggesting that the extended Press-Schechter model for halo growth needs to be modified to incorporate environmental effects. ",
        "introduction": "\\label{introduction}  In studies of cosmological structure formation, the mass of a dark matter halo is a key variable upon which many properties of galaxies and their host haloes depend. For instance, dark matter haloes of lower mass are expected to form earlier on average than more massive haloes in hierarchical cosmological models such as $\\Lambda$CDM. In semi-analytical modelling of galaxy formation (see \\citealt{BaughReview} for a review), properties such as the formation redshift, halo occupation number, galaxy colour and morphology, and stellar vs AGN feedback processes are all assumed to depend on the mass of the halo (sometimes better characterised by the halo circular velocity).  In addition to the halo mass, however, recent work based on numerical simulations has shown that a halo's local environment also affects various aspects of halo formation. For instance, at a {\\it fixed} mass, older haloes are found to cluster more strongly than more recently formed haloes \\citep{Gottlober01, ShethTormen04, Gao2005, Harker06, Wechsler06, JingSutoMo07, WangMoJing07, GaoWhite07, Maulbetsch07}. Other halo properties such as concentration, spin, shape, and substructure mass fraction have also been found to vary with halo environment (e.g., \\citealt{Avila05, Wechsler06, JingSutoMo07, GaoWhite07, Bett07}).  In contrast, no such environmental dependence is predicted in the extended Press-Schechter (EPS) and excursion set models \\citep{PS74, BondEPS, LC93} that are widely used for making theoretical predictions of galaxy statistics and for Monte Carlo constructions of merger trees. The lack of environmental correlation arises from the Markovian nature of the random walks in the excursion set model. This limitation is not built into the model per se, but is an assumption stemming from the use of a tophat Fourier-space window function. When a Gaussian window function is used, for instance, \\citet{ZentnerEPS} finds an environmental dependence in the halo formation redshift, but the dependence is {\\it opposite} to that seen in the numerical simulations cited above. Other attempts at incorporating environmental effects into the excursion set model thus far have not been able to reproduce the correlations in simulations (e.g., \\citealt{Sandvik07, DesJacques07}).  In this paper, we focus on the environmental dependence of the merger rate of dark matter haloes, a topic that has not been studied in detail. The merger rate is an important quantity for understanding and interpreting observational data on galaxy formation, growth, and feedback processes. While the mergers of galaxies and the mergers of dark matter haloes are not identical processes, the two processes are closely related, and quantifying the latter is the first key step in understanding the former. There have been few theoretical studies of merger rates (e.g., \\citealt{Gottlober01, FM08, Stewart08}) probably because mergers are two- (or more-) body processes, and a large ensemble of descendent haloes {\\it and} their progenitor haloes must be identified from merger trees before the rate can be reliably calculated. In comparison, studies of halo properties such as the mass function, density and velocity profiles, concentration, triaxiality, spin, and substructure distribution require only the particle information from a single simulation output.  This paper is an extension of our earlier study \\cite{FM08} (henceforth FM08). There we quantified the global mean merger rates of haloes in the Millennium simulation \\citep{Springel05} over a wide range of descendant halo mass ($10^{12} \\la M_0 \\la 10^{15} M_\\odot$), progenitor mass ratio ($10^{-3} \\la \\xi \\le 1$), and redshift ($0 \\le z \\la 6$). We found that when expressed in units of the mean number of mergers {\\it per halo} per unit redshift, the merger rate has a very simple dependence on $M_0$, $\\xi$, and $z$: the rate depends very weakly on halo mass ($\\propto M_0^{0.08}$) and redshift, and scales as a power law in the progenitor mass ratio ($\\propto \\xi^{-2.01}$) for minor mergers ($\\xi \\la 0.1$), with a mild upturn for major mergers. These simple trends allowed us to propose a universal fitting form for the mean merger rate that is accurate to 10-20\\%.  Here we go beyond the global merger rate and use the rich halo statistics in the Millennium database to quantify the merger rate as a function of halo environment, in addition to descendant mass, progenitor mass ratio, and redshift. We also investigate the environmental dependence of the progenitor (or conditional) mass function. This quantity is closely related to the merger rate and is also the most important ingredient in the EPS and excursion set models for constructing Monte Carlo merger trees.  Several recent environmental studies have used halo clustering, quantified by the halo bias, as a measure of environment (e.g. \\citealt{Gottlober02, ShethTormen04,Gao2005, Harker06, JingSutoMo07, Wechsler06, GaoWhite07}). While halo bias is a powerful statistical quantity, we choose a simpler and more intuitive local environment measure and use the local mass density centred at each halo. The earlier studies that have used local overdensities as measures of halo environment have used a variety of definitions, e.g., the mean density within a sphere of some radius (ranging from 4 to $10\\, h^{-1}$ Mpc) or within a spherical shell (e.g. between 2 and $5\\, h^{-1}$ Mpc) \\citep{LemsonKauffmann, Harker06, WangMoJing07, Maulbetsch07, Hahn08}. In this paper we compare different definitions of the local overdensity, both including and excluding the mass of the central halo itself.  In this paper we also provide an in-depth investigation of the effects of halo fragmentation on the merger rate and its environmental dependence. In FM08, we discussed how fragmentation is a generic feature of all merger trees and compared the {\\it stitching} method with the conventional {\\it snipping} method for handling these events. We will show here that fragmentation occurs more frequently in dense regions than in voids; understanding the effects of fragmentation on merger rates is therefore essential for obtaining robust results in dense environments. There are three general types of approaches to handling fragmentations: do nothing ({\\it snipping}), {\\it stitching} together fragmented haloes, or {\\it splitting} up the common progenitor of the fragmented haloes. We will compare five algorithms for handling fragmentations based on these three approaches and show that except for one algorithm, all the algorithms give similar merger rates to within 20\\%.  This paper is organised as follows. In \\S~\\ref{Definitions} we briefly review how haloes and merger trees are constructed from the particle data in the Millennium simulation. Statistics detailing the distribution of halo mass at different redshifts are summarised in Table~\\ref{table:MassBins}. In \\S~\\ref{MeasuringEnvironment} we compare four local density measures and their distributions in relation to halo mass. Three of the measures use the dark matter mass in a sphere centred at a given halo, either including or excluding the mass of the central halo. The fourth measure is motivated by observables such as luminosity-weighted galaxy counts and uses only the masses of the haloes within a sphere. \\S~\\ref{MainResults} contains the main results of this paper, where we quantify how the merger rate is amplified in denser regions and suppressed in voids for redshifts $z=0$ to 2 over three decades of halo mass ($10^{12} - 5\\times 10^{15} M_\\odot$). A simple power-law fitting function for this environmental dependence is proposed, which can be used in combination with the fit for the global rate presented in FM08. We also show that the progenitor (or conditional) mass function has a similar environmental trend as the merger rate. Even though this is expected given that the two quantities are closely related, this result demonstrates directly that the excursion set model is incomplete. In \\S~5, we present statistics of halo fragmentations, compare five algorithms for handling these events, and illustrate the robustness of the results reported in \\S~4. The Appendix provides a discussion of the self-similarity of the merger rate and its environmental dependence in the context of the choice of mass and environment variables used in the fitting formula. ",
        "conclusions": "\\label{Conclusions}  We have used the dark matter haloes and merger trees constructed from the Millennium simulation to quantify the dependence of halo merger rates on halo environment from redshift $z=0$ to 2. A number of local mass density parameters centred at the haloes, both including and excluding the central halo mass itself, are tested as measures of environment. We have found that $\\dsfof$ defined in equation~(\\ref{delta3}) is a robust measure of the surrounding environment outside of a halo's virial radius. It cleanly subtracts out the contributions to the local density from the central halo and thereby breaks the degeneracy between halo mass and environment for high mass haloes (see Figs.~\\ref{fig:DeltaComparison} and \\ref{fig:DeltaDistribution}).  We have found strong and positive correlations in both the halo merger rate and the progenitor mass function with environmental densities. Figs.~\\ref{fig:Bn}-\\ref{fig:CMF} present our main results, where haloes in the densest regions are seen to experience 2 to 2.5 times higher merger rates than haloes in the voids. Such a density dependence can be approximated analytically by multiplying our earlier fitting formula FM08 for the global merger rates (eq.~\\ref{eqn:Bnfit}) by an additional $\\delta$-dependent factor given by equation~(\\ref{eqn:FIT}). This factor is a simple power-law in both the environmental density and halo mass, and it is redshift-independent. The mass dependence is quite weak, indicating that haloes with different masses but similar values of $1+\\dsfof$ experience similar merger rate amplifications. This is intriguing in light of the fact, discussed in Section~\\ref{Detangling}, that these haloes actually reside in different environments.  The strong correlations of the halo merger rate and progenitor mass function with environment discussed in this paper have important implications for the analytic Press-Schechter \\citep{PS74} and excursion set models \\citep{BondEPS, LC93}. In this popular formalism, halo growth is modelled by the random walk trajectories of dark matter density perturbations smoothed at decreasing scales. Haloes are identified at scales at which these trajectories first cross some critical density threshold, and the Markovian nature of the model allows one to compute the distribution of these first crossings. This distribution is then mapped onto the number-weighted conditional mass function $\\phi(M,z|M_0,z_0)$ discussed in Section~\\ref{CMF} and plays an important role in the Monte Carlo construction of mock merger tree catalogues (see \\citealt{ZFM08} and references therein).  It is generally assumed that the conditional mass function is independent of environment as the excursion set model is Markovian. The Markovian nature of the random walks, however, is not a prediction but rather an assumption resulting from the use of the $k$-space tophat window function to smooth the density perturbations. There have been recent attempts to weaken this assumption or to introduce environmental dependence into other parts of the model \\citep{ZentnerEPS, Sandvik07, DesJacques07}, but these modifications thus far have not been able to reproduce the basic statistical correlation between halo clustering and formation time found in simulation studies: older haloes are more clustered \\citep{Gottlober01, ShethTormen04, Gao2005, Harker06, Wechsler06, JingSutoMo07, WangMoJing07, GaoWhite07, Maulbetsch07}.  How do our environmental results for the merger rates tie in with these simulation and EPS studies? We have shown that the amplification of halo merger rates in denser regions persists at all redshifts (up to at least $z=2$). If mergers were the dominant channel for halo growth, our results would imply that for haloes of a {\\it fixed} mass today, those in denser regions should have formed {\\it more recently} than those in void regions. Interestingly, this is exactly opposite to the trend reported in many recent studies that have found older (i.e. earlier forming) haloes to be more clustered than younger haloes. As we will discuss in the next paper (Fakhouri \\& Ma 2008c), these two results are in fact not in conflict once the other important channel for halo mass growth -- the ``diffuse'' accretion of non-halo material (either unresolved or stripped) -- is taken into account. We will quantify the environmental dependence of this component and show that, when combined with the merger rate results presented in this paper, we recover the formation redshift dependence reported in prior simulation studies."
    },
    "0808/0808.0704_arXiv.txt": {
        "abstract": "\\vspace*{5mm} We discuss the interpretation of the annual modulation signal seen in the DAMA experiment in terms of spin-independent elastic WIMP scattering. Taking into account channeling in the crystal as well as the spectral signature of the modulation signal we find that the low-mass WIMP region consistent with DAMA data is confined to WIMP masses close to $m_\\chi \\simeq 12$~GeV, in disagreement with the constraints from CDMS and XENON. We conclude that even if channeling is taken into account this interpretation of the DAMA modulation signal is disfavoured. There are no overlap regions in the parameter space at 90\\%~CL and a consistency test gives the probability of $1.2\\times 10^{-5}$. We study the robustness of this result with respect to variations of the WIMP velocity distribution in our galaxy, by changing various parameters of the distribution function, and by using the results of a realistic $N$-body dark matter simulation. We find that only by making rather extreme assumptions regarding halo properties can we obtain agreement between DAMA and CDMS/XENON. ",
        "introduction": "The DAMA collaboration has collected an impressive amount of data in their search for the scattering of weakly interacting dark matter particles (WIMPs) off Sodium Iodine. The combined data from DAMA/NaI (7 annual cycles) and DAMA/LIBRA (4 annual cycles) amounts to a total exposure of 0.82~ton~yr~\\cite{Bernabei:2008yi}, in a field where exposure is measured in units of kg~days. DAMA/LIBRA has now provided further evidence for an annual modulation of the event rate in the energy range between 2 and 6~keVee, the claimed statistical confidence of the positive signal being $8.2\\sigma$~\\cite{Bernabei:2008yi}. The phase of the observed modulation (with maximum on day $144\\pm8$) is in striking agreement with the expectation for a modulation in a WIMP scattering signal due to the rotation of the Earth around the Sun, (expected maximum day 152, June 2nd), see e.g.,~\\cite{Jungman:1995df} for a review. An interpretation of this effect in terms of spin-independent interactions of conventional WIMPs with masses $m_\\chi \\gtrsim 50$~GeV is in direct conflict with the constraints from several experiments looking for direct WIMP detection, most notably with the data from CDMS~\\cite{Ahmed:2008eu} and XENON10~\\cite{Angle:2007uj}, which exclude the WIMP cross section consistent with the DAMA modulation for $m_\\chi\\sim 50$~GeV by many orders of magnitude.  In light of this, several alternative explanations of the DAMA annual modulation have been proposed, for example spin-dependent interactions~\\cite{Ullio:2000bv, Savage:2004fn}, light WIMPs with $\\lesssim 10$~GeV masses~\\cite{Bottino:2003iu, Bottino:2003cz, Gondolo:2005hh}, keV scale axion-like dark matter~\\cite{Bernabei:2005ca} (see however, \\cite{Pospelov:2008jk, Gondolo:2008dd}), dark matter interacting only with electrons~\\cite{Bernabei:2007gr}, inelastic WIMP scattering~\\cite{TuckerSmith:2001hy, Chang:2008gd} and mirror dark matter~\\cite{Foot:2008nw}.  In this work we reconsider the possibility of spin-independent elastic scattering of light WIMPs with $\\lesssim 10$~GeV masses~\\cite{Bottino:2003iu, Bottino:2003cz, Gondolo:2005hh}, see~\\cite{Bottino:2007qg, Bottino:2008mf, Petriello:2008jj, Feng:2008dz} for recent studies. The original idea is that light dark matter scattering on the relatively light Sodium nuclei in DAMA could deposit enough energy in the detector to give a signal, whereas the scattering of light halo particles off heavier nuclei, such as for example Ge in CDMS or Xe in XENON would lead to energy depositions below the threshold of those detectors. Recently the importance of the so-called channeling effect~\\cite{Drobyshevski:2007zj} in the crystal structure of the experiment has also been emphasized~\\cite{Petriello:2008jj, Bottino:2008mf}. Specific models for WIMPs with $m_\\chi \\sim 10$~GeV have been studied for example in~\\cite{Bottino:2002ry, Bottino:2007qg, Barger:2005hb, Gunion:2005rw}. Here we do not discuss theoretical implications but focus on the phenomenology of direct detection experiments in a model-independent way by assuming that such light WIMPs can provide the correct relic abundance while any direct collider constraints can be evaded.  In this region of WIMP masses several experiments~\\cite{Altmann:2001ax, Akerib:2003px, Lin:2007ka, Aalseth:2008rx} exclude WIMP--nucleon scattering cross sections in the range $\\sigma_p \\gtrsim 10^{-40}$~cm$^2$. As we will see in the next pages, once we have included channeling as well as the spectral shape of the DAMA modulation signal, the allowed region of our interest is obtained at much small cross sections, around $\\sigma_p \\sim 10^{-41}$~cm$^2$ and $m_\\chi \\sim 10$~GeV. In this region the most relevant constraints come from XENON~\\cite{Angle:2007uj}, the 2008 Germanium data from CDMS~\\cite{Ahmed:2008eu}, and the 2005 CDMS data on Silicon~\\cite{Akerib:2005kh}. Indeed, as we will discuss, the spectral shape of the DAMA annual modulation restricts $m_\\chi$ and $\\sigma_p$ to a region excluded by these experiments.  In our study we elaborate on this result and discuss how robust it is with respect to different assumptions about the dark matter halo of our galaxy. The impact of non-standard halo properties on dark matter direct detection experiments has been discussed by many authors, see for example~\\cite{Belli:2002yt, Fornengo:2003fm, Green:2002ht, Vergados:2007nc}.  At a qualitative level, one would expect that smaller velocity dispersions or truncated velocity distributions would seem to favour the dark matter interpretation of the DAMA signal, as they could lead to more events above the low energy threshold of DAMA but below that of other experiments.  Furthermore, anisotropies in the velocity dispersion could amplify annual modulation signals.  The outline of our work is as follows. In Sec.~\\ref{sec:analysis} we briefly summarise the phenomenology of elastic WIMP scattering in direct detection experiments and give some technical details on our analysis of DAMA, CDMS and XENON data. The results for a standard dark matter halo are presented in Sec.~\\ref{sec:std-halo}.  In Sec.~\\ref{sec:nonstd-halo} we consider deviations from the standard assumptions made about the WIMP velocity distribution: we use results from the Via Lactea $N$-body dark matter simulation~\\cite{Diemand:2006ik}, we vary several parameters of the Maxwellian distribution and consider asymmetric velocity profiles. Sec.~\\ref{sec:conclusions} contains our conclusions.  In Appendix~\\ref{sec:dama-mod} we comment on the DAMA fit using the annual modulation energy spectrum, and in Appendix~\\ref{app:comparison} we briefly compare our results to the ones from other authors. ",
        "conclusions": "\\label{sec:conclusions}  Prompted by recent results from DAMA/LIBRA which establish the annual modulation of their event rate at the 8.2$\\sigma$ level, we have studied the interpretation of this signal in terms of spin-independent elastic WIMP scattering. We have shown that the energy spectrum of the modulation signal strongly restricts the region of WIMP masses below 10~GeV, confining WIMP masses consistent with the DAMA data close to $m_\\chi \\simeq 12$~GeV. This region is excluded by the limits from CDMS and XENON, and therefore we conclude that even if channeling is taken into account this interpretation of the DAMA modulation signal is disfavoured. Applying a stringent test to evaluate the consistency of DAMA with null-result experiments we find consistency only with a formal probability of $10^{-5}$.  We have studied how robust this result is with respect to variations of the WIMP velocity distribution in our galaxy by changing various parameters of the distribution function. We find that decreasing the dispersion of the distribution can somewhat reduce the tension in the fit. Adopting in addition an asymmetric WIMP velocity profile with a larger dispersion in the radial direction than tangential improves the fit considerably.  We conclude that in principle it is possible to reconsile DAMA in the considered framework, at the price of rather exotic properties of the DM halo. The question remains whether such halo properties can be realistic at all. We have checked that a WIMP velocity distribution based on the Via Lactea $N$-body dark matter simulation does not improve the fit considerably with respect to the standard Maxwellian halo model.  Finally we mention that the negative conclusion on the compatibility of DAMA with CDMS and XENON relies crucially on the energy threshold of the latter two. In particular, a shift in the nuclear recoil energy scale in these experiments may change the conclusion. Indeed, the new measurements of the $\\mathcal{L}_\\mathrm{eff}$ parameter in XENON~\\cite{Sorensen:2008ec} (which has not been implemented in the first arXiv version of this work) made the disagreement between DAMA and XENON somewhat less sever."
    },
    "0808/0808.3707_arXiv.txt": {
        "abstract": "We run adiabatic $N$-body/hydrodynamical simulations of isolated self-gravitating gas clouds to test whether conformal gravity, an alternative theory to General Relativity, is able to explain the properties of X-ray galaxy clusters without resorting to dark matter. We show that the gas clouds rapidly reach equilibrium with a density profile which is well fit by a $\\beta$-model whose normalization and slope are in approximate agreement with observations. However, conformal gravity fails to yield the observed thermal properties of the gas cloud: (i) the mean temperature is at least an order of magnitude larger than observed; (ii) the temperature profiles increase with the square of the distance from the cluster center, in clear disagreement with real X-ray clusters. These results depend on a gravitational potential whose parameters reproduce the velocity rotation curves of spiral galaxies. However, this parametrization stands on an arbitrarily chosen conformal factor. It remains to be seen whether a different conformal factor, specified by a spontaneous breaking of the conformal symmetry, can reconcile this theory with observations. ",
        "introduction": "\\label{sec:intro}  On the scale of individual galaxies and larger scales, General Relativity requires large amounts of dark matter to describe the dynamics of cosmic structure. Moreover, the late-time acceleration of the Hubble expansion implies the existence of a cosmological constant, a special case of a dark energy fluid which is suggested by more sophisticated models (see \\citealt{copeland06}, for a review). In principle, we can avoid the dark matter and dark energy solutions to the puzzles posed by the astrophysical data by adopting an alternative theory of gravity, which reduces to General Relativity in the appropriate limit. Independently of the dark matter and dark energy problems, a modification of General Relativity is also highly desirable if we ultimately wish to unify gravity with the other fundamental interactions.  Possible modifications of the Einstein-Hilbert action proposed in the literature are, among others, (i) the introduction of additional scalar and/or vector fields (e.g., \\citealt{fujii03, beken06}); (ii) the assumption of arbitrary functions $f(R)$ of the Ricci scalar $R$ (e.g., \\citealt{capozz-franc07, nojiri08}); (iii) the introduction of additional dimensions to the four dimensions of the General Relativity spacetime manifold (e.g., \\citealt{maart04}).  A different approach was suggested by \\citet{mannh90} who revived  Weyl's theory \\citep{weyl18, weyl19, weyl20} as a possible candidate to solve the dark matter and dark energy problems. When the geometry is kept Riemannian, with a null covariant derivative of the metric tensor, we can obtain a milder version of Weyl's gravity, known as conformal gravity. In this theory, we impose a local conformal invariance on the gravitational field action in the curved four-dimensional spacetime. The Einstein-Hilbert Lagrangian density for the gravitational field is chosen on the requirement that the theory of gravity is a second-order derivative theory. In conformal gravity, the Lagrangian density is chosen on the principle of local conformal symmetry which is uniquely satisfied by the action $I_W = -\\alpha\\int d^4x\\sqrt{-g} C_{\\mu\\nu\\kappa\\lambda}C^{\\mu\\nu\\kappa\\lambda}$, where $C_{\\mu\\nu\\kappa\\lambda}$ is the Weyl tensor, $\\alpha$ is a coupling constant, and $g$ is the determinant of the metric tensor $g_{\\mu\\nu}$. Conformal symmetry is garanteed by the invariance of the Weyl tensor to local conformal transformations $g_{\\mu\\nu}(x) \\to \\Omega^2(x)g_{\\mu\\nu}(x)$, where $\\Omega^2(x)$ is an arbitrary conformal factor that can be specified by a spontaneous breaking of the conformal invariance \\citep{edery06, mannh08b}. The theory of gravity implied by the action $I_W$ is a fourth-order derivative theory. The vacuum exterior solution for a static and spherically symmetric spacetime contains the Schwarzschild solution \\citep{MK89}. The weak-field limit is consistent with the Solar System observational data \\citep{mannh07}, unlike claimed in previous investigations \\citep{barabash99, flanagan06, barabash07}.  \\cite{mannh93, mannh97} shows that conformal gravity can reproduce the rotation curves of disc galaxies without dark matter. Moreover, conformal gravity can explain the current accelerated expansion of the universe without resorting to a fine-tuned cosmological constant or to the existence of dark energy \\citep{mannh01, mannh03, varieschi08}. Unlike the standard theory, where the universe starts accelerating at redshift $z\\la 1$, conformal gravity predicts that the universe is accelerating at all times. Therefore, observational data probing the Hubble plot at very high redshift (e.g., \\citealt{navia08, wei08}) can be a decisive test.  More recently, conformal theory has been proposed as a valid candidate for building a theory of quantum gravity \\citep{mannh08}; in fact, theories based on fourth-order derivative equations of motion have had the long-lasting problem of suffering from the presence of ghosts. \\citet{bender08} have recently shown that this erroneous belief is the result of considering the canonical conjugates ${\\bf p}$ of the generic dynamical variables ${\\bf q}$, when the ${\\bf q}$'s are real, to be Hermitian operators; however, this assumption is incorrect, and when the non-Hermiticity property of the Hamiltonian of these higher-order quantum field theories is correctly taken into account, the states with negative norm disappear.  From an astrophysical perspective, however, the form of conformal gravity that has been proposed in the literature currently has two main shortcomings: the abundance of light elements, and the gravitational lensing phenomenology.  Conformal gravity nicely avoids the requirement of an initial Big Bang singularity, but it still predicts an early universe sufficiently dense and hot to ignite the light element nucleosynthesis. However, conformal gravity predicts a too slow initial expansion rate. This rate favours the destruction of most of the deuterium produced in the early universe \\citep{knox93, elizondo94} and poses conformal gravity in serious difficulties compared to the standard Big Bang nucleosynthesis. Conformal gravity necessarily requires astrophysical processes for the production of the deuterium currently observed, for example neutron radiative capture on protons in the atmospheres of active stars \\citep{mullan99} or gamma-ray bursts \\citep{inoue03}. However, the processes investigated to date do not seem to be efficient enough. For example, significant production of deuterium in the Galaxy seems to be ruled out \\citep{proda03}.  The second open issue is the conformal gravity prediction of gravitational lensing. Early investigations of gravitational lensing in conformal gravity \\citep{walker94, edery98, edery99p, edery99} show that, in the weak-field limit, the deflection angle due to a point mass $M$ is $\\Delta\\alpha=4GM/c^2r - \\gamma r$, where $r$ is the radius of the photon closest approach; $\\Delta\\alpha$ contains the additional term $\\gamma r$ when compared to the General Relativity result. The constant $\\gamma$ has to be positive to fit the galaxy rotation curves, thus implying a repulsive effect in gravitational lensing. It was later realized \\citep{edery01} that the geodesics of photons are independent of the conformal factor $\\Omega^2(x)$ and one can choose an appropriate conformal factor and a radial coordinate transformation to yield attractive geodesics for both massive and massless particles. However, in the strong-field limit, the light deflection might still be divergent or even impossible \\citep{pireaux04a, pireaux04b}.  Until these open questions are completely settled, conformal gravity cannot yet be considered ruled out by observations. Moreover, conformal invariance plays a crucial role in elementary particle physics and a viable theory of gravity that includes this property can at least be suggestive of a relevant route towards the unification of the fundamental interactions.  From the astrophysical point of view, conformal gravity can be further tested by investigating the formation of cosmic structure. To date, nobody has yet explored how the large-scale structure forms in conformal gravity. If structures form by gravitational instability, as in the standard theory, the development of a cosmological perturbation theory, which is still lacking, becomes inevitable. This theory would enable the comparison of conformal gravity with the spectrum of the Cosmic Microwave Background anisotropies and would provide initial conditions for the simulation of the structure evolution into the non-linear regime.  Before building up such a theory, however, it is useful to check whether conformal gravity is able to reproduce the equilibrium configuration of cosmic structures, other than galaxies, without dark matter. \\cite{horne06} has already shown that, if we interpret the observational data of the intracluster medium of X-ray clusters by assuming hydrostatic equilibrium, conformal gravity requires a factor of ten less baryonic mass than inferred from the X-ray surface brightness measures. Here, we extend this analysis by performing hydrodynamical simulations of self-gravitating gas clouds. We modify a standard Tree+SPH code to perform numerical simulations of self-gravitating systems in conformal gravity. The numerical tool we create is extremely relevant if we eventually wish to investigate the formation of the large-scale structure, because we will massively need to resort to numerical integrations when the evolution of the density perturbations reaches the non-linear regime.  In Section \\ref{sec:CG}, we review the basic steps leading to the gravitational potential of a static point source in conformal gravity. In Section \\ref{sec:sphere}, inspired by the analysis of \\citet{horne06}, we compute the gravitational potential energy of a spherical system, and in Section \\ref{sec:xray} we use this result to compute the expected mean temperature and the temperature profile of the intracluster medium. In Section \\ref{sec:numerics}, we derive the same results with  hydrodynamical simulations of self-gravitating gas clouds. ",
        "conclusions": "Conformal gravity can explain the rotation curves of disk galaxies and the current accelerated expansion of the universe without resorting to dark matter and dark energy. We have modified a Tree+SPH code to run hydrodynamical simulations of isolated X-ray galaxy clusters to show that conformal gravity does not share the same success on the scales of clusters.  These simulations confirm our simple analytic estimates that show that gas clouds with mass $\\sim 10^{12}-10^{13}$~M$_\\odot$, which are typical values of the total mass of the hot gas present in real clusters, remain confined with an equilibrium mean temperature $\\sim 10-100$~keV, ten times larger than the observed temperatures; more dramatically, because of the presence of a linear term in the gravitational potential, at large clustrocentric radius $r$, the gas temperature increases with $r^2$, rather than decreasing as in real systems.  Our analysis totally neglects radiative cooling and gas heating from astrophysical sources, as supernovae or active galactic nuclei. The interplay between these processes can in principle provide a way to reconcile conformal gravity with observations. It is however unclear if and how much these processes should be fine-tuned to provide X-ray clusters in agreement with observations.  In addition to this topic, we can see two more open issues whose solution might also reconcile conformal gravity with observations:  \\begin{enumerate} \\item In conformal gravity, all the matter in the universe is expected to affect the local dynamics. The net effect is to contribute a constant inward acceleration $-GM_0/R_0$ in addition to the gravitational acceleration generated by local sources. Because we included this constant acceleration in our simulations, we could neglect the rest of the universe and impose vacuum boundary conditions. It might be possible that assimilating the gravitational influence of the nearby matter surrounding the X-ray cluster in the constant ``universe'' acceleration $-GM_0/R_0$ is inappropriate: in fact, nearby external matter might decrease the gravitational attraction of the interior matter and hopefully reduce the thermal energy of the gas. To appropriately investigate this effect, we should simulate the dynamics of large-scale structure within a full cosmological context. However, this task is not trivial just because the gravitational field is highly non-local. This investigation would also benefit from the implementation into the numerical simulation of the yet unavailable theory of structure formation. \\item The gravitationl potential we implemented in our simulations derives from a metric where the conformal invariance is broken by an arbitrary choice of the conformal factor $\\Omega^2(x)$. It is unclear whether this choice provides a coordinate system whose physics describes the real world or it is an artifact of the reference frame. It also remains to be seen whether a spontaneous breaking of the conformal invariance, in theories where matter and gravity are conformally coupled \\citep{edery06}, can provide a metric, and thus a gravitational potential, where the observed thermal properties of the intracluster medium can be reproduced. \\end{enumerate}   In Section \\ref{sec:intro} we mentioned that the nucleosynthesis of light elements and the phenomenology of gravitational lensing are two open issues that need to be solved before accepting conformal gravity as a viable alternative theory of gravity and cosmology. Here, we have shown that the thermodynamics of X-ray clusters poses a third challenge to this theory."
    },
    "0808/0808.3641_arXiv.txt": {
        "abstract": "{ We discuss the possibility that high-frequency QPOs in neutron-star binary systems may result from forced resonant oscillations of matter in the innermost parts of the accretion disc, excited by gravitational perturbations coming from asymmetries of the neutron star or from the companion star. We find that neutron-star asymmetries could, in principle, be effective for inducing both radial and vertical oscillations of relevant amplitude while the binary companion might possibly produce significant radial oscillations but not vertical ones. Misaligned neutron-star quadrupole moments of a size advocated elsewhere for explaining limiting neutron star periods could be large enough also for the present purpose.} ",
        "introduction": "%  Quasi-periodic oscillations of \\mbox{X-ray} brightness (QPOs) have been observed in a number of accreting binary systems containing compact objects, both with neutron stars \\citep[see][ for a review]{Kli:2000:ARASTRA:,Bar-Oli-Mil:2005:MONNR:} and with black holes \\citep{Rem-McCli:2006:ARASTRA:}. They can have low frequencies ($\\mathrm{Hz}$) or high frequencies ($\\mathrm{kHz}$). The observed $\\mathrm{kHz}$ frequencies are comparable with the Keplerian and epicyclic frequencies in the inner parts of the accretion disc \\citep{Tor:2005:ASTRN:} and the question arises of whether they might be associated with forced resonant oscillations of the inner disc material. In order to initiate these, some perturbation mechanism would be required. Here, we focus on neutron-star systems and investigate the possibility that gravitational perturbations caused either by the binary companion or by asymmetries of the neutron star might provide this mechanism. (The case of the binary companion could be relevant also for black hole systems.) We note that perturbations coming from the neutron star can only be relevant for the innermost parts of the disc (because of the rapid fall-off of the force with distance) and hence are mainly associated with the picture for $\\mathrm{kHz}$ QPO frequencies rather than with that for the lower frequencies. It is necessary that the mechanism should be a resonant one since otherwise the response produced would certainly be much too small to be of interest.  The two types of perturbation (from the neutron star and from the binary companion) clearly induce different behaviours: the frequency of the varying force arising from the influence on the disc of the binary companion is essentially equal to the disc rotation-frequency, whereas the main frequency of the force caused by asymmetries of the neutron star is equal to the difference between the rotation frequencies of the disc and the neutron star \\citep{Petri:2006:ApSS:}. Resonance occurs at those points where one of the intrinsic oscillation frequencies of the disc matter coincides with the forcing frequency. Following the discussion by \\citet{Lan-Lif:1976:Mech:} for test particle motion (which would need to be modified for fluid elements), the growth in amplitude of the oscillations of the particle within the linear regime of forced resonance is given by \\begin{equation} a(t) = \\frac{f_{\\mathrm{p}}}{2m_0\\omega}\\,t\\ , \\label{resonance} \\end{equation} where $f_{\\mathrm{p}}$ is the amplitude of the variations of the force, $\\omega$ is the frequency and $m_0$ is the mass of the particle. This linear regime ends when the oscillation amplitude $a(t)$ becomes large enough so that non-linear phenomena and/or dissipative processes become relevant. Note that $a(t)$ grows linearly with time in this regime and so can become quite large even when the variations in the perturbing force are small. ",
        "conclusions": "\\label{compar}%  Here, we discuss whether the effect of the neutron-star asymmetries or of the binary companion could be large enough to account for excitation of the QPO phenomenon. In our simplified picture (c.f. Eq.\\,(\\ref{resonance})), the excitation time for the amplitude $a$ of resonant oscillations to grow to a particular value is given by \\begin{equation} t_{\\mathrm{ex}} = \\left(\\frac{a}{R}\\right)\\, \\left(\\frac{\\alpha}{\\pi}\\right)\\, \\left(\\frac{f_{\\mathrm{p}}}{f_0}\\right)^{\\!-1}\\, \\tau_{\\sssm{K}}\\,, \\end{equation} where $f_p$ is the amplitude of the perturbing force acting on unit mass, $f_0 = GM_{\\sssm{A}}/R^{2}$ is the main gravitational force from the central object, $\\tau_{\\sssm{K}} = 2\\pi/\\Omega_{\\sssm{K}}$ is the period of circular Keplerian motion at the location being considered and the epicyclic frequency of the perturbation being excited is $\\omega = \\alpha\\Omega_{\\sssm{K}}$. The dimensionless radial and vertical ``epicyclic functions'' satisfy $\\alpha \\le 1$ everywhere. Note that since $\\tau_{\\sssm{K}} \\sim 10^{-3}\\,\\mathrm{s}$, the amplitude amplification in 1 second is $\\sim\\! 10^{3}$. The ratio $a/R$ needs to grow to $> 10^{-3}$ in order to potentially explain the QPO behaviour and it would need to do that within $\\sim\\!10^{3}\\,\\mathrm{s}$ to account for the QPO phenomena seen in atoll sources.  In the case of a binary companion, for the radial perturbing force one has \\begin{eqnarray} \\frac{f_{\\mathrm{p}}}{f_0} & \\sim & \\left(\\frac{M_{\\sssm{B}}}{M_{\\sssm{A}}}\\right) \\left(\\frac{R}{d}\\right)^{2} \\nonumber \\\\ & \\sim & 10^{-9} \\left( \\frac{M_{\\sssm{B}}}{0.1\\,M_{\\sssm{A}}} \\right) \\left( \\frac{R}{10^{6}\\,\\mathrm{cm}}\\frac{10^{10}\\,\\mathrm{cm}}{d} \\right)^{2}\\, , \\end{eqnarray} where we have normalised to typical parameter values. This could produce $a/R \\sim 10^{-3}$ at times $t_{\\mathrm{ex}}\\lesssim 10^{3}\\,\\mathrm{s}$, but there is a problem for it producing resonances in the innermost parts of the disc because of having $\\omega_{\\sssm{B}} \\approx\\Omega_{\\sssm{K}}$. The corresponding vertical force is smaller by at least four orders of magnitude and is clearly irrelevant. (Note that the vertical force is a tidal force whereas the radial one is a direct gravitational attraction; also, $\\theta_{\\sssm{B}}$ is probably rather close to $\\pi/2$.)  For the neutron star asymmetries, we focus on the case of the misaligned quadrupole moments where $f_{\\mathrm{p}}/f_0$ can be $\\sim 10^{-8}$ in the inner parts of the disc for both radial and vertical oscillations (taking $m_{\\mathrm{quad}} = 10^{-8}\\,M_{\\sssm{A}}$). For this, $a/R$ could reach $10^{-3}$ in $\\lesssim 10^2 \\,\\mathrm{s}$, which encourages further investigation of this scenario.  We conclude that at least one of the types of gravitational perturbation considered in this paper might provide a plausible mechanism for inducing kHz QPO behaviour although many details remain to be worked out (in particular concerning the response of the fluid medium and the production of the luminosity variations). The influence of the binary companion could possibly be effective in providing the perturbations but the more likely possibility is that they might be produced by the neutron-star asymmetries. It is striking that the same magnitude for the misaligned quadrupole moment as advocated elsewhere for explaining limiting neutron star periods, seems also to give a plausible mechanism for inducing QPO behaviour.  \\vspace{\\baselineskip}  \\hrule  \\vspace{\\baselineskip}"
    },
    "0808/0808.0254_arXiv.txt": {
        "abstract": "The unexpected high bump in the UV part of the spectrum found in nearby giant elliptical galaxies, a.k.a. the UV upturn, has been a subject of debate. A remarkable progress has been made lately from the observational side, mainly involving space telescopes. The GALEX UV telescope has been obtaining thousands of giant ellipticals in the nearby universe, while HST is resolving local galaxies into stars and star clusters. An important clue has also been found regarding the origin of hot HB stars, and perhaps of sdB stars. That is, extreme amounts of helium are suspected to be the origin of the extended HB and even to the UV upturn phenomenon. A flurry of studies are pursuing the physics behind it. All this makes me optimistic that the origin of the UV upturn will be revealed in the next few years. I review some of the most notable progress and remaining issues. ",
        "introduction": "A review on the UV upturn phenomenon may usually start with a following or similar definition: ``a bump in the UV spectrum between the Lyman limit and 2500\\AA\\, is found {\\em virtually in all bright spheroidal galaxies}'' (e.g., Yi \\& Yoon 2004). This seems no longer true! While earlier studies based on a small sample of nearby galaxies led us to think so, a much greater sample from the recent GALEX database appears to disprove it. Only a small fraction of elliptical galaxies show a strong UV upturn and it is generally limited to the brightest cluster galaxies (Yi et al. 2005). This review is about the recent development on this seemingly-old topic. I recycle some of the contents in my earlier review given in the first Hot Subdwarf and Related Objects workshop held in Keele, UK (Yi \\& Yoon 2004). For a more traditional review, readers are referred to the articles of Greggio \\& Renzini (1999) and O'Connell (1999). ",
        "conclusions": "The UV satellite GALEX is obtaining a valuable UV spectral evolution data for numerous bright cluster galaxies. The apparent trend in redshift vs $FUV-V$ colour seems consistent with the prediction from the single stellar population models. This is comforting while observers feel obliged to build up their database much more substantially in order to make it statistically robust.  Two new issues are notable. Firstly, binary population synthesis community feels odd to find that the simplistic single-star population models are found to be good enough. The in-principle more advanced binary population models are obliged to reproduce the observed CMDs of simple populations (globular clusters) before attempting to model galaxies. For example, I am very eager to see their models reproduce the ordinary HB first, before explaining the EHB.  Secondly, the enhanced helium hypothesis based on the globular clusters in Milky Way and M87 is a very exciting possibility. The deduced value of the helium abundance seems unphysical to be a global property for the galaxy but may be possible for small systems that are vulnerable to a chemical fluctuation in the proto-galaxy cloud. While a more detailed investigation is called for it may be difficult to be influential to the entire stellar population of a galaxy. For instance, adding all spectral energy distributions of the Milky Way globular clusters would not yield anything close to the spectrum of a UV upturn galaxy. Of course, a metallicity difference may act as an added complication.  The secret will be revealed through time and hard work, perhaps very soon."
    },
    "0808/0808.0581_arXiv.txt": {
        "abstract": "A {\\itshape Chandra} study of pulsar wind nebula around the young energetic pulsar PSR B1509$-$58 is presented.  The high resolution X-ray image with total exposure time of 190 ks reveals a ring like  feature 10'' apart from the pulsar.  This feature is analogous to the inner ring seen in the Crab nebula and thus may correspond to a wind termination shock. The shock radius enables us to constrain the wind magnetization, $\\sigma \\geq 0.01$.  The obtained $\\sigma$ is one order of magnitude larger than that of the Crab nebula.  In the pulsar vicinity, the southern jet appears to extend beyond the wind termination shock, in contrast to the narrow jet of the Crab.  The revealed morphology of the broad jet is coincident with the recently proposed theoretical model in which a magnetic hoop stress diverts and squeezes the post-shock equatorial flow towards the poloidal direction generating a jet. ",
        "introduction": "Recent X-ray observations revealed that part of pulsar wind nebulae (PWNe) have common structures of ``Torus'' and ``Jet'', as typified by the crab pulsar \\citep{2002ApJ...577L..49H,2000ApJ...536L..81W}, the Vela pulsar \\citep{2003ApJ...591.1157P,2001ApJ...556..380H}, and PSR B1509$-$58 \\citep{2002ApJ...569..878G}.  However the physical mechanisms which form the ubiquitous features, especially the jets, are still unclear, because the origin of the jet in the Crab is too compact to discriminate its structure even with the current X-ray observatories.  The theoretical approaches toward the axisymmetric PWNe have been aimed to figure out the Crab nebula.  The torus in the Crab had been successfully explained by one-dimensional MHD simulations \\citep{1984ApJ...283..694K,1984ApJ...283..710K}.  A young rotation powered pulsar is believed to lose its rotating energy via a magnetized particle flow from a pulsar, the so-called pulsar wind.  Since the pulsar wind is flowing almost at the speed of light, it is stalled by a termination shock just around the pulsar.  When the wind passes through the termination shock, the kinetic energy of the wind is converted into internal energy. Essentially, the post-shock flow is decelerated yielding to the conservation law of particle flux $nvr^2 = \\mathrm{const}$, where $n$ is the number density (that is almost constant), $v$ is the flow velocity and $r$ is the radius from the pulsar.  At the same time, the frozen-in condition, $rvB = \\mathrm{const}$, induces the compression of the toroidal magnetic field $B$, in which the internal energy is efficiently converted into synchrotron radiation forming an X-ray torus.  In practice the magnetic pressure prevents the deceleration of the post-shock flow, thus the magnetization of the pulsar wind is closely related to the geometric structure of the PWN.  In this way, several methods to calculate the magnetization parameter $\\sigma$ have been proposed \\citep{1974MNRAS.167....1R,1984ApJ...283..694K,1994ApJ...435..230G}. Moreover, recent two-dimensional numerical studies suggested that the magnetization seems to control the formation of the jet \\citep{2002MNRAS.329L..34L,2004A&A...421.1063D,2006cosp...36..191B}. Therefore it is essential to evaluate the magnetization parameters to understand the formation process of PWNe.  \\begin{figure*} \\begin{center} \\FigureFile(170 mm,120mm){overview_closeup.eps} \\end{center} \\caption{(Left)--- Overview image of the PWN accompanied by PSR B1509$-$58 with 190 ks exposure taken by {\\itshape Chandra}.  The image was smoothed with a Gaussian of $\\sigma=2.0''$.  The white rectangle indicates the FoV of the right panel.  (Right)--- Close-up image of the pulsar vicinity smoothed with a Gaussian of $\\sigma=0.5''$.  The inner panel schematically describes the remarkable structures.} \\label{X-ray-image} \\end{figure*}  So far, the shock front of the pulsar wind has been observed only in the Crab nebula \\citep{2000ApJ...536L..81W} because of the small angular sizes of PWNe.  Nevertheless the {\\itshape Chandra} X-ray observatory with sub-arcsec-scale spatial resolution has the potential to detect the wind termination shocks in other PWNe.  In order to obtain a general understanding of the formation process of PWNe, we investigate the fine structures around the pulsar PSR B1509$-$58 utilizing the {\\itshape Chandra} data.  PSR B1509$-$58 is emitting 150 ms radio, X-ray, and Gamma-ray pulsation \\citep{1982ApJ...256L..45S,1982ApJ...262L..31M,1993ApJ...417..738U}. From the spin parameters, a characteristic age $\\tau_\\mathrm{c}=1700$ yr, a spin-down luminosity $\\dot{E} = 1.8 \\times 10^{37}$ ergs s$^{-1}$, and a surface magnetic field $B_\\mathrm{p}=1.5\\times10^{13}$ G have been obtained, making it one of the youngest, the most energetic, and highest field pulsars known \\citep{1982ApJ...256L..45S,1994ApJ...422L..83K}. The pulsar is accompanied by a bright and large PWN with remarkable jet-like features extending to the southeast and northwest \\citep{2002ApJ...569..878G}.  The pulsar and its PWN appears to be embedded in a 30 arcmin radio shell G320.4$-$01.2 \\citep{1981MNRAS.195...89C}.  An H{\\rm I} observation yielded the distance of $d=5.2\\pm1.4$ kpc from the earth \\citep{1999MNRAS.305..724G}.  Adopting this distance and standard parameters of the ISM and supernova explosions yields an age of 6-20 kyr, which is one order of magnitude larger than the pulsar's characteristic age \\citep{1983ApJ...267..698S}.  Moreover \\citet{2005XrU...604..379Y} detected proper motions of X-ray clumps in RCW 89 which coincide with the north radio shell from the 4.3 yr base-line {\\itshape Chandra} observations.  These results imply that the pulsar and the SNR have the same progenitor in a cavity.  If this is the case the termination shock of the pulsar wind can be easily discriminated by {\\itshape Chandra}.  The paper is structured as follows.  The {\\itshape Chandra} observations and their results are shown in \\S2.  The imaging analysis and spectral analysis are described in \\S3 and \\S4, respectively.  Finally, we discuss the obtained results in \\S5. ",
        "conclusions": "\\subsection{Double tori} As shown in Figure \\ref{radial-profile}, the PWN around PSR B1509$-$58 has a centrally concentrated surface brightness profile, which may be caused by its comparatively small angular size or the contamination from the bright jet structure pointing towards us.  For this reason, it is difficult to compare the surface brightness evolution of PSR B1509$-$58 with that of the Crab nebula in a simple way.  Nevertheless PSR B1509$-$58 possesses an apparently different morphology; the nested double ring structure, inner-torus ($R\\sim30''$) and outer-torus ($30''<R<60''$), lying in the equatorial plane relative to the spin axis.  Spectroscopy, these two tori, calling the inner torus and the outer torus, seems to be completely different.  The inner-torus has a constant photon index of $\\Gamma \\sim 1.6$ which is the same as that just at the pulsar.  In contrast, the photon index in the outer-torus obviously increases with the radius from $\\Gamma=1.6$ to 2.1.  The observed interval of the photon index variation is $\\Delta \\Gamma\\sim0.5$, which implies that the observed spectral evolution is caused by synchrotron cooling \\citep{1962SvA.....6..317K}.  In case of the Crab nebula, the post-shock photon index is $\\Gamma=2.1$ and it does not change within the torus.  According to \\citet{1984ApJ...283..694K}, the evolution of the surface brightness in the torus can be explained by the compression of the toroidal magnetic field, which enhances the synchrotron energy loss.  Thus the bright torus of the Crab simply corresponds to the region in which most of the X-ray emitting particles radiate their energy and burn off there. Consequently the photon index outside of the torus is steeper than $\\Gamma=2.1$, and the surface brightness decays rapidly.   In terms of spectroscopy, the counterpart of the X-ray torus in the Crab is the inner torus rather than the outer torus of PSR B1509$-$58.  While in the outer-torus the X-ray emitting charged particles have burnt off and thus the synchrotron cooling break is passing through the observed energy band (8.0 keV to 0.4 keV).  Moreover, it was reported that the outer-torus has a radio counter part \\citep{2002ApJ...569..878G}.  These results of the spectroscopy imply that the outer-torus could not be simply explained by radiative losses as claimed by \\citet{2002ApJ...569..878G}.  \\subsection{Inner Ring} It is essential to know the radius of the pulsar wind termination shock to understand the mechanism which forms the structure of a PWN.  So far, the termination shock of the pulsar wind has been observed only in the Crab nebula.  In this section we discuss the ring feature discovered near around PSR B1509$-$58 comparing it with the inner ring of the Crab nebula.  As the pulsar wind is believed to be ejected almost at the light speed, the free flowing region should not be radiating as seen in the Crab nebula.  In the case of PSR B1509$-$58 the region within the ring feature is also dark.  Comparing it with the tori which are strongly affected by the Lorenz boost, the surface brightness of the ring feature appears to be uniform, so that we can distinguish the south edge of the ring.  For the Crab nebula \\cite{2000ApJ...536L..81W} pointed out that the surface brightness of the inner ring is more uniform in azimuth, indicating that relativistic beaming is less significant in it than in the torus.  The observed apparent properties can distinguish the ring from the tori and then the hypothesis that the ring structure represents the termination shock can be suggested.  By the way, the ring shows an obvious elliptic shape, implying that the shock front in an ellipsoidal surface or a ring viewed with inclination rather than a sphere. Regarding the issue, recent numerical simulations predict a shock front with an oblate shape for an anisotropic pulsar wind \\citep{2003MNRAS.344L..93K,2004A&A...421.1063D}.  At the termination shock, pressure balance between the ram pressure of the pulsar wind and the post-shock plasma is expected, as seen in the Crab nebula.  Based on the X-ray image \\ref{unsharp-masked-image}, we obtained a radius of the ring structure of $r\\sim 9''$, which corresponds to $7\\times 10^{17}$ cm for the distance of 5.2 kpc to the pulsar.  Assuming that most of the pulsar's rotating energy, $\\dot{E}=1.8\\times 10^{37}$ ergs s$^{-1},$ is transferred to the relativistic pulsar wind at the light cylinder, we find the ram pressure of the pulsar wind at the ring structure to be  \\begin{eqnarray} p_\\mathrm{*} &=& \\frac{\\dot{E}}{4 \\pi r_\\mathrm{TS}^{2}\\phi c}\\\\ &=& 9.6\\times 10^{-11} \\phi^{-1} \\left(\\frac{d}{5.2\\;\\mathrm{kpc}}\\right)^{-2} \\left(\\frac{R}{9''.0}\\right)^{-2} \\quad \\mathrm{dyn}\\;\\mathrm{cm}^{-2}, \\end{eqnarray} where $\\phi$ is the fraction of a sphere covered by the wind.   On the other hand, the inner pressure of the nebula is roughly estimated by assuming equipartition conditions (Reference).  From the spectral fits we obtained an X-ray flux of $1.32\\times 10^{-13}$ ergs s$^{-1}$ cm$^{-2}$ and a photon index of $\\Gamma=1.6$ for the region where the ring feature exists ($R=10''$) and we infer an unabsorbed luminosity density at 1 keV $L_\\mathrm{1\\;keV} = 6.6\\times10^{14}$ ergs s$^{-1}$ Hz$^{-1}$.  For simplification we assumed the emitting volume to be a partial spherical shell of radius $R=10''$, thickness $\\Delta R=3''$ and opening angle of $45^{\\circ}$, corresponding to $V=2.6\\times10^{53}$ cm$^{3}$, which may be an overestimate.  For a filling factor $f$ and a ratio of ion to electron energy $\\mu$, the minimum energy present in the source responsible for the synchrotron emission $W_\\mathrm{total} \\sim 2.9(1+\\mu)^{4/7} f^{3/7} \\times 10^{43}$ ergs and a magnetic field of $B\\sim35(1+\\mu)^{2/5}f^{-2/7}\\;\\mu$G are inferred.  The corresponding pressure of the relativistic gas is then $p_\\mathrm{neb} = W_\\mathrm{total}/3V\\sim 3.7(1+\\mu)^{3/7}f^{-4/7} \\times10^{-11}$ dyn cm$^{-2}$ \\footnote{For a torus geometry with opening angle $15^{\\circ}$ corresponding to $f\\sim 0.5$, the internal pressure of $p_\\mathrm{neb} \\sim 5.6(1+\\mu)^{3/7}\\times10^{-11}$ dyn cm$^{-2}$ can be obtained.}. Note that the obtained inner pressure is a lower limit because the post-shock flow of a weakly magnetized pulsar wind is believed to be in particle dominant conditions rather than in equipartition \\citep{1984ApJ...283..694K}.  Taking into account this fact, $p_\\mathrm{*}$ and $p_\\mathrm{neb}$ are comparable, implying that the ram pressure of the pulsar wind and the internal pressure of the PWN is balanced at the ring feature.  It is therefore inferred that the newly found ring feature is the termination shock of the pulsar wind.  \\subsection{Magnetization of the pulsar wind} Adopting that the discovered ring feature corresponds to the wind termination shock, the magnetization parameter $\\sigma$ can be estimated.  As the shock transforms the kinetic energy into thermal energy, $\\sigma$ can be rewritten as \\citep{1998nspt.conf..457S} \\begin{eqnarray} \\sigma&=&\\frac{B_1^2/4\\pi}{[\\mathrm{kinetic \\; energy\\; density}]}\\\\ &\\sim&\\frac{B_1^2/4\\pi}{\\mathrm{[postshock\\;thermal\\;energy\\;density]}}\\\\ &\\sim&\\frac{B_1^2/4\\pi}{B_\\mathrm{Eq}^2/4\\pi}, \\end{eqnarray} where $B_1$ is the toroidal magnetic field just behind the shock, and $E_\\mathrm{Eq}$ is the equipartition magnetic field after the shock. Combining this equation with the equation of continuity ($r_\\mathrm{TS}^2 c/3 = r_\\mathrm{n}^2 v_\\mathrm{n}$) and the flux conservation law ($r_\\mathrm{TS}cB_1=r_\\mathrm{n}B_\\mathrm{Eq}$) yields \\begin{equation} \\frac{r_\\mathrm{n}}{r_\\mathrm{TS}} \\sim \\frac{1}{3\\sqrt{\\sigma}}. \\end{equation} Substituting a shock radius $r_\\mathrm{TS}=9''$ and a radius of the nebula $r_\\mathrm{n}=30''$, corresponding to the inner-torus, we obtain $\\sigma \\sim 0.01$, which is twice larger than that reported by \\citet{2002ApJ...569..878G} although they regarded the inner-torus in Figure \\ref{X-ray-image} as a wind termination shock.  In case of the Crab nebula a magnetization parameter of $\\sigma=0.003$ based on the expansion velocity was calculated by \\citet{1984ApJ...283..694K},  while \\cite{2002phD...Mori.OsakaU} also reported a larger magnetization parameter of $\\sigma \\sim 0.03$ based on imaging analyses utilizing {\\itshape Chandra}.  \\subsection{South jet} The south jet in the pulsar vicinity is remarkably different from that seen in the Crab.  In case of the Crab nebula, the jet seems to flow out directly from the pulsar, even {\\itshape Chandra} cannot resolve the origin of the jet \\citep{2000ApJ...536L..81W,2002ApJ...577L..49H,2004ApJ...609..186M}.  In contrast, the jet of PSR B1509$-$58 extends over $\\sim10''$ at the origin.  Also, the jet itself is broad and brighter than its torus.  If the newly discovered ring feature is a wind termination shock, the south jet of PSR B1509$-$58 may flow out from the vicinity of the shock front or from the inner-torus, namely the post shock region.  The appearance of the south jet at the origin seems to match the recent theoretical scenario in which the hoop stress in the pulsar wind collimates the poloidal flow \\citep{2002MNRAS.329L..34L}.  This idea was of course introduced for the Crab nebula, however, it is difficult to explain the narrow jet because the magnetic pinch works well only downstream of the shock.  Downstream, the magnetic hoop stress is believed to be able to squeeze the ``mildly relativistic'' flow towards the rotating axis.  The observed features indicate that this model can be applied to PSR B1509$-$58.  \\citet{2004A&A...421.1063D} reported two dimensional relativistic MHD simulations of pulsar jets in various configurations.  Interestingly the observed morphology of the PWN around PSR B1509$-$58 closely resembles the results of the simulations for $\\sigma=0.03$.  They argued that producing a jet by magnetic hoop stress requires a comparatively large magnetization parameter, $\\sigma \\geq 0.01$, which is coincident with the calculated value shown above.  It is therefore argued that the south jet of PSR B1509$-$58 originates in the post-shock equatorial flow."
    },
    "0808/0808.2317_arXiv.txt": {
        "abstract": "{} {We study the recently discovered twisting motion of bright penumbral filaments with the aim of constraining their geometry and the associated magnetic field.} {A large sunspot located $40\\degr$ from disk center was observed at high resolution with the 1-m Swedish Solar Telescope. Inversions of multi-wavelength polarimetric data and speckle reconstructed time series of continuum images were used to determine proper motions, as well as the velocity and magnetic structure in penumbral filaments.} {The continuum movie reveals apparent lateral motions of bright and dark structures inside bright filaments oriented parallel to the limb, confirming recent Hinode results. In these filaments we measure upflows of $\\approx 1.1~\\mathrm{km/s}$ on their limbward side and weak downflows on their centerward side. The magnetic field in them is significantly weaker and more horizontal than in the adjacent dark filaments.} {The data indicate the presence of vigorous convective rolls in filaments with a nearly horizontal magnetic field. These are separated by filaments harbouring stronger, more vertical fields. Because of reduced gas pressure, we see deeper into the latter. When observed near the limb, the disk-centerward side of the horizontal-field filaments appear bright due to the \\textit{hot wall} effect known from faculae. We estimate that the convective rolls transport most of the energy needed to explain the penumbral radiative flux.} ",
        "introduction": "The discovery of twisting motions of penumbral filaments seen in time series of filtergrams by Ichimoto et al. (2007) ranks among the most striking recent discoveries in solar physics. In sunspots observed at a heliocentric angle $\\theta$ of $40-50\\degr$ filaments lying roughly parallel to the limb display a twisting motion directed towards solar disk center. The nature and origin of this apparent motion is still unclear and requires further study. An important step in this direction is to determine how these twisting filaments fit into the complex magnetic structure of the penumbra. This is the main aim of the present letter, along with the confirmation of the Hinode-based results of Ichimoto et al. (2007) with the 1-m Swedish Solar Telescope, which allows higher spatial resolution images to be obtained. ",
        "conclusions": "In Fig.~4 we present a sketch of the geometry of the penumbral fine structure, which is consistent with the results of the present letter, as well as those of Ichimoto et al. (2007), Borrero et al. (2008), and Borrero \\& Solanki (2008) based on Hinode data. The $\\tau=1$ level is relevant for the line wings. In the sketched configuration stronger, more vertical fields (spines) are associated with less dense gas, leading to a lowered $\\tau=1$ level there. Interspersed between them are filaments of nearly horizontal, weaker magnetic field pointing into the page (interspines). Important is that the filaments carrying the Evershed flow have a sizable horizontal field (Borrero \\& Solanki 2008), since the flowing gas is observed to be magnetized (Solanki et al. 1994; Bellot Rubio et al. 2004). It is within these filaments that the transverse motions are seen, which are similar to overturning convection. A nearly horizontal magnetic field means that convection takes place in the form of convective rolls, as first proposed by Danielson (1961). Due to the raised $\\tau=1$ level in these filaments (consistent with the findings of Rimmele 2008) the bright part of the filament, which is contoured in Fig.~2, is the penumbral counterpart of the \\textit{hot wall} found in faculae and pores (Spruit 1976; Keller et al. 2004; Carlsson et al. 2004; Cameron et al. 2007). This geometry explains why only lateral motions directed towards disk center are seen and why the observed upflow velocity is higher than the downflow velocity, as can be seen by considering the directions of the flow arrows relative to the line-of-sight in Fig.~4. From these hot walls radiation flows into the neighbouring spines. The parallax effect, discussed in Sect.~3, makes it obvious that the most inclined magnetic field, which is located above the center of the filament, was detected at the limbward border of the bright filament. The contrast between spine and interspine is expected to depend on the relative temperature contrast, width of the spine and the difference in evacuation between spines and interspines. The values plotted in Figs.~2 and~3 are weighted averages along the LOS (the relatively few measured wavelength points do not allow a more detailed analysis). Thus, because a number of the slanted rays pass through both the inclined and the horizontal fields, the difference between them appears less clear in Figs.~2 and 3 than it may be in reality. Whether our measurements contradict the findings of Jur\\v c\\'ak \\& Bellot Rubio (2008), who report that \\textit{bright penumbral filaments show the more vertical fields and weaker flows}, cannot be judged, in part because of the different viewing geometry. They have investigated the limbward part of the outer penumbra, whereas our investigation addresses filaments oriented parallel to the limb, in the mid penumbra.  The energy transport by the convective rolls can be expressed as $F_{\\rm cr}=\\rho\\cdot u\\cdot v$, (neglecting the enthalpy) where $\\rho=2\\times 10^{-7}\\mathrm{g~cm^{-3}}$ is the mass density, $u=3\\times10^{12}~\\mathrm{erg~g^{-1}}$ is the heat deposited by $1~\\mathrm{g}$ of gas as it cools from $12000~\\mathrm{K}$ to $5000~\\mathrm{K}$ (Schlichenmaier et al. 1999), and $v=1~\\mathrm{km~s^{-1}}$ is the upflow velocity observed in the convective rolls. Entering these numbers, we obtain $F_{\\rm cr}\\approx 6\\times10^{10}~\\mathrm{erg~cm^{-2}~s^{-1}}$. This is larger than the radiative flux emitted by the penumbra, $F_{\\rm pen}\\approx 4.7\\times10^{10}~\\mathrm{erg~cm^{-2}~s^{-1}}=0.75\\cdot F_{\\sun}$. However, we must keep in mind that upflows fill at the most half of the surface area of the penumbra; e.g., the spines must be heated radiatively. It is reasonable to assume them to emit the same flux as the umbra in the absence of any lateral heat influx from the interspines, $F_{\\rm umb}=0.2\\cdot F_{\\sun}$. For a penumbra half covered by upflows and half by gas at umbral temperatures, we obtain the following relationship that must be fulfilled: $\\frac{1}{2}\\cdot\\left(\\rho\\cdot u\\cdot v+F_{umb}\\right)=F_{\\rm pen}$. This gives a requirement on $\\rho\\cdot u\\cdot v$ of $\\approx 8 \\times10^{10}~\\mathrm{erg~cm^{-2}~s^{-1}}$, which is only a factor of $1.3$ larger than the estimate obtained from the convective rolls. An underestimation of the partly unresolved velocities of the roll convection can easily account for this factor. The observed roll convection transports far more energy flux than interchange convection or the Evershed flow along flux tubes, as estimated by Schlichenmaier \\& Solanki (2003). We therefore propose that the observed convective rolls are the main form of energy transport in the immediate subsurface layers of the penumbra, carrying most of the energy required to maintain the radiative output of the penumbra.  The configuration of field and flows presented here combines aspects of both (i) the uncombed penumbra model of Solanki \\& Montavon (1993) and (ii) Schlichenmaier et al. (1998a,b) and the gappy penumbra model of Spruit \\& Scharmer (2006), Scharmer \\& Spruit (2006), cf. Scharmer et al. (2008). The geometry of magnetoconvection in the umbral dots models by Sch\\\"ussler \\& V\\\"ogler (2006), cf. Riethm\\\"uller et al. (2008), displays qualitative similarities to our model. Observations by Rimmele (2008) of a sunspot close to disk center reveal that the velocity pattern of dark-cored filaments at lower layers is similar to what he sees in the light bridges, i.e. the dark lanes are correlated with upflows and with downflows to the sides. This also agrees with our results. \\begin{figure} \\includegraphics[width=8cm]{0266fig4.eps} \\centering \\caption{Schematic illustration of the field and flow structure in penumbral filaments. Plotted is a vertical cut placed perpendicularly to the filaments (white semi circular shapes at the bottom of figure). The thick dashed-dotted line designates the $\\tau=1$ level for the plotted view direction. Vertical arrows denote magnetic field lines and arrows inside the lightened filaments denote convective flows. Circled crosses represent magnetic fields inside the filaments oriented along their axes (which also correspond to the direction of the Evershed flow).}\\label{fig4} \\end{figure}"
    },
    "0808/0808.1913_arXiv.txt": {
        "abstract": "*{ The era of exoplanet characterization is upon us. For a subset of exoplanets --- the transiting planets --- physical properties can be measured, including mass, radius, and atmosphere characteristics.  Indeed, measuring the atmospheres of a further subset of transiting planets, the hot Jupiters, is now routine with the {\\it Spitzer Space Telescope}.  The {\\it James Webb Space Telescope} ({\\it JWST}) will continue {\\it Spitzer's} legacy with its large mirror size and precise thermal stability.  {\\it JWST} is poised for the significant achievement of identifying habitable planets around bright M through G stars---rocky planets lacking extensive gas envelopes, with water vapor and signs of chemical disequilibrium in their atmospheres. Favorable transiting planet systems, are, however, anticipated to be rare and their atmosphere observations will require tens to hundreds of hours of {\\it JWST} time per planet. We review what is known about the physical characteristics of transiting planets, summarize lessons learned from {\\it Spitzer} high-contrast exoplanet measurements, and give several examples of potential {\\it JWST} observations. }  \\abstract{ The era of exoplanet characterization is upon us. For a subset of exoplanets --- the transiting planets --- physical properties can be measured, including mass, radius, and atmosphere characteristics.  Indeed, measuring the atmospheres of a further subset of transiting planets, the hot Jupiters, is now routine with the {\\it Spitzer Space Telescope}.  The {\\it James Webb Space Telescope} ({\\it JWST}) will continue {\\it Spitzer's} legacy with its large mirror size and precise thermal stability.  {\\it JWST} is poised for the significant achievement of identifying habitable planets around bright M through G stars---rocky planets lacking extensive gas envelopes, with water vapor and signs of chemical disequilibrium in their atmospheres. Favorable transiting planet systems, are, however, anticipated to be rare and their atmosphere observations will require tens to hundreds of hours of {\\it JWST} time per planet. We review what is known about the physical characteristics of transiting planets, summarize lessons learned from {\\it Spitzer} high-contrast exoplanet measurements, and give several examples of potential {\\it JWST} observations. } ",
        "introduction": "\\label{sec:intro} \\begin{figure}[ht] \\centering \\includegraphics[scale=.45]{f1.eps} \\caption{Known planets as of July 2008. We have defined planet to have a maximum mass of 13 Jupiter masses. The symbols indicate the discovery technique; see text for details. Data from \\cite{exoe}.} \\label{fig:exocensus} \\end{figure}  The existence of exoplanets is firmly established with over 300 known to orbit nearby, sun-like stars. Figure~\\ref{fig:exocensus} shows the known exoplanets as of July 2008 with symbols indicating their discovery techniques \\cite{exoe}.  The majority of the known exoplanets have been discovered by the Doppler technique which measures the star's line-of-sight motion as the star orbits the planet-star common center of mass \\cite{butl2006, udry2007}. While most planets discovered with the Doppler technique are giant planets, the new frontier is discovery of super Earths (loosely defined as planets with masses between 1 and 10~$M_{\\oplus}$).  About a dozen radial velocity planets with $M$$<$10~$M_{\\oplus}$ and another dozen with 10~$M_{\\oplus}$$<$$M$$<$30~$M_{\\oplus}$ have been reported. The transit technique finds planets by monitoring thousands of stars, looking for the small drop in brightness of the parent star that is caused by a planet crossing in front of the star. At the time of writing this article, around 50 transiting planets are known. Due to selection effects, transiting planets found from ground-based searches are limited to small semi-major axes \\cite{char2007a}. Gravitational microlensing has recently emerged as a powerful planet-finding technique, discovering 6 planets, two belonging to a scaled down version of our own solar system \\cite{gaud2008}. Direct imaging is able to find young or massive planets with very large semi-major axes. The mass of directly imaged planets (e.g.,  \\cite{chau2004} and references therein) is inferred from the measured flux based on evolution models, and is hence uncertain. The timing discovery method includes both pulsar planets \\cite{wols1992} and planets orbiting stars with stable oscillation periods \\cite{silv2007}.  Many fascinating properties of exoplanets have been uncovered by the initial data set of hundreds of exoplanets.  A glance at Figure~\\ref{fig:exocensus} shows one of the most surprising features: that exoplanets exist in an almost continuous range of mass and semi-major axis.  Not shown in Figure~\\ref{fig:exocensus} are the equally wide range of eccentricities; several different theories for the origin of planet eccentricities have been proposed.  Because there is not enough solid material close to the star in a protoplanetary disk, the giant planets are believed to have formed further out in the disk and migrated inwards. The migration stopping mechanisms, and even the details of planet migration are not fully understood.  Out of the $\\sim$50 known transiting exoplanets, several have very large radii, and are too big for their mass and age according to planet evolution models (see Figure~\\ref{fig:massradius}). These ``puffed-up'' planets must have extra energy in their core that prevents cooling and contraction, but no satisfactory explanation yet exists.  The next step beyond discovery is to characterize the physical properties of exoplanets by measuring densities, atmospheric composition, and atmospheric temperatures.  There are two paths to exoplanet characterization. The first is direct imaging where the planet and star are spatially separated on the sky. Direct imaging has been successful for discovering hot or massive planetary candidates with large ($\\sim$ 50--100 AU) orbital and projected spatial separation \\cite{chau2004, chau2005, neuh2005}. Although {\\it JWST} will incorporate several coronagraphic modes, neither the telescope nor the instruments were optimized for coronagraphy. A relatively large inner working angle and limited planet-star contrast restrict {\\it JWST} to studying young or massive Jupiters with large semi-major axes.  The case for {\\it JWST} coronagraphic observations of exoplanets is presented in \\cite{gard2006}.  Solar-system aged small planets are not observable via direct imaging with current technology, even though an Earth at 10 pc is brighter than the faintest galaxies observed by the {\\it Hubble Space Telescope} ({\\it HST}). The major impediment to direct observations is instead the adjacent host star; the Sun is 10 million to 10 billion times brighter than Earth (for mid-infrared and visible wavelengths, respectively).  No existing or planned telescope is capable of achieving this contrast ratio at 1~AU separations.  The second path to exoplanet characterization is via the transit technique. A subset of exoplanets cross in front of their stars as seen from Earth (``primary eclipse'' or ``transit''). Planets that cross in front of their star, also pass behind the star (secondary eclipse), provided that the transiting planet is on a circular orbit. The probability to transit is $\\sim R_*/a$, where $R_*$ is the stellar radius and $a$ the semi-major axis. Transits are therefore most easily found for planets orbiting close to the star. Indeed, all but one of the $\\sim$50 known transiting exoplanets have $a < 0.09$~AU (See \\cite{exoe} and references therein).  Observations of transiting planets exploit separation of photons in time, rather than in space (see Figure~\\ref{fig:transitchar}).  That is, observations are made in the combined light of the planet-star system.  When the planet transits the star as seen from Earth the starlight gets dimmer by the planet-to-star area ratio. If the size of the star is known, the planet size can be derived.  During the planet transit, some of the starlight passes through the optically thin part of the planet atmosphere, picking up spectral features from the planet. A planetary transmission spectrum can be obtained by dividing the spectrum of the star and planet during transit by the spectrum of the star alone (the latter taken before or after transit).   \\begin{figure}[ht] \\centering \\includegraphics[scale=.5]{f2.eps} \\caption{Schematic of a planet transit. Not to scale.} \\label{fig:transitchar} \\end{figure}  When the planet disappears behind the star, the total flux from the planet-star system drops. The drop is related to both relative sizes of the planet and star and their relative brightnesses (at a given wavelength). The flux spectrum of the planet can be derived by subtracting the flux spectrum of the star alone (during secondary eclipse) from the flux spectrum of both the star and planet (just before and after secondary eclipse). The planet's flux gives information on the planet composition and temperature gradient (at infrared wavelengths) or albedo (at visible wavelengths).  Finally, non-transiting exoplanets can in principle also be observed in the combined planet-star flux as the planet goes through illumination phases as seen from Earth. In the thermal infrared the phase observations provide information on energy redistribution of absorbed stellar radiation (see Section~\\ref{sec-atmsummary}) and at visible wavelengths the phase observations give information on scattering particles (gas or clouds).  Primary and secondary eclipses enable high-contrast measurements because the precise on/off nature of the transit and secondary eclipse events provide an intrinsic calibration reference.  This is one reason why the {\\it Hubble Space Telescope} and the {\\it Spitzer Space Telescope} have been so successful in measuring high-contrast transit signals that were not considered in their designs. ",
        "conclusions": "Several different exoplanets have published {\\it Spitzer} secondary eclipse atmosphere measurements. These secondary eclipse measurements have detection significances ranging from 60$\\sigma$ \\cite{knut2007} down to 5$\\sigma$. Rather than describe each planet individually, we present a summary based on an important question related to atmospheric circulation.  Hot Jupiters are expected to be tidally locked to their parent stars---presenting the same face to the star at all times.  This causes a permanent day and night side. A long standing question has been about the temperature difference from the day to night side. Are the hot Jupiters scorchingly hot on the day side and exceedingly cold on the night side? Or, does atmospheric circulation efficiently redistribute the absorbed stellar radiation from the day side to the night side?  Surprisingly, {\\it Spitzer} has found that both scenarios are possible. {\\it Spitzer} has measured the flux of the planet and star system as a function of orbital phase for several hot Jupiter systems \\cite{harr2006, cowa2007, knut2007}. Assuming that the star is constant in flux, the resulting brightness change is due to the planet alone.  The HD~189733 star and planet shows some variation at 8~$\\mu$m during the 30 hour continuous observation of half an orbital phase \\cite{knut2007}. This variation corresponds to about 20\\% variation in planet temperature (from a brightness temperature of 1212 to 973~K).  In contrast, the non-transiting exoplanet Ups And shows a marked day-night contrast suggesting that the day and night side temperatures differ by over 1000~K \\cite{harr2006}.  Once the stellar radiation is absorbed on the planet's day side, there is a competition between reradiation and advection. If the radiation is absorbed high in the atmosphere, the reradiation timescale is short and reradiation dominates over advection. In this case the absorbed stellar radiation is reradiated before it has a chance to be advected around the planet, resulting in a very hot planet day side and a correspondingly very cold night side. If the radiation penetrates deep into the planet atmosphere where it is finally absorbed, the advective timescale dominates and the absorbed stellar radiation is efficiently circulated around the planet. This case would generate a planet with a small temperature variation around the planet. See also \\cite{seag2005, harr2007, burr2008, fort2008}. See \\cite{show2007} and references therein for a review discussion of atmospheric circulation models.  In Figure~\\ref{fig:twoatm} we plot measured brightness temperatures of seven hot exoplanets together with two equilibrium temperature ($T_{eq}$) curves. One of the $T_{eq}$ curves is for a planet with evenly redistributed absorbed stellar radiation ($f=1/4$ below), corresponding to a planet with little temperature difference between the day and night sides.  The other $T_{eq}$ curve is for a planet with instantaneous reradiation of absorbed stellar radiation ($f=2/3$ below), corresponding to a planet with a strong day-night temperature difference. The cooler of the hot exoplanets in Figure~\\ref{fig:twoatm} lie along the evenly redistributed energy curve, while the hotter exoplanets lie nearer to the instantaneous reradiation $T_{eq}$ curve.  \\begin{figure}[ht] \\centering \\includegraphics[scale=.55]{f4.eps} \\caption{ Brightness temperature (8~$\\mu$m) as a function of the day side equilibrium temperature for six hot Jupiters. Brightness temperature is the measured flux converted to a temperature. The day side equilibrium temperature $T_{eq}$, is defined in equation~(\\ref{eq:Teq}) and the accompanying text.  The parameter $f$ is used to approximate atmospheric circulation: in this formulation of equation~(\\ref{eq:Teq}), $f=1/4$ corresponds to a uniform temperature around the planet (solid line), whereas $f=2/3$ corresponds to instantaneous reradiation on the planet's dayside hemisphere (dashed line).  The cooler planets lie near the uniform temperature line whereas the hotter planets lie near the hot day side line.  From left to right, the brightness temperatures are GJ 436b \\cite{demi2007}, HD 189733 b \\cite{knut2007}, TrES-1 \\cite{char2005}, HD 209458b \\cite{knut2008}, Ups And \\cite{harr2006}, HD 149026 \\cite{harr2007}. (Note that the Ups And day side temperature is estimated from its thermal phase curve).} \\label{fig:twoatm} \\end{figure}  Physically, $T_{eq}$ is the effective temperature attained by an isothermal planet after it has reached complete equilibrium with the radiation from its parent star. $T_{eq}$ is described by \\begin{equation} \\label{eq:Teq} T_{eq} = T_{*} \\left(\\frac{R_{*}}{a}\\right)^{1/2} [f(1-A_{\\rm B})]^{1/4} , \\end{equation} where $T_{*}$ and $R_{*}$ are the effective temperature and the radius of the star, $a$ is the planet semi-major axis, and $f$ and $A_{\\rm B}$ are the re-radiation factor and the Bond albedo of the planet.  Here we explain how hot Jupiters can exist both with and without large day-night temperature variations. See also \\cite{hube2003, harr2007, burr2008, fort2008}.  Hot planets such as Ups And are on one side of a temperature-driven chemical composition boundary, while cooler planets such as HD~209458b are on the cooler side. Specifically, if the hot Jupiter planet atmosphere is relatively cool, TiO is locked into solid particles that have little absorbing power in the atmosphere. In the hotter atmosphere, TiO is a ``deadly'' gas that absorbs so strongly it puts the planet in the reradiation regime leading to a large day-night contrast. At the temperature of these hot day side exoplanets, some elements will be in atomic (instead of molecular) form and atomic line opacities may also play a significant absorbing role.  What evidence do we have for the temperature-induced two-atmosphere hypothesis? Cool stars (M stars) have visible-wavelength spectra that are dominated---and indeed dramatically suppressed---by TiO gas.  Hot brown dwarfs also have spectra with TiO absorption features wheras cooler brown dwarfs do not, implying that for cooler brown dwarfs Ti is sequestered in solid particles. Temperature and pressures in hot Jupiter atmospheres are similar to brown dwarfs (although the temperature gradient is different) so that we expect a similar temperature-induced chemical composition.  Other notable exoplanet atmosphere discoveries by {\\it Spitzer} come from spectrophotometry and include discovery of a temperature inversion on the day side of HD~209458b \\cite{burr2007, knut2008} and a tentative detection of water vapor in transmission spectra during primary transit \\cite{tine2007,ehre2007}.  This interesting ``two types of hot Jupiter atmospheres'' hypothesis shows just how complex hot Jupiters are. The results also imply that 3D coupled radiative transfer-atmospheric circulation models are needed to fully understand hot Jupiters. Next generation data with {\\it JWST} (Section~\\ref{sec-JWSTatm}) in terms of high SNR low-resolution spectra as a function of orbital phase will lead to a deeper understanding of hot Jupiter atmospheres.  \\subsection{Lessons Learned from {\\it Spitzer}} \\label{sec-lessonsspitzer} {\\it Spitzer} and its instruments were not designed for high contrast observations. Here we describe some of the instrumental effects that become important at the part-per-thousand level and consequently affect exoplanet observations. These may be useful to consider when planning {\\it JWST} exoplanet transit observations.  The most notable instrumental effect is the ``ramp'': a gradual detector-induced rise of up to 10\\% in the signal measured in individual pixels over time. This rise is illumination-dependent; pixels with high levels of illumination converge to a constant value within the first two hours of observations and lower-flux pixels increase linearly over time.  \\cite{knut2007} have attributed this to charge trapping.  The ramp is present at the long wavelength detectors (8 and 16 $\\mu$m and possibly also at 5~$\\mu$m).  The charge trapping is likely caused by the ionized impurities in the arsenic-doped silicon detector. The first electrons that are released by photons get trapped by the ions and therefore they are not immediately read out.  The ramp can be removed from a data set by a fitting a linear plus logarithmic function.  A method to avoid the ramp is to ``preflash'' the detector before an observation, by observing a bright star. We note that at 24~$\\mu$m no ramp effect is observed \\cite{demi2005}; this may be because the detector is always illuminated by the zodiacal background. The ramp effect may be important for {\\it JWST} observations because similar detector materials are being used on MIRI.  We recommend pre-flashing before transit observations.  A second instrumental effect that is significant for exoplanet transit observations is the IRAC intrapixel sensitivity variation for the short-wavelength channels (3.5 and 4.5 $\\mu$m) \\cite{mora2006}. This intrapixel sensitivity variation is a property of the detector, is spatially asymmetric, and does not change with time. One possible explanation of this detector effect is that there are physical gaps between the pixels which are unresponsive to light. When an image is centered on a pixel the highest sensitivity arises. A second possibility is related to the bump-bond contact between the detector and the underlying multiplexer. Each pixel makes electrical contact with the multiplexer at pixel center.  In this scenario, the electrons generated close to the contact point may be collected more efficiently than electrons generated at pixel edges.  The IRAC intrapixel sensitivity variation can be corrected for by a natural mapping of the pixels, by exploiting the telescope pointing jitter \\cite{demi2008}.  A third significant instrumental effect found from exoplanet observations is a telescope pointing oscillation with a period of roughly one hour. This pointing oscillation affects IRS slit spectra. As the star becomes uncentered and recentered on the slit the number of photons going through the slit changes. This change is wavelength-dependent, because of the wavelength-dependent PSF, and therefore can generate erroneous features in an exoplanetary spectrum.  The cause of the pointing oscillation is still a matter of debate.  A fourth effect is a longer-term telescope drift in telescope pointing. This telescope drift is less certain than the well-documented 1 hour oscillation. This affects the IRS spectral flux in terms of the broad distribution of energy, because the slit width is comparable to the diffraction width of the telescope PSF. This is also a wavelength-dependent effect.  Despite not being designed for high-contrast observations, {\\it Spitzer} has revolutionized the study of exoplanet atmospheres by routinely detecting secondary eclipses of hot Jupiters and measuring their brightness temperatures at different wavelengths. Because of {\\it JWST}'s larger mirror diameter and higher spectral resolution, and by taking care to understand and remove instrumental effects, transit observations with {\\it JWST} hold even more promise.    The {\\it Hubble Space Telescope} and the {\\it Spitzer Space Telescope} (Section~\\ref{sec-lessonsspitzer}) were not designed to achieve the very high SNRs necessary to study transiting exoplanets, but they have succeeded nonetheless.  We have learned from these observations that with enough photons, systematics that were unknown in advance can often be corrected for (as long as the errors are uncorrelated). Given the expected thermal stability of {\\it JWST} and the differential nature of transit observations, we are optimistic that {\\it JWST}, too, will succeed in high-contrast transit observations.  We have described a few examples where {\\it JWST} will have a significant impact, including measuring precise exoplanet radii for all sizes of exoplanets and spectra for giant planets at a variety of semi-major axes. Individually, these observations will be ``cheap'' in terms of telescope time, with single transit measurements being sufficient.  The most significant exoplanet observations {\\it JWST} is poised to make are those for potentially habitable exoplanets. In the most optimistic case, {\\it JWST} is able to identify planets with atmospheric water vapor, or even chemical disequilibrium indicative of biological origin. {\\it JWST} has the capability to do so if the tens to hundreds of hours per target are allocated, and if such rare transiting planets can be discovered in sufficient numbers.  \\begin{acknowledgement} We thank Mark Clampin, George Ricki, Dave Charbonneau, and Heather Knutson for useful discussions.  \\end{acknowledgement}"
    },
    "0808/0808.1122_arXiv.txt": {
        "abstract": "Observations have ruled out the presence of significant amounts of antimatter in the Universe on scales ranging from the solar system, to the Galaxy, to groups and clusters of galaxies, and even to distances comparable to the scale of the present horizon.  Except for the model-dependent constraints on the largest scales, the most significant upper limits to diffuse antimatter in the Universe are those on the $\\sim$~Mpc scale of clusters of galaxies provided by the EGRET upper bounds to annihilation gamma-rays from galaxy clusters whose intracluster gas is revealed through its x-ray emission.  On the scale of individual clusters of galaxies the upper bounds to the fraction of mixed matter and antimatter for the 55 clusters from a flux-limited x-ray survey range from $5 \\times 10^{-9}$ to $<  1 \\times 10^{-6}$, strongly suggesting that individual clusters of galaxies are made entirely of matter or, of antimatter.  X-ray and gamma-ray observations of colliding clusters of galaxies, such as the Bullet Cluster, permit these constraints to be extended to even larger scales.  If the observations of the Bullet Cluster, where the upper bound to the antimatter fraction is found to be $< 3 \\times 10^{-6}$, can be generalized to other colliding clusters of galaxies, cosmologically significant amounts of antimatter will be excluded on scales of order $\\sim 20$ Mpc ($M \\sim 5 \\times 10^{15}M_{\\odot}$). ",
        "introduction": "\\label{intro}  Had the Universe had been matter-antimatter (baryon-antibaryon) symmetric during the later stages of its early evolution, when the temperature was below that of the quark-hadron transition (and well below the nucleon mass), it would have experienced an ``annihilation catastrophe\" in the sense that the number of post-annihilation nucleons would be a billion -- or more -- times less abundant than is observed in our present Universe (for a review and an extensive list of references, see \\cite{steig76}).  Even worse, for a symmetric universe the annihilation which ceased during its early evolution due to the low density of the surviving nucleon-antinucleon pairs, would resume when gravitationally-collapsed objects formed (if, indeed, collapsed objects could form in such a baryon-poor universe), further reducing the baryon density and inhibiting the formation of stars, planets, etc.  The problems of a symmetric Universe were appreciated by Sakharov~\\cite{sakharov}, who outlined the necessary conditions for generating a matter-antimatter asymmetry during the early evolution of the Universe.  Later, when it was understood that Grand Unified Theories, along with the expected, early evolution of the standard, hot big bang cosmological model, contained the ingredients of Sakharov's recipe for generating a baryon asymmetry \\cite{yoshimura,susskind,toussaint,weinberg}, it became generally accepted that our Universe is matter-antimatter asymmetric, consisting predominantly of ordinary matter (by definition!) and containing, at most, only trace amounts of antimatter.  For a recent, contrary point of view, see \\cite{bambi1,bambi2,dolgov}.  Over the years, the theoretical expectations of a baryon asymmetric Universe have been tested and confirmed by a wide variety of observations which strongly limit the observationally allowed amount of cosmological antimatter \\cite{steig76,wolfendale,cdg,bambi1,bambi2,dolgov}.  During the collapse of stars and, in particular, of the pre-solar nebula, any relic antimatter would have annihilated to extremely low levels.  Even so, it is useful to note that the absence of annihilation gamma rays, which could have been produced as the solar wind sweeps over the planets, strongly restricts the presence of significant amounts of antimatter in the solar system \\cite{steig76}. Beyond the solar system, only the cosmic rays provide direct evidence of the composition of the stars and gas in the Galaxy.  The absence of complex antinuclei (\\eg antihelium) in the cosmic rays at a level $< 10^{-6}$ \\cite{ams1} provides an interesting upper bound to antimatter in the Galaxy.  New cosmic ray experiments such as PAMELA \\cite{pamela} and AMS \\cite{ams2} have the potential to reduce this limit further or, to find evidence for antimatter from observations of galactic or extragalactic cosmic ray antinuclei.  While a collapsed object, such as a star, may hide appreciable amounts of antimatter\\footnote{In its journey through the Galaxy, an antistar would accrete interstellar gas, leading to the production of annihilation gamma-rays.  Observations of discrete Galactic gamma-ray sources limit the fraction of antistars in the Galaxy to $< 10^{-4}$ \\cite{steig76}.}, stars are formed from interstellar gas and, in the course of their evolution and, at the end of their evolution, they return substantial amounts of material to the interstellar medium.  Since an antinucleon (or, antinucleus) in the diffuse interstellar medium (ISM) of the Galaxy has a very short lifetime against annihilation, $\\sim 300$~yr -- 200 kyr \\cite{steig76}, any antimatter present when the Galaxy formed wouldn't have survived to the present epoch\\footnote{The $\\sim 3 - 5$~order of magnitude longer lifetime in Bambi and Dolgov \\cite{bambi1}, while still very small compared to the age of the Galaxy, fails to account for the $\\sim 3 - 5$ order of magnitude enhancement of the annihilation cross section at the very low collision energies in the ISM.  As a consequence, the Bambi and Dolgov gamma-ray flux estimates \\cite{bambi1} may need to be corrected upwards by 3 -- 5 orders of magnitude.}.  Indeed, the observed Galactic gamma rays indirectly limit the ratio of antimatter to matter in the ISM to $< 10^{-15}$ \\cite{steig76}.  On scales larger than that of the Galaxy, upper bounds to the observed gamma-ray flux provide indirect limits on the presence of diffuse regions of mixed matter and antimatter.  In \\S\\ref{clusters}, the gamma-ray limits on matter-antimatter annihilation in the hot, x-ray emitting gas of clusters of galaxies are reviewed, leading to bounds on the antimatter fraction in systems of size $\\sim$~few Mpc and mass $\\sim 10^{15}M_{\\odot}$.  In \\S\\ref{bullet} it is noted that observations of the ``Bullet Cluster\" \\cite{markevitch} and of other colliding clusters of galaxies \\cite{jee,mahdavi,bradac} permit these bounds to be extended to larger distance/mass scales. The results presented here are summarized in \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  Direct observations in the solar system and of the galactic cosmic rays set stringent constraints on antimatter in our local vicinity.  Gamma rays produced when matter and antimatter meet and annihilate provide indirect evidence for regions of mixed matter and antimatter.  There is no evidence for such mixed regions on scales from galaxies, to groups and clusters of galaxies, limiting the antimatter fraction to $< 1\\times 10^{-6}$, or smaller, on scales up to $\\sim$~Mpc and $M \\sim 10^{15}M_{\\odot}$.  Recent observations of clusters in collision \\cite{markevitch,jee,mahdavi,bradac} permit the x-ray cluster bounds on the fraction of antimatter in the Universe to be extended to even larger scales.  Observations of the colliding galaxy clusters which constitute the Bullet Cluster limit the antimatter fraction to $f_{\\rm Bullet} < 3\\times 10^{-6}$, on a scale of $M \\sim 6\\times 10^{15}h_{50}^{-1}M_{\\odot}$.  If, indeed, there are regions of antimatter in the Universe, they must be separated from regions of ordinary matter by distances on the order of tens of Mpc (mass scales of order $10^{16}M_{\\odot}$).  Evidence for galaxy clusters in collision suggest that these scales can be probed in the near future."
    },
    "0808/0808.3916_arXiv.txt": {
        "abstract": "We apply the joint lensing and dynamics code for the analysis of early-type galaxies, {``\\cauldron''}, to a rotating N-body stellar system with dark matter halo which significantly violates the two major assumptions of the method, i.e. axial symmetry supported by a two-integral distribution function. The goal is to study how {\\cauldron} performs in an extreme case, and to determine which galaxy properties can still be robustly recovered. Three data sets, corresponding to orthogonal lines of sight, are generated from the N-body system and analysed with the identical procedure followed in the study of real lens galaxies, adopting an axisymmetric power-law total density distribution. We find that several global properties of the N-body system are recovered with remarkable accuracy, despite the fact that the adopted power-law model is too simple to account for the lack of symmetry of the true density distribution. In particular, the logarithmic slope of the total density distribution is robustly recovered to within less than $10\\%$ (with the exception of the ill-constrained very inner regions), the inferred angle-averaged radial profile of the total mass closely follows the true distribution, and the dark matter fraction of the system (inside the effective radius) is correctly determined within $\\sim 10\\%$ of the total mass. Unless the line of sight direction is almost parallel to the total angular momentum vector of the system, reliably recovered quantities also include the angular momentum, the $V/\\sigma$ ratio, and the anisotropy parameter $\\delta$. We conclude that the {\\cauldron} code can be safely and effectively applied to real early-type lens galaxies, providing reliable information also for systems that depart significantly from the method's assumptions. ",
        "introduction": "\\label{sec:introduction}  Determining the structure of early-type galaxies and reliably pinning down their dark matter content is a crucial step in order to fully understand the formation and evolution processes of these systems.  Within the currently favoured cosmological scenario, the $\\Lambda$CDM paradigm, early-type galaxies are thought to be formed via hierarchical merging of lower mass galaxies \\citep{Toomre1977, White-Frenk1991, Barnes1992, Cole2000}. While very successful in reproducing many observational features of elliptical galaxies, including complex ones \\citep[see e.g.][]{Jesseit2007}, these formation models are still encountering difficulties in explaining the origin of the empirical scaling laws that correlate the global properties of early-type galaxies \\citep[see e.g.][]{Robertson2006}. Providing a reliable and detailed description of the mass density distribution, orbital structure and intrinsic properties of early-type galaxies is therefore critical in order to enable stringent tests of galaxy formation models.  For this reason, considerable effort has been devoted during the last decades towards the observation and the modelization of nearby early-type galaxies, by means of both stellar dynamics and X-ray studies, finding more or less strong evidence for a dark matter halo component (e.g. \\citealt{Fabbiano1989}, \\citealt{Mould1990}, \\citealt*{Saglia1992}, \\citealt{Bertin1994}, \\citealt{Franx1994}, \\citealt{Carollo1995}, \\citealt{Arnaboldi1996}, \\citealt{Rix1997}, \\citealt{Matsushita1998}, \\citealt{Loewenstein1999}, \\citealt{Gerhard2001}, \\citealt{Borriello2003}, \\citealt{Romanowsky2003}, \\citealt{Humphrey2006}, \\citealt{Forbes2008} and more recently the SAURON collaboration: see e.g. \\citealt{SauronII}, \\citealt{SauronIII}, \\citealt{SauronIV}). Both methods, however, present some difficulties. In the case of stellar dynamics it is believed that some degeneracy can be present between the mass profile of the galaxy and the anisotropy of the stellar velocity dispersion tensor, which can be alleviated when higher order velocity moments are available \\citep[see e.g.][]{Gerhard1993} or by using physically motivated distribution functions \\citep[see e.g.][for a discussion of this point]{Bertin2000}. X-ray analyses, on the other hand, can seriously overestimate the total mass of the system if the assumption of hydrostatic equilibrium for the hot gas does not hold, especially near the center \\citep[see e.g.][]{Pellegrini2006, Ciotti-Pellegrini2008}.  A full understanding of the evolution of early-type galaxies cannot be achieved without extending the study also to the mass density profile of objects at higher redshift ($z \\ga 0.1$). This, however, has not been attempted until recently, due to observational limitations and to the increased difficulty in extracting detailed kinematic information, which hinders traditional analyses based on stellar dynamics only. An effective solution in order to overcome these issues is constituted by a joint analysis which combines the constraints from stellar dynamics with the information obtained from gravitational lensing, when the early-type galaxy also happens to act as a lens with respect to a background source at higher redshift \\citep{Treu-Koopmans2002, Koopmans-Treu2002, Treu-Koopmans2003, Treu-Koopmans2004, vandeVen2008}.  \\citet{Koopmans2006} have successfully used this combined approach to analyse fifteen early-type lens galaxies (within a redshift range $z = 0.06 - 0.33$) discovered in the Sloan Lens ACS Survey (SLACS, \\citealt{Bolton2006}) as well as six systems between $z \\sim 0.5$ and $1$ from the Lenses Structure and Dynamics (LSD) Survey, showing that all of the examined systems are well described by a power-law total density distribution very close to $r^{-2}$. The technique for the joint lensing and dynamics analysis has been expanded by \\citet[][hereafter BK07]{Barnabe-Koopmans2007} into a general and self-consistent method, completely embedded within the framework of Bayesian statistics, which puts constraints on the total density distribution of the lens galaxy by taking advantage of all the available data, i.e. not only the lensed image and a single stellar velocity dispersion measurement, but also the surface brightness distribution and the 2D kinematic maps (first and second projected velocity moments).  Similar to other methods for the determination of the structure and internal dynamics of early-type galaxies, already mentioned above, the joint lensing and dynamics analysis also relies on a certain number of assumptions. For example, the simple and robust approach of \\citet{Koopmans2006} treats the gravitational lensing and the stellar dynamics as independent problems. The projected mass distribution of the lens galaxy is modelled as a singular isothermal ellipsoid in order to determine the total mass within the Einstein radius, which is then used as a constraint for the dynamical model, where spherical symmetry and a specific prescription for the stellar orbital anisotropy are assumed. The more sophisticated framework of BK07 is designed to be very general and allows in principle for an arbitrary choice of the total potential. In practice, however, such freedom must balance against technical and computational limitations. Therefore, in order to have a fast and efficient algorithm, the current implementation of the method, the {\\cauldron} code\\footnote{Combined Algorithm for Unified Lensing and Dynamics \\mbox{ReconstructiON}}, is restricted to axisymmetric potentials and two-integral stellar phase-space distribution functions (DFs). Under these hypotheses, it has been shown in BK07 that the method is capable of recovering with considerable accuracy the correct potential parameters and inclination angle, even in the presence of realistic noise.  The point above raises the question of whether (and to what extent) the simplifying assumptions can be deemed valid for the astrophysical systems to which such methods are applied.  In fact, real galaxies are not idealized objects, and there is no reason to expect them to be exactly axisymmetric (and neither triaxial ellipsoids) or to have two or three integrals of motion. Whereas axial symmetry generally seems to constitute a fairly good approximate description for most early-type galaxies, a more detailed inspection, such as that allowed by the SAURON observations \\citep[see][]{SauronIII, SauronVIII}, reveals a multitude of features indicating departure from axisymmetry, e.g. the presence of isophotal twist in the surface brightness distribution, minor axis rotation and kinematically decoupled cores.  For this reason, in the present paper we apply our algorithm to the end-product of a two-component (stars plus dark matter) N-body simulation of a merger process, i.e. to a system which does not obey any restrictive prescription of symmetry, and therefore violates the assumptions of the method. We aim to study how the {\\cauldron} algorithm performs when subjected to this kind of ``crash-test'', and which quantities can be robustly recovered even in such an extreme case. A similar approach has been followed by \\citet{Thomas2007}, who have applied their three-integral axisymmetric Schwarzschild code to the study of non axisymmetric N-body merger remnants, although without any gravitational lensing information, and by \\citet{Meneghetti2007} in the case of clusters of galaxies.  The paper is organized as follows. In Section~\\ref{sec:code} we provide an overview of the {\\cauldron} algorithm for combined lensing and dynamics analysis. In Section~\\ref{sec:simulation} we summarize the properties of the N-body system that we use as lens galaxy. In Section~\\ref{sec:observables} we describe how the 2D maps of the simulated data with added realistic noise are obtained from the particle distribution. In Section~\\ref{sec:results} we apply the joint lensing and dynamics analysis to the simulated data and we present the results, which are then further discussed in Section~\\ref{sec:conclusions}, where also conclusions are drawn. ",
        "conclusions": "\\label{sec:conclusions}  We have applied {\\cauldron}, currently the most advanced code for joint gravitational lensing and stellar dynamics analysis of early-type galaxies, to a galaxy model with dark matter halo resulting from a numerical N-body simulation of galaxy merging. Such a N-body system, which we use as lens, significantly violates the two major assumptions upon which the algorithm is based, namely axial symmetry of the total density distribution and two-integral stellar distribution function. The purpose of this crash test is to investigate how the code will perform or fail in an extreme case, and to identify the quantities which can still be reliably recovered. Such robust quantities are expected to be recovered with at least comparable accuracy when {\\cauldron} is employed in the analysis of real galaxies which, while not necessarily axisymmetric or described by a two- (or three-) integral DF, will hardly depart from the code's assumptions more severely than the simulated system studied here.  Further complications can also arise, in general, due to the effects of the environment on the lens galaxy. However, the SLACS galaxies, despite living in overdense regions as expected for massive early-type galaxies, are not found to be significantly affected by the contribution of the environment or of line-of-sight contaminants \\citep[see the in-depth study of][]{Treu2008}. Therefore, since the environment does not play a major role for lens galaxies at least in the redshift range $z \\sim 0.1 - 0.4$, we have not considered this issue further in the present paper.  From the N-body system we have generated three data sets corresponding to three orthogonal lines of sight, one of which has been chosen to be approximately oriented along the total angular momentum of the system, in order to obtain a ``face-on'' projection data set with little or no rotation discernible in the kinematic maps. An elliptical Gaussian surface brightness distribution has been constructed and then gravitationally lensed by the mock galaxy in order to create the lensing data set. We have also taken into account the effect of the PSF and added realistic noise to the simulated data, using as a reference the real data set for the SLACS lens galaxy {\\galaxy} studied in C08.  These data sets have been analysed with {\\cauldron}, assuming as a model an axisymmetric power-law total density distribution, with the identical procedure followed in the study of real lens galaxies. In the three cases, the recovered best models, obtained via maximization of the Bayesian evidence, show clear difficulties in reconstructing the observables (in particular the lensed image and the velocity dispersion map) up to the noise level. This is a consequence of having adopted a very simple model which cannot account for the complexity and the lack of symmetry of the true density distribution.  However, the method is still capable of recovering with remarkable accuracy several global structural and dynamical characteristics of the examined system, provided that some information about the galaxy rotation is available from the kinematic maps. In particular:  \\begin{enumerate} \\item The logarithmic slope $\\slope$ of the total density distribution is a robust quantity which is recovered with remarkable accuracy (within less than $10 \\%$), even in the case of the face-on projection data set. While a power-law model cannot account for a break in the actual density profile, indications of its presence will show up in the observables as features which are not reproduced by the best model (e.g. a much flatter central velocity dispersion). \\item The angle-averaged total density and total mass radial profiles for the edge-on projection best models are found to closely follow the corresponding true distributions, within approximately an effective radius. Discrepancies are larger for the face-on projection case. \\item The best reconstructed inclination angle and flattening of the total density distribution have little or no relation with the corresponding quantities of the N-body system. We conclude that when the axisymmetric model is applied to a system with a more complex geometry, one can not expect to recover reliable information about its shape. \\item By adopting the maximum bulge approach (i.e. maximizing the contribution of the luminous component, assuming a position-independent stellar mass-to-light ratio), it is possible to estimate within approximately $10 \\%$ (in total mass) the dark matter fraction of the analysed system, whether the mass ratio is calculated within a sphere or within a line-of-sight oriented cylinder of radius $\\simeq \\Reff$. \\item When rotation is present in the kinematic maps, global quantities such as the angular momentum $\\Lz$ and the ratio $V/\\sigma$ (a measure of ordered vs chaotic motions) recovered from the best model describe quite reliably (i.e. within $\\sim 10 \\%$ and $25 \\%$, respectively) the dynamical properties of the system under analysis. The global anisotropy parameters $\\beta$ and $\\gamma$ are not correctly estimated, i.e. the anisotropy distribution is not robustly recovered by the method unless the analyzed system effectively respects the assumptions of axisymmetry and two-integral distribution function. The anisotropy parameter~$\\delta$, on the contrary, is found to be a robust quantity even when such assumptions are violated: we recover the correct value within $\\la 15 \\%$ for the edge-on projection data sets. \\end{enumerate}  The major conclusion of our study is that the joint lensing and dynamics code {\\cauldron} can be effectively applied also to the analysis of galaxies which deviate from the (quite restrictive) assumptions of axial symmetry and two-integral stellar DF. This result is very relevant for the analysis of galaxies at $z \\gtrsim 0.1$ for which, due to the data limitations, the more powerful methods available at lower redshifts are not viable techniques.  Several fundamental structural and dynamical quantities, in particular the total density slope and the dark matter fraction within the region where the data are available, can be recovered with good accuracy. Other quantities, such as $\\Lz$ and the anisotropy parameter $\\delta$, can be reliably recovered provided that rotation is detected in the kinematic maps: when this is not the case (either because of intrinsic properties of the galaxy or because the system is observed face-on) the constraints are much looser and one needs to be very sceptical about the outcomes of the best model relative to these quantities. We point out that special care should be taken when the best model does not manage to reproduce the observables satisfactorily (i.e. at or close to the noise level), in particular the lensed image. This is a strong indication that the true density distribution deviates significantly from the assumptions of the adopted model family, and therefore, while there are still reliable recovered parameters (as listed above), more delicate quantities such as flattening and inclination angle, as well as the reconstructed source, should not be trusted. In these cases, the {\\cauldron} code can provide a robust but only quite general description of the structure of the galaxy, paving the way for the more sophisticated modelling techniques which are necessary (together, if possible, with a better data set) to achieve a deeper and more detailed knowledge of the system.  Interestingly, large residuals in the lensed images reconstruction such as the ones found in this study have never been encountered so far in the analysis of the SLACS sample of lens galaxies (\\citealt{Koopmans2006}, C08) indicating that an axisymmetric single power-law total density distribution constitutes a satisfactory model for these systems."
    },
    "0808/0808.3850_arXiv.txt": {
        "abstract": "We present a catalogue of $^{12}$CO(J=1-0) and $^{13}$CO(J=1-0) molecular clouds in the spatio-velocity range of the Carina Flare supershell, GSH 287+04-17. The data cover a region of $\\sim66$ square degrees and were taken with the NANTEN 4m telescope, at spatial and velocity resolutions of $\\sim2.6'$ and $0.1$ km s$^{-1}$. Decomposition of the emission results in the identification of 156 $^{12}$CO clouds and 60 $^{13}$CO clouds, for which we provide observational and physical parameters. Previous work suggests the majority of the detected mass forms part of a comoving molecular cloud complex that is physically associated with the expanding shell. The cloud internal velocity dispersions, degree of virialization and size-linewidth relations are found to be consistent with those of other Galactic samples. However, the vertical distribution is heavily skewed towards high-altitudes. The robust association of high-$z$ molecular clouds with a known supershell provides some observational backing for the theory that expanding shells contribute to the support of a high-altitude molecular layer. ",
        "introduction": "The mechanical feedback from OB clusters has a major impact on the structure and evolution of the interstellar medium. Supershell shocks are known to drive correlated episodes of triggered star formation -- either through the compression of existing molecular material to an unstable state \\citep[e.g.][]{boss95, vanhala98}, or through the manufacture and subsequent collapse of parent molecular clouds from the accumulated gas \\citep{bergin04, hartmann01, mccray87, koyama00}. Observational studies of molecular supershells have the power provide constraints on theories of triggered cloud and star formation, and to accurately pinpoint locations at which such triggering may be occurring. Yet, with a handful of exceptions \\citep{fukui99, jung96, matsunaga01, moriguchi02, normandeau96, rizzo98,  yamaguchi99, yamaguchi01b}, supershells with a significant molecular component remain relatively understudied.\\par  The CO molecule has been widely and successfully used as a tracer of the molecular ISM, and has formed the basis of many large-scale surveys of the Galactic molecular disk \\citep[e.g. ][]{dame87, heyer98, mizuno04, sanders86}. Uniform, wide-coverage CO datasets are indispensable for the information they provide on the distribution and physical properties of the molecular ISM, and the decomposition of such datasets into isolated areas of emission enables the statistical analysis of large populations of clouds in a variety of environments \\citep[e.g.][]{fukui08, heyer01, kawamura98, solomon87, yonekura97}.  This paper presents a catalogue of CO molecular clouds in the region of the `Carina Flare' supershell (GSH 287+04-17). The Carina Flare is an H{\\sc i}-H$_2$ supershell located in the near side of the Carina Arm at $D=2.6\\pm0.4$ kpc. It was first discovered in $^{12}$CO(J=1-0) with the NANTEN 4m telescope by \\citet{fukui99}, who reported the discovery of an extensive, expanding complex of molecular clouds reaching as high as $z\\sim450$ pc above the Galactic Plane. Comparisons with H{\\sc i} 21 cm survey data confirmed that the majority of the mass in this molecular complex -- as well as much additional lower-latitude material -- forms co-moving parts of a gently expanding atomic supershell \\citep{dawson08}. This supershell extends from $b\\sim0.5^{\\circ}$ to $b\\sim10^{\\circ}$, covering a total projected area of $\\sim70$ square degrees, with a main body $360\\times260$ pc in size, and a capped, high-latitude extension. The total associated molecular mass is estimated as $M(\\mathrm{H}_2)=2.0\\pm0.6\\times10^5 M_{\\solar}$, which amounts to $\\sim20\\%$ of the total neutral gas mass, and this molecular component is associated with a rich variety of substructural features in the atomic shell. Moreover, several molecular clouds show evidence for active star formation, suggesting triggering may be occurring \\citep{dawson07, fukui99}.  Previous work has discussed the macroscopic distribution and properties of the Carina Flare molecular ISM. However, detailed data on individual clouds, as well as the statistical properties of the sample as a whole, have not previously been published. Here we present the first comprehensive catalogue of the cloud population, intended both as a complement to past work, and as a resource upon which future studies of the region may draw. The present study includes observations in both $^{12}$CO(J=1-0) and $^{13}$CO(J=1-0). The $^{12}$CO line, which is generally assumed to be optically thick, is an excellent tracer of the overall molecular gas distribution, down to diffuse, low-luminosity envelopes. This is complemented by the $^{13}$CO line, which is optically thin up to H$_2$ column densities of $N(\\mathrm{H}_2)\\sim10^{22}$ cm$^{-2}$, and correlates well with denser material and star forming regions.  As well as the obvious interest that comes from investigating the molecular component of a known supershell, the present dataset also has several advantages that make it particularly attractive for this kind of large-scale cloud catalogue. Firstly, the association of the molecular clouds with a single object means that their distances are both uniform and well-determined, reducing what would otherwise be a substantial source of uncertainty in the determined cloud properties. Secondly, the majority of the complex is at high enough latitudes that line-of-sight confusion is minimal. This means that emission may be assigned to individual clouds without the need to resort to artificially high intensity thresholds, and leads to the inclusion of low-luminosity cloud envelopes that may be missed by other studies. In addition, the present dataset is both sensitive and relatively well resolved, with rms noise levels per 0.1 km s$^{-1}$ channel of $< 0.5$ K in $^{12}$CO(J=1-0) and $< 0.2$ K in $^{13}$CO(J=1-0), and telescope half power beam widths of $\\sim2.6'$ and $\\sim2.7'$ respectively.  The paper is organized as follows. Section \\ref{observations} describes the instrumentation and observing strategy. Section \\ref{cat} presents the $^{12}$CO and $^{13}$CO cloud catalogue tables, and   also describes the cloud decomposition method, the derivation of cloud properties and the completeness and coverage of the catalogue. Section \\ref{results} examines some of the statistical properties of the cloud sample as a whole, and section \\ref{discussion} briefly discusses the vertical distribution of the molecular gas, as well as some of the issues surrounding mass estimation. We summarize in section \\ref{summary}.\\\\\\par ",
        "conclusions": "\\label{discussion}  \\subsection{Vertical Distribution of the Molecular Gas}  Figure \\ref{vmassfig} shows the variation of mean $M_{\\mathrm{CO}}$ per $2'\\times2'$ pixel with altitude above the midplane, $z$. Also shown is a histogram of cloud count versus $z$. The vertical mass distribution may be discussed quantitatively in terms of a standard Galactic molecular gas scale height. The Carina Flare clouds are located at a Galactic radius of $R\\approx8$ kpc, assuming $R_{\\solar}=8.5$ kpc \\citep{dawson08}. For the molecular disk near the solar circle, the standard value for the Gaussian scale height of the gas is $\\sigma_z\\approx60$ pc \\citep{bronfman88, clemens88, malhotra94a}. For a distribution that follows this model we would therefore expect to find only $\\sim0.1\\%$ of the total mass at latitudes greater than $3.3\\sigma_z$ ($\\approx200$ pc). For the present dataset the figure is $\\sim7\\%$ -- nearly two orders of magnitude greater. Moreover, this fraction includes a large contribution from sub-complexes as high as $\\sim6\\sigma_z$ and above.\\par  Previous studies report some evidence for two different molecular cloud populations in the Galaxy, with differing scale heights, and possibly also with different physical properties \\citep{dame94, malhotra94b}. \\citet{dame94} find that large-scale CO survey data from the first quadrant is best fitted by two Gaussians rather than one, with a thin layer of $\\sigma_z\\sim30$ pc and a thick, less massive component  of $\\sigma_z\\sim100$ pc. Similarly, \\citet{malhotra94b} has examined the properties of a distinct population of high-$z$ clouds that are outliers to her single-component gaussian fits. Though the data were less sensitive than the present observations, and the $z$ coverage lower, we find that these outliers share many similarities with our high-$z$ catalogue clouds in terms of size, mass and typical velocity dispersions.\\par  Supershells have of course been suggested as one potential means of supporting this high-$z$ population -- either through the transport of existing clouds, or through in-situ formation from swept up atomic gas. The present catalogue provides solid observational proof of high-altitude molecular clouds that are not only robustly associated with a known supershell, but which also share similarities with the previously identified thick-disk component. Although a single object can offer no insight into the relative contribution of supershells to the support of high altitude clouds on a Galaxy-wide scale, the present work does offer a strong indication that they play a role in that support.  It should be noted that the limited longitude coverage of the catalogue at $b < 2^{\\circ}$ results in the selective inclusion of the very mass-rich region of the GMC and the exclusion of the region around it, potentially biasing the results. An alternative treatment might choose to include the full range $284^{\\circ} < l < 292^{\\circ}$, the overall CO distribution of which is shown in figure 10 of \\citet{fukui99}. This would have the effect of increasing the contribution from emission physically distant from the supershell, but would avoid unnaturally inflating the average mass at the latitude of the GMC. We find that the inclusion of these non-catalogue regions does not affect the overall proportion of the mass found at high latitudes, which remains at $\\sim7\\%$ for $z > 200$ pc. Since the low latitude contribution as a whole is likely to be inflated by the contribution of distant emission, this figure is a lower limit. \\\\\\par   \\subsection{Remarks on Mass Estimation}  It is worth making a few comments on the uncertainties associated with molecular cloud mass estimation, which are often passed over in CO cloud studies.  The empirical `X-factor', used to convert $^{12}$CO luminosity to H$_2$ column density, has been derived from a number of independent methods, and is believed to be accurate to a factor of around two, when averaged over large cloud populations \\citep[see e.g.][]{dame01, hunter97, solomon91}. For individual clouds the uncertainty will be larger \\citep[e.g.][]{magnani95}. In particular, \\citet{heyer01} have demonstrated how the correct X-factor for a cloud depends on the gravitational state of the gas, and argue that it may significantly overestimate the masses of objects that are not bound by self-gravity. Conversely, several authors have used the results of PDR modeling to argue that the standard Galactic X-factor may significantly \\textit{underestimate} the masses of many molecular clouds \\citep{bell06, kaufman99, pringle01}, especially those that are particularly diffuse.  $^{13}$CO LTE-based mass estimates are also subject to their own uncertainties. \\citet{padoan00} use model clouds with realistic temperature and density structures to argue that LTE analysis may underestimate true column density by a factor of 1.3 to 7, with the most extreme differences occurring in regions of low column density. The LTE values of $N(^{13}\\mathrm{CO})$ observed in the present dataset range from $\\sim5\\times10^{14}$ cm$^{-2}$ to $\\sim8\\times10^{15}$ cm$^{-2}$, which corresponds to a factor of $N(^{13}\\mathrm{CO})_{true}/N(^{13}\\mathrm{CO})_{LTE}~\\lesssim~3.5$ in their analysis. The unavoidable assumption of a constant H$_2$ to $^{13}$CO abundance ratio may also introduce error into the results.  CO catalogues such as the one presented here offer an attractive means of calibrating molecular cloud mass conversion factors, in particular the X-factor, through comparison with data from the next generation of gamma ray telescopes such as GLAST. Molecular clouds act as passive targets for hadronic cosmic ray interactions, which result in the production of gamma rays in the MeV to TeV regime. The mass of the target cloud may be calculated directly from the gamma-ray luminosity, and this mass may be compared to CO luminosity without the need for many of the assumptions that plague other calibration methods.  Unfortunately, the estimated masses and angular extents of the catalogue clouds place most of them well below the detectability threshold for GLAST \\citep{torres05}. Nevertheless, several high-latitude objects -- notably clouds 74 and 16 -- fall on the border of predicted range and may eventually prove to be detectable. In the near future these clouds, and others like them, may offer the opportunity to calibrate the X-factor for high altitude clouds away from the immediate solar neighborhood."
    },
    "0808/0808.3582_arXiv.txt": {
        "abstract": "We present the first detection and mapping of the HD 32297 debris disk at 1.3 mm with the Combined Array for Research in Millimeter-wave Astronomy (CARMA). With a sub-arcsecond beam, this detection represents the highest angular resolution (sub)mm debris disk observation made to date. Our model fits to the spectral energy distribution from the CARMA flux and new Spitzer MIPS photometry support the earlier suggestion that at least two, possibly three, distinct grain populations are traced by the current data.  The observed millimeter map shows an asymmetry between the northeast and southwest disk lobes, suggesting large grains may be trapped in resonance with an unseen exoplanet. Alternatively, the observed morphology could result from the recent breakup of a massive planetesimal.  A similar-scale asymmetry is also observed in scattered light but not in the mid-infrared.  This contrast between asymmetry at short and long wavelengths and symmetry at intermediate wavelengths is in qualitative agreement with predictions of resonant debris disk models.  With resolved observations in several bands spanning over three decades in wavelength, HD 32297 provides a unique testbed for theories of grain and planetary dynamics, and could potentially provide strong multi-wavelength evidence for an exoplanetary system. ",
        "introduction": "Debris disks provide the principal means of studying the formation and evolution of planetary systems on timescales of 10$-$100 Myr. Evidence for exoplanets in these systems can be found by matching density variations in debris disks to theoretical models of the gravitational perturbations caused by planets (e.g., Reche et al. 2008).  A modest sample of debris disks have now been imaged in the visible, and many show substructure such as clumps, warps, and offsets, consistent with dynamical perturbations by massive planets. However, a wide variety of other mechanisms can produce similar structures (Moro-Martin et al. 2007, and references therein).  As different wavebands are sensitive to different grain sizes, which are in turn subject to different dynamical influences, multi-wavelength observations offer the most promising path towards definitively classifying the physical mechanisms at work in these systems.  At present, only a few debris disks have resolved observations spanning more than a decade in wavelength.  A particularly critical, though technologically challenging, deficit of observations lies at (sub)millimeter wavelengths, which trace large grains primarily affected by gravitational forces.  To date, bolometer arrays have resolved 8 debris disks in the (sub)mm (e.g., Holland et al. 1998, Greaves et al. 1998). However, such low-resolution ($\\theta_{\\mathrm{beam}} \\gtrsim 10''$) single-dish measurements are limited to the largest, nearest disks.  Higher resolution interferometric observations are needed to access the larger debris disk population already imaged at shorter wavelengths. Some pioneering work has been done in this area; OVRO and PdBI have detected and resolved two debris disks (Vega: Koerner et al. 2001, Wilner et al. 2002; HD 107146: Carpenter et al. 2005).  Recently, \\citet{Corder08} mapped HD 107146 with the Combined Array for Research in Millimeter-wave Astronomy (CARMA) at 1.3 mm, providing the highest fidelity interferometric debris disk map to date.  Here, we report the near-simultaneous CARMA detection of HD 32297, the third debris disk mapped with a (sub)mm interferometer.  With a sub-arcsecond beam, this detection is the highest angular resolution (sub)mm debris disk observation made to date.  HD 32297 is a $\\sim$30 Myr A-star at $112_{-12}^{+15}$ pc \\citep{Perryman97}, first discovered to host a resolved debris disk with HST/NICMOS near-infrared (NIR) imaging \\citep{Schneider05}.  The discovery image showed an edge-on debris disk extending to 400 AU (3.3$''$), with an inner-disk brightness asymmetry inward of 60 AU ($0.5''$).  \\citet{Kalas05} subsequently imaged HD 32297 in the optical, revealing an asymmetric, extended outer disk ($\\sim$1700 AU, 15$''$) likely interacting with the interstellar medium.  Later, \\citet{Redfield07} detected circumstellar gas in this system, reporting the strongest Na I absorption measured toward any known debris disk. Most recently, \\citet{Fitzgerald07}, hereafter F07, and \\citet{Moerchen07} resolved HD 32297 in mid-infrared (MIR), thermal emission.  Detailed analysis of the spectral energy distribution (SED) by F07 showed that multiple grain populations may be present in the disk.  The lack of long wavelength data needed to characterize the large grain properties of HD 32297 motivated the CARMA observations presented here. ",
        "conclusions": "F07 attempted to model the observed SED and $N'$-band image of HD 32297 with a single ring of grains of characteristic size, but found that a second population of grains was needed to adequately fit the observed SED for $\\lambda \\gtrsim 25 \\mu$m.  Indeed, the contrast between the observed mm and $N'$-band morphologies (see below and Figure \\ref{multiwav}) suggests that the grains responsible for the emission at each wavelength constitute separate populations.  To test whether the population of mm-emitting grains is consistent with the second, larger grain population proposed by F07 to fit the SED at 25 $\\mu$m $\\lesssim \\lambda \\lesssim$ 60 $\\mu$m, we revisit the F07 model, adopting their data and fitting method, and incorporating the Qa-band flux of \\citet{Moerchen07} and the new MIPS and CARMA fluxes. The free parameters in the single population model of F07 are the disk inclination ($i$), position angle (PA), radii of the inner and outer edges ($\\varpi_0$, $\\varpi_1$), surface density power-law index ($\\gamma$), vertical optical depth to absorption at the inner edge ($\\tau_0\\equiv\\tau_\\perp^{\\mathrm{abs}}(\\varpi_0)$), stellar flux factor ($\\xi$), and effective grain size ($\\lambda_\\mathrm{sm}$). To fit the long-wavelength SED ($\\lambda \\gtrsim 25 \\mu$m), we augment this model with a population of larger grains of effective size $\\lambda_\\mathrm{lg}$ and total emitting area $A_\\mathrm{lg}$, located in a narrow ring at the small-grain inner-disk edge, $\\varpi_0$. These grains contribute flux, $F_{\\nu,\\mathrm{lg}}$, according to: \\begin{eqnarray} \\epsilon_{\\nu,\\mathrm{lg}} &=&\\left\\{ \\begin{array}{ll} \\lambda_\\mathrm{lg}/\\lambda & \\mbox{if $\\lambda>\\lambda_\\mathrm{lg}$,} \\\\ 1 & \\mbox{otherwise}, \\end{array}\\right. \\\\ T_\\mathrm{lg}(r) &=& 468 \\left(\\frac{L_*/L_\\sun}{\\lambda_\\mathrm{lg}/1\\,\\micron}\\right)^{1/5} \\left(\\frac{r}{1\\,\\mathrm{AU}}\\right)^{-2/5} \\mathrm{K}, \\\\ F_{\\nu,\\mathrm{lg}} &=& \\left(\\frac{A_\\mathrm{lg}}{d^2}\\right) \\epsilon_{\\nu,\\mathrm{lg}} B_\\nu[T_\\mathrm{lg}(\\varpi_0)]. \\end{eqnarray}  Following F07, we ran three Monte Carlo Markov chains with $3\\times 10^4$ samples, and simulataneously fit both the large and small grain populations.  The results of this procedure are listed in Table \\ref{sed_table}, and the corresponding range of allowed dust emission is plotted in Figure \\ref{sed} (red: small grains, blue: large grains, grey: photosphere, green: composite).  While the two-population model provides a satisfactory fit to the mid- and far-infrared data, the mm flux is underestimated at the 4$\\sigma$ level, suggesting that the second grain population proposed by F07 is not responsible for the majority of the mm flux.  Nevertheless, the new MIPS data lend further evidence to the F07 suggestion that two populations are needed to fit the observed SED for $\\lambda \\lesssim 160 \\mu$m.  Therefore, the fit suggests at least three distinct populations are traced by the current observations.  However, from the two-population model, only the 1.3 mm flux appears to trace the putative third population.  Since this population is described by both a mass/emitting area and a size/temperature, we lack sufficient data to fully characterize it via modeling. We note, though, that this population likely traces $\\gtrsim 95$\\% of the total dust mass.  Assuming a characteristic stellocentric distance of $\\sim$50 AU (\\S 3), $L_{*}=5.4L_{\\sun}$ (F07), and an effective grain size of 1.3-mm, the implied mm grain temperature is $\\sim$30 K, suggesting a dust mass of $M_{\\mathrm{mm}}\\sim M_{\\earth}$ (adopting an opacity of 1.7 cm$^2$ g$^{-1}$).  This estimated mass is among the highest observed for debris disks detected in the (sub)mm and is two orders of magnitude larger than that implied for the ``large-grain'' population in the SED fit: $M_{\\mathrm{lg}}\\sim0.02M_{\\earth}$ (using the fitted parameters in Table \\ref{sed_table} and assuming spherical grains with a density of 1 g cm$^{-3}$).  Future far-infrared / sub-mm observations are needed to confirm the three populations proposed here and better constrain their properties.  \\begin{deluxetable}{lcc} \\tablecaption{\\label{sed_table} Best-Fit Model Parameters} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter}  & \\colhead{Best-fit} & \\colhead{Description} } \\startdata $i$ (deg) & $90  \\pm 5$ & disk inclination \\\\ PA (deg) & $46 \\pm 3$ & disk position angle \\\\ $\\varpi_0$ (AU) & $70_{-10}^{+20}$ & inner edge \\\\ $\\varpi_1$ (AU) & $>$ 1200 & outer edge \\\\ $\\log_{10}\\tau_0$ & $-2.4_{-0.2}^{+0.3}$ & vertical optical depth at $\\varpi_0$ \\\\ $\\gamma$ & $<$ -1.63 & surf density power-law index \\\\ $\\xi / 7.22 \\times 10^{-20}$ & $1.03 \\pm 0.03$ & $(R_{*}/d)^2$, stellar flux factor \\\\ $\\log_{10}(\\lambda_\\mathrm{sm}/1\\,\\micron)$ & $-1.4_{-1.3}^{+0.6}$ & small grain effective size \\\\ $\\log_{10}(A_\\mathrm{lg}/1\\,\\mathrm{cm}^2)$ & $28.9 \\pm 0.2$ & large grain emitting area \\\\ $\\log_{10}(\\lambda_\\mathrm{lg}/1\\,\\micron)$ & $1.3_{-0.1}^{+0.2}$ & large grain effective size \\enddata \\tablecomments{Confidence intervals are 95\\% for marginal posterior distributions.  Adopted priors are as in F07, except for $\\gamma$, which is constrained to be in the domain [-4,0] due to the inconsistency of rising surface density with the scattered-light image and CARMA map; the large grain effective size was constrained by a log-uniform prior from 1\\,nm to 100\\,mm and a requirement that $\\lambda_{\\mathrm{lg}} > \\lambda_{\\mathrm{sm}}$.} \\end{deluxetable} \\begin{figure} \\centering \\includegraphics[width=0.23\\textwidth,angle=90]{f4.eps} \\caption{{\\it Left:} CARMA contours of HD 32297, overlaid on the NIR scattered-light image from Schneider et al. (2005). {\\it Right:} Photosphere-subtracted MIR contours from F07 overlaid on the same image. The asymmetry in the CARMA data between the northeast and southwest lobes suggests the large, mm-sized grains may be trapped in resonance with an unseen exoplanet.  A similar asymmetry is also observed in scattered light but not in the MIR.  The contrast between asymmetry at short and long wavelengths and symmetry at intermediate wavelengths is a direct prediction of the resonant debris disk models of Wyatt (2006).} \\label{multiwav} \\end{figure}  To address the observed mm morphology, in Figure \\ref{multiwav}, we qualitatively compare the CARMA mm map (contours, left panel) to the MIR image of F07 (contours, right panel) and the NIR scattered-light image of \\citet{Schneider05} (color, both panels).  The images at all three wavelengths are consistent with an edge-on disk.  However, while both the NIR and CARMA data exhibit a brightness asymmetry between the northeast and southwest lobes inward of $\\sim$0.5$''$, the F07 MIR image is consistent with azimuthal symmetry.  We note that \\citet{Moerchen07} found evidence for asymmetry in their Qa-band data, though their data were not PSF subtracted, and the asymmetry was in the opposite sense as observed in the NIR / CARMA data (NE lobe brighter than SW).  Recently, \\citet{Grigorieva07} used predictions from their numerical model of collisional avalanches to suggest that the observed scattered-light asymmetry in HD 32297 results from the breakup of a large planetesimal.  However, while a massive collision can explain the observed NIR and mm morphology, it is not clear why the MIR image would not show a similar asymmetry.  An alternative hypothesis is that the structure results from a planetary-induced resonance.  In this case, the dust is either trapped in resonance as it drifts inward due to Poynting-Robertson drag, or it remains locked in resonance after being generated by parent planetesimals in resonance themselves (e.g., Krivov et al. 2007, and references therein).  Interestingly, a recent study of the latter mechanism by \\citet{Wyatt06} directly predicts a contrast between asymmetry at short and long wavelengths and symmetry at intermediate wavelengths, as observed in HD 32297.  In this model, (sub)mm emission is dominated by large grains, which have the same clumpy resonant distribution as the parent planetesimals.  Small grains, traced at short wavelengths, exhibit a similar asymmetry, as they are preferentially born in the high density, resonant structures before being rapidly expelled from the system.  Lastly, moderately-sized grains sampled at intermediate wavelengths remain bound to the star, but have fallen out of resonance due to radiation pressure and are subsequently scattered into an axisymmetric morphology.  This proposed scenario of \\citet{Wyatt06} conveniently explains the qualitative picture for HD 32297 depicted in Figure \\ref{multiwav}, yet we caution that this suggestion is highly speculative, and future rigorous modeling of this system is needed to draw firm conclusions from the available data.  One potential problem with this hypothesis is that the SED model predicts that the small ($N'$-band-emitting) grains have sizes $\\lesssim$ 1 $\\mu$m (Table \\ref{sed_table}), similar to that expected to produce the NIR scattered-light image (see discussion in F07).  However, the interpretation of Figure \\ref{multiwav} in terms of the \\citet{Wyatt06} models requires that the NIR-scattering and MIR-emitting grains have different sizes.  This ambiguity could potentially be resolved through simultaneous modeling of scattering and emission, incorporating the optical image presented in \\citet{Kalas05}. Most importantly, and independent of speculation, HD 32297 is currently one of only a few debris disks with resolved observations in four wavelength regimes (optical, NIR, MIR, mm). Taken together, these observations can provide a unique testbed for theories of grain and planetary dynamics."
    },
    "0808/0808.1855_arXiv.txt": {
        "abstract": "The microscopic quantum field theory origins of warm inflation dynamics are reviewed.  The warm inflation scenario is first described along with its results, predictions and comparison with the standard cold inflation scenario.  The basics of thermal field theory required in the study of warm inflation are discussed.  Quantum field theory real time calculations at finite temperature are then presented and the derivation of dissipation and stochastic fluctuations are shown from a general perspective.  Specific results are given of dissipation coefficients for a variety of quantum field theory interaction structures relevant to warm inflation, in a form that can readily be used by model builders.  Different particle physics models realising warm inflation are presented along with their observational predictions. ",
        "introduction": "It has been fourteen years since warm inflation was introduced and with it the first and still only alternative dynamical realisation of inflation to the standard scenario.  The standard picture of inflation introduced in 1981 relied on a scalar field, called the inflaton, which during inflation was assumed to have no interaction with any other fields.  During inflation, this field rolls down its potential and due to it being coupled to the background metric, a damping-like term is present which slows down its motion.  As this inflaton field was assumed to not interact with other fields, there was no possibility for radiation to be produced during inflation, thus leading to a thermodynamically supercooled phase of the Universe during inflation.  Getting out of this inflation phase and putting the Universe into a radiation dominated phase was a key issue, termed the ``graceful exit\" problem \\cite{oldinf,ni,reheatu}.  The first successful solution of the graceful exit problem and so the first successful cold inflation model, was new inflation \\cite{ni}.  The solution was to picture particle production as a distinct separate stage after inflationary expansion in a period called reheating.  In this phase, couplings to other fields were assumed to be present and the inflaton would find itself in a very steep potential well in which it would oscillate.  These oscillations would lead to a radiative production of particles.  There have been many variants of the original cold inflation picture, first introduced in the context of the new inflation model and shortly afterwards in the chaotic inflation model~\\cite{ci}, with many other models that followed.  The warm inflation picture differs from the cold inflation picture in that there is no separate reheating phase in the former, and rather radiation production occurs concurrently with inflationary expansion.  The constraints by General Relativity for realising an inflationary phase simply require that the vacuum energy density dominates and so this does not rule out the possibility that there is still a substantial radiation energy density present during inflation.  Thus on basic principles, the most general picture of inflation accommodates a radiation energy density component.  The presence of radiation during inflation implies the inflationary phase could smoothly end into a radiation dominated phase without a distinctively separate reheating phase, by the simple process of the vacuum energy falling faster than the radiation energy, so that at some point a smooth crossover occurs.  This is the warm inflation solution to the graceful exit problem.  Dynamically warm inflation is realised if the inflaton were interacting with other fields during the inflation phase.  In fact in any realistic model of inflation, the inflaton must be coupled to other fields, since eventually the inflaton must release its vacuum energy to other fields thereby creating particles which form the subsequent radiation dominated era in the Universe. Thus the idea that these couplings to other fields somehow are inactive during inflation, as pictured in the cold inflation picture, is something that does require verifying by detailed calculation.  When such calculations are done, the result is that there are regimes in which particle production during inflation occurs.  This review will present the calculations which demonstrate particle production during inflation, thereby leading to a warm inflationary expansion.  The idea of particle production concurrent with inflationary expansion was first suggested in the pre-inflation inflation paper by L.Z. Fang in 1980 \\cite{Fang:1980wi}.  His paper proposed using a scalar field with the origin of inflationary expansion due to a claimed anomalous dissipation term that would be generated based on Landau theory if the field was undergoing a second order phase transition.  This was dynamically very different from the scalar field inflation that eventually became successful, and the source of dissipation was also different from that in warm inflation.  However this model captured the basic idea of concurrent particle production and inflation. Then in the mid-80s two papers proposed adding a local $\\Upsilon {\\dot \\phi}$ type dissipation term into the evolution equation of the inflaton field, Moss \\cite{im} and then Yokoyama and Maeda~\\cite{Yokoyama:1987an}.  In both cases the dissipative term generated a source of radiation production during inflation.  The idea of a dissipative term was re-discovered independently by Berera and Fang~\\cite{Berera:1995wh} almost a decade later.  They went further by proposing that the consistent dynamics of the inflaton field was a Langevin equation, in which a fluctuation-dissipation theorem would uniquely specify the fluctuations of the inflaton field.  That paper by Berera and Fang provided the foundations for the theory of fluctuations in warm inflation and the Langevin equation has since been the fundamental equation governing inflaton dynamics.  {}Following that work, in~\\cite{Berera:1995ie} Berera proposed that a separate reheating phase, as standard in all inflation models up to then, could be eliminated altogether.  This paper proposed a new picture of inflation, which it termed warm inflation, in which the process of inflationary expansion with concurrent radiation production could terminate simply by the radiation energy over-taking the vacuum energy, thus going from an inflationary to a radiation dominated era.  This work presented an alternative solution to the graceful exit problem to the one given by the standard inflation scenario.  This warm inflation picture was verified explicitly in \\cite{Berera:1996fm}, where the {}Friedmann equations for a Universe consisting of vacuum and radiation energy were studied and gave many exact warm inflation solutions to the graceful exit problem. {}Finally in~\\cite{Berera:1999ws} the calculation of fluctuations were done by Berera for the inflaton evolving by a Langevin equation in a thermal inflationary Universe.  Alongside the development of the basic scenario, the first principles quantum field theory dynamics of warm inflation was developed.  This started in \\cite{Berera:1996nv} with a quantum mechanical model which demonstrated the origin of the fluctuation-dissipation relation in warm inflation.  The key step in deriving warm inflation from quantum field theory is in realising an overdamped regime for the evolution of the background inflaton field.  The initial attempt at this was done by Berera, Gleiser and Ramos in~\\cite{BGR}. In this work, it was proposed that the overdamped evolution should occur under adiabatic conditions in which the microscopic dynamical processes operated much faster than all macroscopic evolution, in particular the scalar field motion and Hubble expansion.  Based on this criteria, a set of consistency conditions were formulated in \\cite{BGR}, which would be required for a self-consistent solution.  However in~\\cite{BGR} no explicit warm inflation solutions were found.  This point was further highlighted by Yokoyama and Linde in~\\cite{Yokoyama:1998ju}, in which several models were studied from which the conclusions of~\\cite{BGR} were verified.  The problem in these early works was that dissipation effects were being looked at in a high temperature regime and it proved too difficult to keep finite temperature effective potential corrections small, so that the inflaton potential remained relatively flat, and at the same time obtain a large dissipative coefficient. One type of model was shown that could realise such requirements~\\cite{BGR2} and this was the first quantum field theory model of warm inflation, although it was not a very compelling model.  Subsequently Berera and Ramos in \\cite{BR1} suggested a solution for getting around the mutual constraints of obtaining a large dissipative coefficient and yet small effective potential corrections.  The main observation was that supersymmetry can cancel local quantum corrections, such as zero temperature corrections to the effective potential, whereas temporally non-local quantum effects, such as those that underly the dissipative effects, will not be cancelled.  This led to \\cite{BR1} proposing a two-stage interaction configuration, in which the inflaton was coupled to heavy ``catalyst'' fields with masses larger than the temperature of the Universe and these fields in turn were coupled to light fields.  The evolution of the inflaton would induce light particle production via the heavy catalyst fields.  Since these heavy catalyst fields were basically in their ground state, the quantum corrections associated with them could be cancelled in supersymmetric models.  The calculation of the low temperature dissipative coefficients for this two-stage mechanism were first done by Moss and Xiong \\cite{mx}.  There is an earlier review which covered the basics of the warm inflation scenario \\cite{Berera:2006xq}.  In this review full details will be developed of the quantum field theory dynamics of warm inflation.  This will first start in Sec.  \\ref{wipicture} with a summary of the warm inflation scenario, including a comparison of it to cold inflation.  In Sec. \\ref{sec TFT} a basic introduction to thermal field theory is given including the real time formalism for interacting field theories.  In Sec. \\ref{effeom} the effective evolution equation of the inflaton field is derived, in which all fields it interacts with are integrated out, leading to a Langevin type nonconservative equation which contains a dissipative term and a noise force term.  The detailed properties of these dissipation and fluctuation terms is then studied in Sec.  \\ref{fd}.  In addition the physical picture of the dissipation effects in warm inflation are discussed.  In Sec. \\ref{FRWspacetime} the calculations are extended to curved spacetime. In Sec. \\ref{particlemod} various particle physics models of warm inflation are presented.  Finally Sec. \\ref{concl} presents some concluding remarks and future work being done on the subject. Our conventions are as follows. We use spacetime metrics with $(+---)$ signature, and we use natural units for the Planck's constant, Boltzmann's constant and the velocity of light $\\hbar=k=c=1$. ",
        "conclusions": "\\label{concl}   Generically, the inflaton interacts with other fields in any typical inflation model and so its dynamics is dissipative.  As such, inflation, like most dynamics in nature is an open system phenomenon, thus requiring a much more complex analysis of its dynamics than the one typically formulated for the cold inflation picture.  This general point has been voiced by B. L. Hu and coworkers ~\\cite{hu1} and more specifically in the initial motivating papers of warm inflation \\cite{Berera:1995wh,Berera:1995ie,Berera:1996nv,BGR}.  In all these cases, the point has been made that the problem of inflation encompasses many different branches of Physics, from nonequilibrium statistical dynamics to particle physics phenomenology.  Warm inflation dynamics is a rich area of study for both quantum field theory real time dynamics and for particle physics model building.  The study of the warm inflation dynamics has almost exclusively motivated the understanding of strong dissipative behavior in quantum field theory.  This started initially with the work of Berera, Gleiser and Ramos~\\cite{BGR}, in which dissipation was examined in the overdamped regime for the first time in quantum field theory using extended linear response calculations.  Since then, specific progress has been made to underpin interaction structures in interacting quantum field theory models which lead to strong dissipative behavior under warm inflationary conditions~\\cite{BR1,BRplb1,BRfrw,BRplb2}.  At a more general level, this work has motivated the first calculations of dissipation in the low temperature regime~\\cite{mx}.  In this review, a physical picture has also been developed in Sec.~\\ref{fd}, for explaining the dissipation behavior found in warm inflation from the quantum field theory calculations, and this is further developed in~\\cite{grahammoss}.  Up to now all these studies has been based on variants of linear response methods, in which case the dynamics is supposed to happen close to equilibrium, or in a quasi-adiabatic regime, including various types of resummations.  Several extensions to this work, which will give a more accurate understanding of dissipation, are under way. {}For instance, when departing from the quasi-adiabatic regime for the field dynamics, it is expected that the local Markovian approximation commonly used to analyse the field equations can differ significantly from the exact nonlocal equations. Preliminary tests have shown that this difference can be very large for a period of time starting from the initial period, but with differences between the dynamics getting smaller at longer times, depending on the model parameters~\\cite{FRS}.  The methods used in this review are all based on the effective equation of motion derived from the action functional.  In order to study strong nonequilibrium dynamics, it requires methods beyond these.  In this case, full kinetic set of equations for the relevant fields have to be studied, which contain information not only of the inflaton effective dynamics but also about thermalisation and equilibration. Work in this direction, also can help to better understand the physics of particle production during the system dynamics and be used to check the reliability of using the approximation of thermal initial conditions for the bath fields.  Initial work in this direction, though not directly related to the type of models relevant to warm inflation as discussed extensively e.g. in Sec.~\\ref{particlemod}, has appeared~\\cite{Aarts:2007ye}, while a work more on the warm inflation dynamics motivated side is under way~\\cite{BMR}.  This review has presented in detail the methods used so far to understand warm inflation dynamics as well as the limitations of these studies.  In particular these methods used so far are primarily quasi-adiabatic approximations for the fields with the assumption of near thermal equilibrium evolution.  Despite these limitations, these results find parameter regimes in which physically acceptable solutions exist over time periods sufficiently long to be of use in studies of warm inflation. During this time interval in which these approximations apply, and where a local Markovian dynamics can be used, in contrast to the full non Markovian one, in the typical model implementations discussed in Sec.~\\ref{particlemod}, the amount of radiation production was seen to be sufficient to change the cold inflationary picture predictions regarding the density perturbations, thus requiring its description in terms of the warm inflation picture. Moreover, once dissipative effects are strong enough, inflation can be sustained and driven longer than when these effects are neglected.  Consequently, parameter values typically required in the cold inflation case can be relaxed, which can help to evade various problems that plaque the standard scenario of inflation, like the graceful exit and $\\eta$-problem, discussed in details in this review, as well as the problems of quantum-to-classical transition~\\cite{Berera:1995ie,bel} and the initial conditions for inflation~\\cite{BG,ROR}.  The development of the quantum field theory dynamics of warm inflation has in turn been applied to particle physics model building, in which warm inflation dynamics in realised.  Early on it was recognisied in \\cite{Berera:1999ws} that warm inflation has some appealing and unique model building features.  In particular, it offers a simple solution to the $\\eta$-problem and for monomial potentials, observationally consistent inflation can occur for the inflaton amplitude below the Planck scale $\\langle \\phi \\rangle < m_P$.  These features have been realised in explicit first principles quantum field theory models of warm inflation in \\cite{BGR2,Berera:1999ws,bb4,BuenoSanchez:2008nc}. These studies are now being extended to develop a complete particle cosmology in which not only is warm inflation realised, but in addition other features such as leptogenesis and gravitino abundances are addressed, developing in depth some of the work already started on these topics \\cite{Taylor:2000jw,bb1,bb4,Lambiase:2006md,BuenoSanchez:2008nc}.  In a separate direction, the warm inflation models developed so far in \\cite{bb4,BuenoSanchez:2008nc} have been in the low temperature regime.  In \\cite{bbrecent} this is being extended to higher temperatures.  This requires calculating all the thermal loop corrections in the SUSY models that have the two-stage interaction structure Eq. (\\ref{wi2stage}) relevant for warm inflation.  Some initial work has been done in \\cite{Hall:2004zr}, and a more detailed analysis is now underway \\cite{bbrecent}."
    },
    "0808/0808.0455_arXiv.txt": {
        "abstract": "We present new \\XMM EPIC observations of the nuclei of the nearby radio galaxies 3C\\,305, DA\\,240, and 4C\\,73.08, and investigate the origin of their nuclear X-ray emission. The nuclei of the three sources appear to have different relative contributions of accretion- and jet-related X-ray emission, as expected based on earlier work. The X-ray spectrum of the FRII narrow-line radio galaxy (NLRG) 4C\\,73.08 is modeled with the sum of a heavily absorbed power law that we interpret to be associated with a luminous accretion disk and circumnuclear obscuring structure, and an unabsorbed power law that originates in an unresolved jet. This behavior is consistent with other narrow-line radio galaxies. The X-ray emission of the low-excitation FRII radio galaxy DA\\,240 is best modeled as an unabsorbed power law that we associate with a parsec-scale jet, similar to other low-excitation sources that we have studied previously. However, the X-ray nucleus of the narrow-line radio galaxy 3C\\,305 shows no evidence for the heavily absorbed X-ray emission that has been found in other NLRGs. It is possible that the nuclear optical spectrum in 3C\\,305 is intrinsically weak-lined, with the strong emission arising from extended regions that indicate the presence of jet--environment interactions. Our observations of 3C\\,305 suggest that this source is more closely related to other weak-lined radio galaxies. This ambiguity could extend to other sources currently classified as NLRGs. We also present \\XMM and VLA observations of the hotspot of DA\\,240, arguing that this is another detection of X-ray synchrotron emission from a low-luminosity hotspot. ",
        "introduction": "\\label{intro}  Radio galaxies consist of twin jets of particles that are ejected from a compact region in the vicinity of a supermassive black hole, feeding into large-scale `plumes' or `lobes'. There are two principal morphological classes of radio galaxies, low-power (Fanaroff-Riley type I, hereafter FRI) sources and high-power (FRII) sources \\cite{fr74}. FRI sources exhibit `edge-darkened' large-scale radio structure, and modeling implies that initially supersonic jets in these sources decelerate to transonic speeds on $\\sim$kpc scales before flaring into large plumes (e.g., \\citealt{per07}). FRII sources appear `edge-brightened', and in these cases highly supersonic jets propagate out to large distances (often $>100$ kpc) from the core before terminating in bright hotspots and accompanying radio lobes. Observationally, the Fanaroff-Riley divide occurs at a 178-MHz radio power of $\\sim$$10^{25}$ W~Hz$^{-1}$~sr$^{-1}$.  It is important to understand whether the kpc-scale Fanaroff-Riley dichotomy is determined by the interaction between the jet and its external hot-gas environment (e.g., \\citealt{bic95}), or rather is nuclear in origin and governed by differences in the properties of the accretion flow (\\citealt{rey96}). The first observations of large samples of $z<0.1$ 3CRR radio-galaxy {\\it nuclei} with \\Ch and \\XMM (\\citealt{don04,bal06,evans06}) showed that FRI nuclei show no signs of heavily absorbed X-ray emission that would be expected from standard AGN unification models (\\citealt{up95}), are dominated by emission from an unresolved jet (e.g.,~\\citealt{wb94,bal06,evans06}), and have highly radiatively inefficient accretion flows. Narrow optical-line FRII sources show evidence for heavily obscured ($N_{\\rm H}>10^{23}$~cm$^{-2}$) nuclear X-ray emission that is associated with a radiatively efficient accretion flow, together with an unabsorbed component of jet-related emission (\\citealt{evans06,hec06}). FRII radio galaxies at higher redshift are consistent with such behavior (\\citealt{bel06}).  A significant breakthrough for understanding the physical origin of the FRI/FRII dichotomy came from \\Ch and \\XMM observations of the population of {\\it low-excitation radio galaxies} (LERGs), which have weak or no emission lines in their optical spectra (\\citealt{hine79,jac97}). Almost all FRI radio galaxies are LERGs, but there is a significant population of FRII LERGs at $0.1 < z < 0.5$. The X-ray spectra of LERGs, irrespective of their FRI or FRII morphology, are dominated by unabsorbed emission that can be associated with a parsec-scale jet, with no obvious contribution from accretion-related emission. These sources are likely to accrete in a radiatively inefficient manner (\\citealt{hec06}). On the other hand, high-excitation radio galaxies (HERGs -- i.e., NLRGs, BLRGs, and quasars), which display prominent narrow or broad optical emission lines, have X-ray spectra that are consistent with standard unification models: they show evidence for luminous, radiatively efficient accretion disks, together with circumnuclear tori when the source is oriented close to edge-on with respect to the observer. HERGs tend to show evidence for additional hot dust over and above that of LERGs in their mid-IR spectra (e.g., \\citealt{ogle06}; Birkinshaw et al., in preparation), which is again consistent with reprocessing of luminous accretion-related emission by torus-like structure. Most high radio-power (FRII) sources are high-excitation radio galaxies. The distinct X-ray nuclear properties of low- and high-excitation radio galaxies, regardless of their large-scale FRI or FRII morphology, could be interpreted as implying that the Fanaroff-Riley dichotomy remains principally influenced by jet power and environment. The excitation dichotomy, on the other hand, is interpreted to be attributed to the radiative efficiency of the accretion flow (e.g., \\citealt{hec06}) and possibly related to the nature of the accreting material (\\citealt{hec07}).   Here, we report new \\XMM observations of the nuclei of three $z<0.1$ 3CRR radio galaxies --- 3C\\,305, DA\\,240, and 4C\\,73.08. The three sources have 178-MHz radio powers that lie close to the FRI/FRII dividing luminosity (Table~\\ref{sourcessummary}), plus a range of radio morphologies and optical emission-line characteristics. They are therefore good candidates for examining possible connections between the central engine and large-scale radio characteristics. This paper is organized as follows. In Section 2, we describe the optical and radio properties of the three sources. Section 3 contains a description of the data and a summary of our analysis. In Section 4, we report the results of our spectroscopic analysis of the sources. In Section 5, we describe VLA and \\XMM observations of the bright NE hotspot in DA\\,240. In Section 6, we interpret the observations in the context of our previous \\Ch and \\XMM observations of 3CRR radio galaxies and discuss the optical emission-line characteristics of the sources. We end with our conclusions in Section 7. All results presented in this paper use a cosmology in which $\\Omega_{\\rm m, 0}$ = 0.3, $\\Omega_{\\rm \\Lambda, 0}$ = 0.7, and H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$. Errors quoted in this paper are 90 per cent confidence for one parameter of interest (i.e., $\\chi^2_{\\rm min}$ + 2.7), unless otherwise stated. ",
        "conclusions": "We have presented results from \\XMM observations of the nuclei of the radio galaxies 3C\\,305, DA\\,240, and 4C\\,73.08. We have shown the following:  \\begin{enumerate} \\item The X-ray spectrum of the narrow-line FRII radio galaxy 4C\\,73.08 can be modeled as the sum of a heavily absorbed power law associated with a luminous accretion disk and circumnuclear obscuring structure, together with an unabsorbed component of X-ray emission that has a common origin with the radio emission at the base of an unresolved jet. This behavior is consistent with the other narrow-line FRII radio galaxies studied by \\cite{evans06} and \\cite{hec06}. \\item The nuclear X-ray spectrum of the FRII giant radio galaxy DA\\,240, optically classified as a low-excitation radio galaxy, can be modeled as a single, unabsorbed power law that is likely associated with emission from the parsec-scale jet. The upper limit to the X-ray luminosity of any additional, accretion-related emission suggests that the accretion process in DA\\,240 is substantially sub-Eddington and likely radiatively inefficient in nature. \\item The X-ray emission in the nucleus of the narrow-line radio galaxy 3C\\,305 can be modeled as an unabsorbed power law that originates at the base of the jet. However, it shows no evidence for heavily absorbed X-ray emission was found in the NLRGs studied by \\cite{evans06}. \\item We have discovered an X-ray counterpart to the NE hotspot of the giant radio galaxy DA\\,240. We argue that the emission process is overwhelmingly likely to be synchrotron emission. Because of the high X-ray flux of the hotspot, it is a good candidate for followup high-resolution X-ray observations. \\item We have discussed the different origins of optical emission lines in the nuclear and circumnuclear gaseous environments of radio galaxies. These include photoionization from the AGN accretion flow or parsec-scale jet, shock-excitation by the radio jet, or cooling gas in the centers of clusters. This may lead to occasional misclassification of genuinely weak-lined sources such as 3C\\,305 as high-excitation sources. \\item We therefore argue that there is not necessarily always a one-to-one correspondence between optical emission-line class (low- vs. high-excitation) and accretion-flow state (inefficient flow vs. standard thin disk), especially when low angular-resolution optical spectroscopy is used. We suggest that only the combination of high-resolution optical, X-ray, and infrared observations can reliably uncover the nature of the central engine in radio-loud AGN. \\end{enumerate}"
    },
    "0808/0808.0349_arXiv.txt": {
        "abstract": "s{ A striking concentration of ultra-high energy cosmic ray (UHECR) events observed by the Pierre Auger Observatory around the direction of the nearby radio galaxy Centaurus A revives the idea that radio galaxies may be dominant sources of UHECR. In this paper, we give a brief overview about processes which may accelerate protons and nuclei in radio galaxies, and their relation to jet power, radio morphology and cosmic source density. We argue that, except for the most powerful FR-II radio galaxies, processes in radio lobes are unlikely to explain the origin of UHECR. However, Fermi acceleration of protons at internal shocks in the ``blazar-zone'' of \\textit{all} radio galaxies, and their photohadronic conversion into neutrons, may lead to the ejection of ``UHECR-beams'', which remain collimated over several Mpc. Consequences of this hypothesis for the interpretation of the UHECR event distribution, in particular for the special case of Centaurus A, are discussed.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.2385_arXiv.txt": {
        "abstract": "We present an analysis of six 12 ${\\umu}$m selected Seyfert 2 galaxies that have been reported to be unabsorbed in the X-ray.  By comparing the luminosities of these galaxies in the mid-IR (12 ${\\umu}$m), optical ([O\\,{\\sc iii}]) and hard X-ray (2-10 keV), we show that they are all under-luminous in the 2-10 keV X-ray band. Four of the objects exhibit X-ray spectra indicative of a hard excess, consistent with a heavily obscured X-ray component and hence a hidden nucleus.  In these objects the softer X-rays may be dominated by a strong soft scattered continuum or contamination from the host galaxy, which is responsible for the unabsorbed X-ray spectra observed, and accounts for the anomalously low 2-10 keV X-ray luminosity. We confirm this assertion in NGC 4501 with a {\\it Chandra} observation, which shows hard X-ray emission coincident with the nucleus, consistent with heavy absorption, and a number of contaminating softer sources which account for the bulk of the softer emission. We point out that such ``Compton thick\" sources need not necessarily present iron K$\\alpha$ emission of high equivalent width. An example in our sample is IRASF 01475-0740, which we know must host an obscured AGN as it hosts a hidden broad line region seen in scattered light \\citep{tran03}. The X-ray spectrum is nonetheless relatively unobscured and the iron K$\\alpha$ line only moderate in strength ($\\sim 160$ eV). These observations can be reconciled if the hidden nuclear emission is dominated by transmitted, rather than reflected X-rays, which can then be weak compared to the soft scattered light or galactic emission even at 6.4 keV.  Despite these considerations, we conclude that two sources, NGC 3147 and NGC 3660, may intrinsically lack a broad line region (BLR), confirming the recent results of \\cite{bianchi08} in the case of NGC 3147.  Neither X-ray spectrum shows signs of hidden hard emission  and both sources exhibit X-ray variability leading us to believe we are viewing the nucleus directly. ",
        "introduction": "The current model of AGN unification (e.g. \\cite{antonucci93}) states that the absence of broad lines in the optical spectra of Seyfert 2 galaxies is due to their obscuration by an optically thick structure.  Following this, the obscuration that blocks the BLR from view should also absorb soft X-rays from the continuum source in the same region as the BLR.  This is well represented and documented in the prototypical Seyfert 2, NGC 1068, which shows an X-ray reflection spectrum and strong iron K$\\alpha$ emission \\citep{iwasawa97}, typical of Compton thick obscuration of the X-ray source, as well as displaying polarised broad lines (PBLs) in its optical spectrum, indicative of a hidden BLR \\citep{antonucci85}. The X-ray line of sight column density ($N_{\\rm{H}}$) has been measured to be in excess of $10^{23}$ cm$^{-2}$ in many Seyfert 2s \\citep{awaki91}, greatly exceeding the typical values in Seyfert 1s and lending strong support to the unification model. However, there remains an enigmatic group of Seyfert 2s that appear unabsorbed in X-rays (e.g. \\cite{pappa01}; \\cite{panessa02}), whose properties therefore seem to  contradict the unified scheme. This apparent mismatch may result from a genuine absence of the BLR. This intriguing possibility has been proposed to pertain in Seyfert 2 galaxies with particularly low accretion rates, with concomitant implications for the origin of the broad optical lines \\citep{nicastro03}. Apparent confirmation of such a BLR-free AGN has recently been presented by Bianchi et al. (2008) who considered the case of NGC 3147. They were able to reject an alternative interpretation that the mismatch between X-ray and optical properties was due to variability. They also argued against a hidden nucleus based on the {\\it XMM-Newton} spectrum of NGC 3147, which shows a relatively modest iron K$\\alpha$ emission line. Most objects with very heavily obscured nuclei show very strong iron K$\\alpha$ lines \\citep{turner97} with the most extreme examples being the Compton thick sources, with lines often in excess of 1 keV equivalent width \\citep{matt96_2}.  On the other hand, it is conceivable that in these unabsorbed sources we are simply being mislead by the X-ray data. The emission from the nucleus could be contaminated by diffuse or extra-nuclear point-source emission from the host galaxy or by AGN X-rays scattered by hot electrons.  It is therefore possible that there are many Seyfert 2s which from their X-ray spectra, appear to be unabsorbed, but in fact are hiding a deeply buried AGN and that the unabsorbed profile is not necessarily nuclear. The observed equivalent width of the iron K$\\alpha$ line is strongly dependent on the relative contributions of the various components and, in addition, the geometry and physical parameters of the obscuring material.  In this paper we present an analysis of six apparently unabsorbed Seyfert 2 galaxies, aiming to probe the nature of these objects.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Positional and observational data on the six Seyfert 2 galaxies featured.  Spectropolarimetry data are taken from Tran (2003)  and positional, redshift and distance data are taken from NED.}\\label{obsdat} \\begin{tabular}{|l l l l l l l c l l l} \\hline Object name \t& RA \t\t& Dec \t\t& z \t& Distance \t& $N_{\\rm{H}}$ \t\t& PBLs? & $\\frac{\\rm{[N\\,II]}}{\\rm{H\\alpha}}$ & $\\frac{\\rm{[O\\,III]}}{\\rm{H\\beta}}$ & Ref. \\\\ (1) \t\t& (2) \t\t& (3) \t\t& (4) \t& (5) \t\t& (6) \t\t\t& (7) \t& (8) \t& (9) \t& (10)\\\\ \\hline IRASF 01475-0740& 01h50m02.7s \t& -07d25m48s \t& 0.0177& 69.7 \t\t& $2.0 \\times 10^{20}$ \t& yes \t& 0.49 \t& 5.21 \t& 2 \\\\ NGC 3147 \t& 10h16m53.6s \t& +73d24m03s \t& 0.0094& 39.8 \t\t& $2.9 \\times 10^{20}$ \t& no  \t& 2.71 \t& 6.11 \t& 1 \\\\ NGC 3486 \t& 11h00m23.9s \t& +28d58m29s \t& 0.0023& 13.5\t\t& $1.9 \\times 10^{20}$ \t& -  \t& 1.05 \t& 4.57 \t& 1 \\\\ NGC 3660 \t& 11h23m32.3s \t& -08d39m31s \t& 0.0123& 56.2\t\t& $3.4 \\times 10^{20}$ \t& no \t& 0.82 \t& 2.63 \t& 3 \\\\ NGC 3976\t& 11h55m57.6s \t& +06d45m03s \t& 0.0083& 39.4\t\t& $1.1 \\times 10^{20}$ \t& -  \t& 1.96 \t& 3.52 \t& 1 \\\\ NGC 4501\t& 12h31m59.2s \t& +14d25m14s \t& 0.0076& 36.0\t\t& $2.6 \\times 10^{20}$ \t& no \t& 2.10 \t& 5.25 \t& 1 \\\\ \\hline \\end{tabular} Col. (1) Galaxy name as given by \\cite{rush93} in the extended 12 micron galaxy sample; Col. (2) Right ascension (NED, J2000); Col. (3) Declination (NED, J2000); Col. (4) Redshift (NED); Col. (5) Luminosity distance (NED, Mpc); Col. (6) Galactic $N_{\\rm{H}}$ (NED, cm$^{-2}$); Col. (7) Indicates whether the objects have polarised broad lines in the optical spectrum if data exists \\citep{tran03}; Cols. (8-10) Line ratios as measured by (1) \\cite{ho97}; (2) \\cite{degrijp92}; (3) \\cite{gu06}. \\end{minipage} \\end{table*}  \\section[]{Data Analysis}  \\subsection[]{Sample Selection} The parent sample for our study is the extended {\\it IRAS} 12 $\\umu$m sample of \\cite{rush93}. From this we select objects defined as Seyfert 2s  by the NASA/IPAC Extragalactic Database (NED). We perform our own post-hoc analysis of the optical spectra below.  We require there to be good quality X-ray data ({\\it XMM-Newton, Chandra or ASCA}) with a column density indicating that they are unabsorbed in X-rays ($N_{\\rm{H}} < 10^{22}$ cm$^{-2}$), from the literature or our own analysis.  The NED optical classifications may not necessarily be robust, so we searched the literature to deselect sources with incorrect or ambiguous optical classification. The remaining six objects, IRASF 01475-0740, NGC 3147, NGC 3486, NGC 3660, NGC 3976, NGC 4501 form our sample: they are all unambiguously classified as  AGN using line ratio diagnostics (Line ratios and references given in Table \\ref{obsdat} and BPT diagram presented in Fig. \\ref{bptdiag}).  \\subsection[]{Optical data} We compiled optical spectroscopy, spectropolarimetry and line ratio data for our objects from the literature and present these data with references in Table \\ref{obsdat} and Fig. \\ref{bptdiag}, which plots the ratio [O\\,{\\sc iii}] $\\lambda5007$/H$\\beta$ versus the ratio [N\\,{\\sc ii}] $\\lambda6584$/H$\\alpha$.  We plot these on top of the catalogue of low redshift SDSS galaxies of \\cite{kauffmann03}. The demarcation line is that defined by \\cite{kewley01} separating AGN and starbursts, hence showing that their AGN classification is not in doubt.  Additionally, we plot the line ratios of NGC 6810, a Seyfert 2 also unabsorbed in X-rays, but shown to have a dubious Seyfert 2 classification due to broader-than-normal optical lines produced by a super-wind \\citep{strickland07}.  The authors used an $XMM-Newton$ observation to show that the X-ray emission from this galaxy was probably due to X-ray binaries.  \\begin{figure} \\includegraphics[width=67mm,angle=90]{bptdiag.ps} \\caption{BPT diagram (Baldwin et al. 1981) of the six featured Seyfert 2 (green diamonds) sources plus NGC 6810 (red triangle), which has previously but incorrectly been classified as a Seyfert 2 (Strickland 2007), plotted on top of 55, 757 low redshift SDSS galaxies (blue dots) of Kauffmann et al. (2003). The demarcation line is that defined by Kewley et al. (2001) and separates AGN from starbursts.}\\label{bptdiag} \\end{figure}  \\subsection[]{X-ray Data} We carried out the X-ray analysis on {\\it XMM-Newton} observations of IRASF 01470-0740, NGC 3147, NGC 3486, NGC 3976 and NGC 4501; {\\it Chandra} observations of NGC 3147 and NGC 4501 and an {\\it ASCA} observation of NGC 3660 (X-ray observational information listed in table \\ref{xobsdat}).  For {\\it XMM-Newton} data we use {\\sc sas v7.0} tasks to perform spectral extractions on the EPIC-pn data, for the {\\it Chandra} ACIS-S data we use {\\sc ciao v3.4} tasks and use {\\it ASCA} data products from the {\\it Tartarus} database.  All spectra were grouped with a minimum of 20 counts per bin, with the exception of the {\\it Chandra} observation of NGC 4501, where the spectrum was grouped using a minimum of 7 counts.  Spectral fitting was carried out using {\\sc xspec v11.3}.  \\begin{table} \\centering \\caption{Information on the X-ray observations used in this analysis.  Exposure time in brackets give the time after filtering for background flares in {\\it XMM-Newton} data.}\\label{xobsdat} \\begin{tabular}{l l l l l} \\hline Object name\t& Observatory\t\t& Date\t\t& Obs. ID\t\t& exp.\t\t\\\\ &\t\t\t&\t\t&\t\t\t& (ks)\t\t\\\\ \\hline F01475-0740 \t& {\\it XMM-Newton}\t& 2004-02-01\t& 0200431101\t\t& 12(9) \t\\\\ NGC 3147 \t& {\\it XMM-Newton}\t& 2006-10-06\t& 0405020601\t\t& 17(14)\t\\\\ & {\\it Chandra}\t\t& 2001-09-19\t& 1615\t\t\t& 2\t\t\\\\ NGC 3486 \t& {\\it XMM-Newton} \t& 2001-05-09\t& 0112550101\t\t& 15(4)  \t\\\\ NGC 3660 \t& {\\it ASCA}\t\t& 1995-06-09\t& 73039000\t\t& 26 \t\t\\\\ NGC 3976\t& {\\it XMM-Newton}\t& 2006-06-16\t& 0301651801\t\t& 14(4) \t\\\\ NGC 4501\t& {\\it XMM-Newton} \t& 2001-12-04\t& 0112550801\t\t& 14(3) \t\\\\ & {\\it Chandra}\t\t& 2002-12-09\t& 2922 \t\t\t& 18 \t\t\\\\ \\hline \\end{tabular} \\end{table}  \\section[]{Estimates of Bolometric Luminosity} We estimate the bolometric luminosities, $L_{\\rm{Bol}}$, of the six AGN from their 12 ${\\umu}$m and [O\\,{\\sc iii}] fluxes using bolometric corrections, $\\kappa_{\\rm 12\\umu m}$ calculated from template SEDs of \\cite{mrr08} and $\\kappa_{\\rm [O\\,III]}$ published by \\cite{heckman04}.  We also estimate bolometric luminosities from the unabsorbed 2-10 keV (HX) luminosity using the mean bolometric correction, $\\kappa_{\\rm HX}$, of \\cite{vasudevan07}.  Table \\ref{lumins} presents the observed luminosities with these estimates and Fig. \\ref{lumfig} plots the estimated bolometric luminosities from the [O\\,{\\sc iii}] luminosity against the estimated bolometric luminosities from the 2-10 keV luminosity.  From this analysis, all six objects appear to be significantly under-luminous in the 2-10 keV X-ray band by factors of 10-100.  For a typical Seyfert 2 this is easily understood, as we expect the 2-10 keV X-rays to be suppressed by absorption, but of course for our sample it seems the evidence for that absorption in terms of the spectral shape is absent. As discussed below, despite the lack of any obvious soft X-ray absorption in these objects, some of the X-ray spectra present evidence for hidden hard components. This allows an additional estimate of the intrinsic AGN bolometric luminosity show as the red arrows in Fig. \\ref{lumfig}. These are discussed in more detail below.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Observed luminosities of the sample using distances from Table \\ref{obsdat}. $12 \\umu$m luminosity is calculated from the {\\it IRAS} flux density presented by Rush et al. 1993 and [O\\,{\\sc iii}] luminosities have been corrected for absorption using the formulation of Bassani et al. (1999). The respective predicted bolometric luminosity calculated from these measured luminosities are also given.}\\label{lumins} \\begin{tabular}{|l | c c c|c c c|} \\hline Object name & \\multicolumn{3}{|c|}{Observed luminosities (erg s$^{-1}$)} & \\multicolumn{3}{|c|}{$L_{\\rm{Bol}}$ (erg s$^{-1}$) predicted from:}\\\\ &$L_{\\rm{12{\\umu}m}}$ & $L_{\\rm [O\\,III]}$ & $L_{\\rm{HX}}$ & $L_{\\rm{12{\\umu}m}}$ & $L_{\\rm [O\\,III]}$ & $L_{\\rm{HX}}$\\\\ \\hline IRASF 01475-0740& $7.68\\times 10^{43}$ & $4.80\\times 10^{41}$ & $4.71\\times 10^{41}$ & $8.54\\times 10^{44}$ & $1.68\\times 10^{45}$ & $1.28\\times 10^{43}$\\\\ NGC 3147 \t& $5.02\\times 10^{43}$ & $1.50\\times 10^{40}$ & $2.68\\times 10^{41}$ & $5.58\\times 10^{44}$ & $5.08\\times 10^{43}$ & $7.25\\times 10^{42}$\\\\ NGC 3486 \t& $2.83\\times 10^{42}$ & $3.89\\times 10^{38}$ & $3.26\\times 10^{39}$ & $3.14\\times 10^{43}$ & $1.36\\times 10^{42}$ & $8.82\\times 10^{40}$\\\\ NGC 3660 \t& $3.15\\times 10^{43}$ & $9.38\\times 10^{40}$ & $8.72\\times 10^{41}$ & $3.50\\times 10^{44}$ & $3.28\\times 10^{44}$ & $2.36\\times 10^{43}$\\\\ NGC 3976 \t& $1.59\\times 10^{43}$ & $4.30\\times 10^{39}$ & $2.18\\times 10^{40}$ & $1.76\\times 10^{44}$ & $1.50\\times 10^{43}$ & $5.90\\times 10^{41}$\\\\ NGC 4501 \t& $6.79\\times 10^{43}$ & $9.57\\times 10^{39}$ & $1.26\\times 10^{40}$ & $7.54\\times 10^{44}$ & $3.35\\times 10^{43}$ & $3.42\\times 10^{41}$\\\\  \\hline \\end{tabular} \\end{minipage} \\end{table*}  \\begin{figure} \\includegraphics[width=67mm,angle=90]{lumplot.ps} \\caption{Estimates of the bolometric luminosities from $L_{{\\rm [O\\,III]}}$ with $1\\sigma$ variance (green lines) against the estimated bolometric luminosities from $L_{\\rm{HX}}$.  The solid black line displays where on the plot these estimates would lie should they be equal, as expected, and the dashed black lines represent the $1\\sigma$ spread in $\\kappa_{\\rm HX}$.  As these objects are all positioned under the solid line, they are all under-luminous in X-rays.  Red arrows show how the bolometric luminosity predicted from $L_{\\rm{HX}}$ changes when we try to estimate the intrinsic $L_{\\rm{HX}}$ of the nucleus based on the X-ray spectral fitting.}\\label{lumfig} \\end{figure}   \\section[]{X-Ray Analysis} \\subsection{IRASF 01475-0740} The {\\it XMM-Newton} spectrum of IRASF 01475-0740 is well fitted by a power-law absorbed by a column of $N_{\\rm{H}} = 4.1 \\times 10^{21}$ cm$^{-2}$.  While this will produce some reddening and extinction in the optical, it is unlikely to be sufficient to suppress the optical broad lines entirely unless the absorber has an anomalously high dust-to-gas ratio (c.f. MCG-6-30-15; \\cite{reynolds97}).  The X-ray spectrum also shows line emission between 6 and 7 keV.  The centroid energy of this line is not well constrained, so we assume that the emission is from neutral iron K$\\alpha$ at a rest energy of 6.4 keV, though it may originate from ionised iron.  A gaussian fit to this feature with the centroid energy fixed at 6.4 keV gives an equivalent width (EW) of 160 eV. On the face of it, this relatively modest  iron K$\\alpha$ EW argues against a hidden Compton thick nucleus (Bianchi et al. 2008). We found, however, that it was possible to fit a second, heavily absorbed component from new Monte-Carlo models of  X-rays in heavily obscured AGN, incorporating line emission (Brightman et al., in preparation), similar to those presented by e.g. \\cite{ghisellini94} and \\cite{krolik94}. For simplicity given the limitations of the data we fit the extreme case of 4$\\pi$ coverage (i.e. a spherical distribution of matter), and we constrained the column density of the heavily absorbed component to $N_{\\rm{H}}=2 \\times 10^{24}|^{+65}_{-0.8}$ cm$^{-2}$. Although it is not formally required in the fit, this demonstrates a clear scenario in which the multiwavelength data can be reconciled, as such a column is easily sufficient to suppress the direct nuclear emission in the optical. By correcting for absorption in this heavily obscured component, we can make another prediction of the intrinsic luminosity of the nucleus and calculate the corresponding bolometric luminosity. This now agrees well with estimations from $L_{\\rm [O\\,III]}$ and $L_{\\rm{12{\\umu}m}}$.  The {\\it XMM-Newton} observation also shows no clear signs of variability. The spectropolarimetric survey of \\cite{tran03} reveals that this AGN hosts a HBLR, in full agreement with our conclusion of a heavily buried nuclear contribution. The same electrons which scatter the broad optical lines into the line of sight can then also scatter nuclear X-rays accounting for the apparently unobscured nature of the soft spectrum.  \\subsection{NGC 3147} Our analysis of the {\\it XMM-Newton} spectrum of NGC 3147 shows no absorption above the Galactic column, however it features a small (130 eV) iron K$\\alpha$ line at 6.4 keV. These results are in full agreement with those of Bianchi et al. (2008).  \\cite{bianchi08} used the small EW of the iron K$\\alpha$ line to argue against the Compton thick nature of this source and came to the conclusion that NGC 3147 has an intrinsic absence of a BLR.  However, as we have shown in IRASF01475-0740, this conclusion is not necessarily robust. Fitting the Monte Carlo models to the spectrum once again shows that a very heavily obscured component is consistent with the spectrum. Here the column density is constrained to be $N_{\\rm{H}}=9 \\times 10^{23}|^{+419}_{-1.0}$ cm$^{-2}$, with the main constraint coming from the iron K$\\alpha$ line.  We can again use this second component to predict the intrinsic luminosity of this source and calculate the corresponding bolometric luminosity, which also agrees well with estimations from $L_{\\rm [O\\,III]}$ and $L_{\\rm{12{\\umu}m}}$.  Other evidence suggests, however, that this may not be the correct interpretation in this case. The 2ks {\\it Chandra} observation of NGC 3147 shows that the measured flux of the nucleus is $\\sim2$ times that measured by {\\it XMM-Newton} indicating that it is variable in nature and hence it must be being observed directly, rather than in scattered light. Furthermore,  the {\\it Chandra} imaging reveals no evidence for extra-nuclear X-ray sources that could be producing the unabsorbed X-ray profile or the variability, and spectral extraction from the nuclear region in the {\\it Chandra} image does not reveal a significantly harder X-ray spectrum then the larger {\\it XMM-Newton} beam (Table~4).  \\subsection{NGC 3486} The {\\it XMM-Newton} spectrum of NGC 3486 is well fitted by a power-law absorbed by the Galactic column only and so appears to be another unabsorbed Seyfert 2 galaxy.  However, an excess at hard energies points to a different scenario. The spectrum is also well fitted by the addition of a Compton reflection model ({\\tt pexmon}, \\cite{nandra07}) with an underlying thermal component ({\\tt raymond}), so this could in fact be a Compton thick Seyfert 2 which looks unabsorbed.  We also attempt to add a strongly absorbed transmission component to the model, which also improves the fit from a simple power-law.  However, the column density of this component is poorly constrained, so we fix this to an arbitrary log($N_{\\rm{H}}$) = 24.5. Estimates of the intrinsic $L_{\\rm{HX}}$ from the reflection component predict a $L_{\\rm{Bol}}$ which agrees well with $L_{\\rm{Bol}}$ estimated from $L_{\\rm [O\\,III]}$, suggesting that NGC 3486 is under-luminous in X-rays due to Compton thick obscuration. Finally, there is no evidence for variability of the X-ray flux in the {\\it XMM-Newton} observation.  \\subsection{NGC 3660} As there have been no observations of NGC 3660 with {\\it XMM-Newton} or {\\it Chandra}, we use the {\\it ASCA} observation and its documentation in the {\\it Tartarus} catalogue.  The {\\it ASCA} spectrum is fitted well by a power-law with no absorption above the Galactic column. Fig. \\ref{ngc3660lc} shows the light-curve for the 26 ks observation which shows significant variability on short time-scales.  \\begin{figure} \\centering \\includegraphics[width=84mm]{ngc3660_lc.ps} \\caption{{\\it ASCA} lightcurve of NGC 3660 showing the X-ray variable nature of the source.}\\label{ngc3660lc} \\end{figure}   \\subsection{NGC 3976} The {\\it XMM-Newton} spectrum of NGC 3976 is very similar to that of NGC 3486 as it shows no intrinsic absorption, but a hard excess present above the simple power-law fit allows us to add a heavily absorbed power-law component. A strongly absorbed transmission component produces an overall better fit, suggesting that NGC 3976 is also hiding a heavily obscured nucleus beneath its apparently unabsorbed soft X-ray spectrum. Again, there is also no variability detected in this observation.  \\subsection{NGC 4501} The {\\it XMM-Newton} spectrum of NGC 4501 reveals another apparently unabsorbed Seyfert 2 galaxy in X-rays.  It is well fitted by a power-law with no intrinsic absorption plus emission from a thermal plasma component, but shows no clear hard excess or other spectral evidence supporting a deeply buried AGN (such as iron K$\\alpha$ emission).  An entirely different picture emerges, however, when one considers the {\\it Chandra} data, which have higher spatial resolution. The {\\it Chandra} image of NGC 4501 reveals a hard X-ray emission coincident with the optically defined nucleus (Fig. \\ref{chanimg}) consistent with heavy X-ray absorption. It also shows that there are also multiple extra-nuclear sources present, including diffuse soft emission close to the nucleus, which {\\it XMM-Newton} could not resolve. The implication is that the unabsorbed nature of the {\\it XMM-Newton} spectrum, which has a larger beam size, is due to contamination, and that the true nuclear emission is heavily obscured. We performed a spectral extraction of the optically defined nucleus using the region identified in Fig \\ref{chanimg} which did indeed reveal a hard excess above the unabsorbed power-law, as in NGC 3486 and NGC 3976, which we fit with a Compton reflection component ({\\tt pexmon}). Using the reflection component to estimate the intrinsic $L_{\\rm{HX}}$ shows that the X-ray faintness of NGC 4501 is probably due to heavy absorption (Fig. \\ref{lumfig}).  There is also no evidence for variability of NGC 4501, either between subsequent observations by {\\it XMM-Newton} and {\\it Chandra}, or during them.  \\begin{figure} \\centering \\includegraphics[width=50mm]{NGC4501_img.ps} \\caption{Chandra image of NGC 4501 smoothed by a $\\sigma=3$ pixels Gaussian. Red represents 0.5-1.0 keV emission, green represents 1.0-3.0 keV emission and blue represents 3.0 - 10.0 keV in a log scale.  The black cross marks the position of the optical nucleus, the smaller green circle marks the {\\it Chandra} extraction region and larger green circle marks the equivalent {\\it XMM-Newton} extraction region.  The white area coincident with the nucleus indicates that it is the source of hard X-ray emission, and therefore likely to be heavily absorbed.}\\label{chanimg} \\end{figure}   \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{X-ray spectral fitting data for the sample, where R is the ratio of the underlying transmitted or reflected component to the power-law}\\label{fitdat} \\begin{tabular}{l l l l l l l l l l r@{}l} \\hline Object name \t& Model\t& {$N_{\\rm{H,1}}$} \t& {$N_{\\rm{H,2}}$} \t& $\\Gamma$ \t\t& EW$_{6.4}$ \t& kT \t\t\t& R\t\t\t& $\\chi_r^2/\\nu$& {$F_{\\rm HX}$}& \\multicolumn{2}{c}{$L_{\\rm int}$}\\\\ (1)\t\t& (2)\t& (3)  \t\t\t& (4)\t \t\t& (5)  \t\t\t& (6) \t\t& (7)\t\t\t& (8)\t\t\t& (9)\t\t& (10) \t\t& & (11)\t\\\\ \\hline IRASF01475-0740 & A\t& $0.41^{+0.07}_{-0.05}$& -\t\t\t& $2.02^{+0.06}_{-0.06}$& $160^{+300}_{-160}$& -\t\t& -\t\t\t& 65.5/58\t& 0.84\t\t& 49 &\t\t\\\\ ({\\it XMM-Newton})& C\t& $0.41^{+2.14}_{-0.18}$& -\t\t\t& $2.15^{+0.45}_{-0.24}$& - \t\t& - \t\t\t& $1.0^{+32}_{-1.0}$\t& 64.4/56 \t& 0.84\t\t& 48 &\t\t\\\\ & D\t& $0.46^{+0.17}_{-0.17}$& $200^{+6500}_{-80}$\t& $2.16^{+0.26}_{-0.23}$& - \t\t& - \t\t\t& $10^{+1600}_{-10}$\t& 63.3/56\t& 0.82\t\t& 465 &\t\t\\\\ \\hline NGC 3147 \t& A\t& $0.03^{+0.01}_{-0.01}$& -\t\t\t& $1.60^{+0.06}_{-0.06}$& $130^{+80}_{-100}$& -\t\t\t& -\t\t\t& 103/113\t& 1.4 \t\t& 27 &\t\t\\\\ ({\\it XMM-Newton})& C\t& $0.03^{+0.04}_{-0.01}$& -\t\t\t& $1.63^{+0.09}_{-0.07}$& - \t\t& -    \t\t\t& $1.6^{+8.3}_{-1.6}$\t& 103/112\t& 1.4 \t\t& 42 &\t\t\\\\ \\vspace{0.1in}\t& D\t& $0.04^{+0.01}_{-0.02}$& $90^{+4210}_{-10}$\t& $1.68^{+0.04}_{-0.11}$& -\t\t& -    \t\t\t& $1.9^{+10}_{-1.9}$\t& 102/112\t& 1.4 \t\t& 49 &\t\t\\\\ NGC 3147\t& A\t& $<0.07$\t\t& -\t\t\t& $1.50^{+0.19}_{-0.12}$& -\t\t& -\t\t\t&  -\t\t\t& 62.2/49\t& 3.2\t\t& 61 &\t\t\\\\ ({\\it Chandra})\t&\t&\t\t\t& \t\t\t&\t\t\t&\t\t&\t\t\t&\t\t\t&\t\t&\t\t& &\t\t\\\\ \\hline NGC 3486 \t& A\t& $0.03^{+0.17}_{-0.03}$& -\t\t\t& $2.44^{+1.17}_{-0.61}$& - \t\t& - \t\t\t& -\t\t\t& 28.3/18\t& 0.04\t\t& 0 & .10\t\\\\ ({\\it XMM-Newton})& B\t& $<0.10$\t\t& -\t\t\t& $1.73^{+1.29}_{-0.83}$& -\t\t& $0.23^{+0.13}_{-0.09}$& -\t\t\t& 24.8/16\t& 0.09\t\t& 0 & .20\t\\\\ & C\t& $<0.06$\t\t& -\t\t\t& $1.9^{\\dag}$\t\t& -\t\t& $0.24^{+0.19}_{-0.12}$& $34^{+220}_{-34}$\t& 21.5/15\t& 0.15\t\t& 4 & .8\t\\\\ & D\t& $<0.06$\t\t& 320$^{\\dag}$\t\t& $1.9^{\\dag}$\t\t& -\t\t& $0.24^{+0.66}_{-0.10}$& $200^{+650}_{-200}$\t& 23.2/16\t& 0.09\t\t& 31 & \t\t\\\\  \\hline NGC 3660 \t& A\t& $0.03^{\\dag}$\t\t& -\t\t\t& $1.82^{+0.04}_{-0.04}$& - \t\t& -    \t\t\t& -\t\t\t& 301/304\t& 2.3\t\t& 87 &\t\t\\\\ ({\\it ASCA})\t&\t&\t\t\t& \t\t\t&\t\t\t&\t\t&\t\t\t&\t\t\t&\t\t&\t\t& &\t\t\\\\ \\hline NGC 3976\t& A\t& $0.06^{+0.15}_{-0.06}$& -\t\t\t& $1.80^{+1.15}_{-0.60}$& - \t\t& -    \t\t\t& -\t\t\t& 38.0/33\t& 0.07\t\t& 1 & .3\t\\\\ ({\\it XMM-Newton})& B\t& $0.10^{+0.10}_{-0.07}$& -\t\t\t& $1.9^{\\dag}$\t\t& -\t\t& $0.44^{+0.29}_{-0.18}$& -\t\t\t& 26.2/32\t& 0.05\t\t& 0 & .97\t\\\\ & C \t& $0.09^{+0.09}_{-0.06}$& -\t\t\t& $1.9^{\\dag}$\t\t& - \t\t& $0.43^{+0.83}_{-0.14}$& $43^{+49}_{-43}$\t& 27.0/30  \t& 0.10\t\t& 37 &\t\t\\\\ & D\t& $0.10^{+0.06}_{-0.04}$& $17^{+190}_{-12}$\t& $1.9^{\\dag}$\t\t& -\t\t& $0.44^{+0.18}_{-0.11}$& $3.0^{+25}_{-3.0}$\t& 24.8/33\t& 0.12\t\t& 4 & .2\t\\\\  \\hline NGC 4501\t& A \t& $0.21^{+0.33}_{-0.14}$& -\t\t\t& $3.91^{+2.76}_{-1.14}$& - \t\t& - \t\t\t& -\t\t\t& 21.2/13\t& 0.02\t\t& 0 & .36\t\\\\ \\vspace{0.1in} ({\\it XMM-Newton})& B\t& $0.02^{+0.10}_{-0.02}$& -\t\t\t& $2.13^{+1.14}_{-0.63}$& -\t\t& $0.36^{+0.20}_{-0.08}$& -\t\t\t& 6.82/11\t& 0.08\t\t& 1 & .2\t\\\\ NGC 4501\t& A\t& $<0.02$\t\t& -\t\t\t& $1.49^{+0.29}_{-0.36}$& -\t\t& -\t\t\t& -\t\t\t& $44.7^C$\t& 0.05\t\t& 0 & .78\t\\\\ ({\\it Chandra})\t& B\t& $1.75^{+0.12}_{-0.31}$& -\t\t\t& $1.65^{+0.25}_{-0.31}$& -\t\t& $0.05^{+0.01}_{-0.01}$& -\t\t\t& $22.7^C$\t& 0.07\t\t& 1 & .1\t\\\\ & C \t& $0.66^{+1.11}_{-0.66}$& -\t\t\t& $1.9^{\\dag}$\t\t& -\t\t& $0.39^{+0.16}_{-0.10}$& $29^{+32}_{-20}$\t& $8.64^C$\t& 0.08\t\t& 13 &\t\t\\\\ & D \t& $0.16^{+0.47}_{-0.16}$& $18^{+47}_{-15}$\t& $1.9^{\\dag}$\t\t& -\t\t& $0.35^{+0.34}_{-0.17}$& $8.9^{+28}_{-6.4}$\t& $10.7^C$\t& 0.09\t\t& 2 & .6\t\\\\    \\hline \\end{tabular} Notes: HX = 2-10 keV; $^\\dag$ indicates a fixed parameter; $^C$ indicates use of Cash statistics in spectral fitting; Col. (1) Galaxy name as given by \\cite{rush93}; Col. (2) Model used for fitting, A=pl, B=pl+therm, C=pl+therm+refl, D=pl+therm+trans; where pl = simple power-law; trans = Monte-Carlo model of transmitted component; therm = thermal component (raymond); refl = pure reflection component (pexmon); Col. (3) Column density of simple power-law, units of $10^{22}$ cm$^{-2}$; Col. (4) Column density of transmitted component, units of $10^{22}$ cm$^{-2}$; Col. (5) Power-law index, pegged for simple power-law and transmitted component; Col. (6) Equivalent width of the neutral iron line at 6.4 keV if present in eV; Col. (7) Temperature of the thermal component in keV; Col. (8) The normalisation ratio of the reflection or transmitted components to the simple power-law; Col. (9) Reduced chi-squared and number of degrees of freedom; Col. (10) Absorption corrected observed flux in the hard (2-10 keV) band, units of $10^{-12}$ erg cm$^{-2}$ s$^{-1}$; Col. (11) Hard X-ray intrinsic luminosity, units of $10^{40}$ erg s$^{-1}$, calculated from reflection or transmitted component where present. \\end{minipage} \\end{table*}   \\begin{figure*} \\begin{minipage}{160mm} \\includegraphics[width=80mm]{IRASF01475_mo.ps} \\includegraphics[width=80mm]{NGC3147_mo.ps} \\includegraphics[width=80mm]{IRASF01475_data.ps} \\includegraphics[width=80mm]{NGC3147_data.ps} \\includegraphics[width=80mm]{NGC3147_chan_mo.ps} \\includegraphics[width=80mm]{NGC3486_mo.ps} \\includegraphics[width=80mm]{NGC3147_chan_data.ps} \\includegraphics[width=80mm]{NGC3486_data.ps} \\end{minipage} \\end{figure*}  \\begin{figure*} \\begin{minipage}{160mm} \\includegraphics[width=80mm]{NGC3660_mo.ps} \\includegraphics[width=80mm]{NGC3976_mo.ps} \\includegraphics[width=80mm]{NGC3660_data.ps} \\includegraphics[width=80mm]{NGC3976_data.ps} \\includegraphics[width=80mm]{NGC4501_mo.ps} \\includegraphics[width=80mm]{NGC4501_chan_mo.ps} \\includegraphics[width=80mm]{NGC4501_data.ps} \\includegraphics[width=80mm]{NGC4501_chan_data.ps}  \\caption{{\\it XMM-Newton} (unless otherwise stated) spectra of IRASF 01475-0740, NGC 3147, NGC 3486, NGC 3660, NGC 3976 and NGC 4501, shown with best fit model from Table \\ref{fitdat}.  In the plots of the best fit model, blue dashed lines represent the unabsorbed power-law, red dot-dashed lines represent the thermal component if present, and the magenta dotted lines represent the reflection or transmitted component if present.} \\label{specfig} \\end{minipage} \\end{figure*} ",
        "conclusions": "\\subsection{Main Findings}  We have presented X-ray and multi-waveband data for a sample of bona fide Seyfert 2 galaxies, as defined by their optical line ratios, which appear unabsorbed in the X-ray. Our first key finding is that the 2-10 keV X-ray luminosities of the objects are low with respect to their mid-IR and optical [O\\,{\\sc iii}] luminosities. Under normal circumstances this would be interpreted straightforwardly as being due to suppression of the X-ray flux by an absorbing column, however a naive analysis of the X-ray spectra is at odds with this interpretation. It may be that the simplest interpretation of the spectra is misleading us: if scattered or host galaxy emission dominates the spectrum below $\\sim 5$ keV, the spectrum may appear unabsorbed when in fact a harder nuclear components is present. We favour the latter interpretation in four objects. These show tentative evidence for a hard excess which is likely to be a hidden reflection or transmission spectrum indicative of a heavily obscured nucleus, accounting for the low hard X-ray luminosity of these objects. The most compelling case is NGC 4501, where the high resolution {\\it Chandra} image reveals a hard source co-incident with the nucleus embedded in a number of softer emitting regions. The clear conclusion in this case is that is that  2-10 keV {\\it XMM-Newton} data are  dominated by non-nuclear emission, accounting for the unabsorbed X-ray spectrum. We infer a similar conclusion for IRAS F01475-0740, NGC 3486 and NGC 3976. For two of our sources (NGC 3147 and NGC 3660) we see no signs of hidden reflection or transmission in the {\\it XMM-Newton} X-ray spectra. NGC 3147 shows a single soft X-ray source co-incident with the nucleus in the Chandra image, and both exhibit X-ray variability. These observations clearly point to the idea that we are seeing the nuclei directly in these sources, so the evidence suggest that they might intrinsically lack a BLR.  It therefore appears from our analysis that there are two populations of unabsorbed Seyfert 2 galaxies and that they appear unabsorbed in X-rays for different reasons entirely.  It may be that our view of the AGN is genuinely unobscured, and that the lack of broad lines in their optical spectra is due to an intrinsic absence or weakness of the BLR.  Alternatively though, it may be that the X-ray spectrum is misleading, and that the unabsorbed spectra are due to scattered or galactic soft X-rays which dominate the emission from a much more powerful obscured nucleus.  We discuss each in turn.  \\subsection{Genuinely unobscured Seyfert 2s}  \\subsubsection{NGC 3147} NGC 3147 has no hidden broad line region and a {\\it Chandra} observations shows no discernible extra-nuclear sources.  Between observations by {\\it Chandra} and {\\it XMM-Newton}, the hard X-ray flux drops by a factor of $\\sim 2$ indicating that it is variable in X-rays which means that we are likely to be seeing the nucleus of NGC 3147 directly. \\cite{bianchi08} carried out simultaneous optical/X-ray observations of NGC 3147 to show that the apparent mismatch between optical and X-ray classification was not due to differential variability.  By noting a small equivalent width of the iron line and a large ratio between hard X-ray and [O {\\sc iii}] fluxes, they come to the same conclusion that the nucleus of NGC 3147 is genuinely unobscured, and that this AGN must therefore intrinsically lack a broad line region.  We cannot however rule out the Compton thick nature of this source without observing it above 10 keV, where transmission from a Compton thick medium would dominate.  \\subsubsection{NGC 3660} The X-ray variability of this source on short time-scales and unabsorbed profile of its X-ray spectrum leads us to the conclusion that we are genuinely viewing the nucleus of NGC 3660 directly.  This leads us to believe that NGC 3660 also lacks a broad line region as we believe with NGC 3147. However, the origin of the X-ray variability is not necessarily nuclear as variability was recently discovered in the ultra-luminous X-ray (ULX) source, M82 X-1 \\citep{mucciarelli06}.  A {\\it Chandra} observation of NGC 3660 should be able to resolve any ULXs as we have shown for NGC 4501.  We also cannot rule out differential variability as the cause of the mismatch, as \\cite{bianchi08} did for NGC 3147.  As has been observed in other X-ray variable Seyferts (e.g. NGC 4151, \\cite{gaskell86}), the broad line flux has also been seen to vary.  If the broad line flux is observed in its low state, the object will be seen as type 2, despite being unobscured.  NGC 3660 will also need to be observed simultaneously in the optical and in the X-rays to rule out this possibility.  \\subsection{Misidentified Heavily Obscured Seyfert 2s}  \\subsubsection{IRASF 01475-0740} \\cite{bianchi08} use the small size of the iron line in NGC 3147 to argue against the Compton thick nature of that source.  However, IRASF01475-0740 also shows a small iron line, but in this object we know that heavy nuclear obscuration is occurring as shown by \\cite{tran03}.  This spectropolarimetric study showed that IRASF01475-0740 has broad lines in its polarised optical spectrum, which are missing in its normal optical spectrum, confirming that our line of sight to the nucleus is blocked by optically thick material.  We show that a Compton thick source may exhibit only a moderate iron line if the continuum emission below 10 keV is dominated by a strong scattered continuum or extra-nuclear emission.  We demonstrated this by adding a second component to the fit of the X-ray spectrum, with a much larger column density than the first component,  by using the iron line as a constraint.  The intrinsic hard X-ray luminosity of this second component is a factor of $\\sim 10$ greater than the observed luminosity.  Assuming that the observed hard X-ray spectrum is indeed dominated by scattered light, we calculate an upper limit on the scattering fraction by fixing the column density of the second component to its lower limit.  The upper limit on the scattering fraction turns out to be $50\\%$, far higher than the typical $\\sim 3\\%$ seen in Seyfert  2 galaxies \\citep{cappi06}.  \\cite{ueda07} observed a $<1\\%$ scattered fraction in two Compton thick AGN, SWIFT J0601.98636 and SWIFT J0138.64001. They used that fraction to suggest that these objects have a large covering fraction of the torus, or a low abundance of the gas responsible for the scattering in those objects.  The inverse could be true for IRASF0175-0740 - it may have a low covering fraction of the torus, or a high abundance of the gas responsible for the scattering.  Scattered components are also often seen in reflection dominated sources (eg. NGC 1068; \\citep{pier94}). If there is an underlying reflection component in IRASF0175-0740, this may also be dominated by scattered emission if the inclination angle of the torus is very high so that only a small part of the inner wall of the torus is directly visible and not self obscured.  This would give a very high scattered/reflected fraction as noted for model C in Table \\ref{fitdat}.  If the line emission seen here between 6 and 7 keV in fact originates from ionised iron, then we could be seeing reflection from ionised matter, as seen by \\cite{nandra07_2} in IRAS 00182-7112.  The effect would be to increase the soft X-ray flux compared to neutral reflection as the lighter elements responsible for absorbing the soft X-rays will have been stripped of most of their electrons.  This would fit well with the soft spectrum seen in IRASF01475-0740.  The case of IRASF01745-0740 shows that Compton thick/intermediate AGN do not have to be reflection dominated with intense iron lines as they can be transmission dominated with a strong scattered component and no intense iron lines.  \\subsubsection{NGC 4501} By constraining hard X-ray emission to the nucleus in the {\\it Chandra} observation of NGC 4501 we confirmed that this source is in fact Compton thick.  Since the extra-nuclear sources seen in the {\\it Chandra} image lie outside the {\\it XMM-Newton} spectral extraction region used, we surmise that the power-law and thermal emission which dominate over the reflection component in the {\\it XMM-Newton} spectrum does not originate from ULXs or the AGN, but from diffuse galactic emission close to the nucleus.  If this is scattered emission from the nucleus, the scattered fraction is $\\sim 3\\%$ in this Seyfert.   \\subsubsection{NGC 3486} Fitting the hard excess in the X-ray spectrum of NGC 3486 with a reflection model yields an improved fit to this data, but the significance is too low to conclude that the source is Compton thick, as a reflection spectrum would imply. However, the {\\it XMM-Newton} spectrum here, is in all ways similar to the {\\it Chandra} spectrum of NGC 4501, which we showed to originate from a hard nucleus which is very likely Compton thick. As NGC 3486 is severely under-luminous in hard X-rays when compared to optical and mid-infrared luminosities, this conclusion seems likely for NGC 3486 too.  Again we need to observe this Seyfert 2 above 10 keV to confirm our hypothesis, or with the high angular resolution capabilities of {\\it Chandra}.   \\subsubsection{NGC 3976} As with NGC 3486, the hard excess present in this X-ray spectrum is well fitted by the addition of a reflection component.  However, it is also well fitted by a transmission component.  Again, this spectrum is very similar to the {\\it Chandra} spectrum of NGC 4501, so we conclude that this source is also heavily obscured.   \\subsection{Wider Implications}  The discovery of AGN which do not have a broad line region has important consequences for the AGN unification model.  We can no longer exclusively invoke orientation based effects to explain the difference between type 1 and type 2 AGN.  There must be some other physical effect which means that the BLR is absent.  \\cite{nicastro03} studied a sample of Seyfert 2s extracted from the spectropolarimetric study of \\cite{tran03} and concluded that the presence of a hidden broad line region in Seyfert 2s is dependent on the rate at which matter is accreted by the central black hole.  They find that those which do not have a hidden broad line region have a mass accretion rate less than $10^{-3}$ Eddington units.  Thus, low accretion rate systems may not form a stable broad line region so they would not show broad lines in their spectra, yet will still be unobscured.  On the other hand, our multiwavelength analysis has shown a significant number (4 out of 6) Seyfert 2 galaxies in which the absorbing column appears to have been completely underestimated. The distribution of absorbing column densities is a key ingredient in models of the cosmic X-ray background, especially the fraction of Compton thick AGN, which is required to fit the 30 keV peak in the spectral energy distribution \\citep{gilli07}. The derived $N_{\\rm{H}}$ distribution from these models generally over-predicts the number of Compton thick AGN observed by many authors (e.g. \\cite{risaliti99}) and also under-predicts the number of unobscured AGN.  Our work potentially provides an answer to this discrepancy, at least for the 12 ${\\umu}$m sample.   From our small but well selected sample, we find that 2/3 of Seyfert 2 galaxies with an $N_{\\rm{H}}$ reported to be $<10^{21}$ cm$^{-2}$ should in fact be designated with an $N_{\\rm{H}}>10^{23}$ cm$^{-2}$, and in reality are probably Compton thick. The 12 ${\\umu}$m sample contains 79 Seyfert 2s with good quality X-ray coverage, and of these 19 have been reported to be unabsorbed (24\\%), compared to 4\\% found by  \\cite{risaliti99} for their sample of local Seyfert 2s. Combined with the consideration that many unabsorbed  Seyfert 2s of the 12 ${\\umu}$m sample have a dubious classification as Seyfert 2, our findings alter the $N_{\\rm{H}}$ distribution of the 12 ${\\umu}$m sample to better agree with those derived from the X-ray background and other observations.  Misidentifying these heavily obscured objects due to contamination by the host galaxy, as may be the case in NGC 3486, 3976 and 4501, should be an effect only seen in local AGN, as in deep surveys, only the brightest AGN are selected, where the nucleus will dominate over any galactic emission, but if scattered nuclear emission dominates the soft X-ray then this could be a significant issue even at high redshift."
    },
    "0808/0808.3013.txt": {
        "abstract": "Extremely red objects, identified in the early {\\it Spitzer} Space Telescope observations of the bright-rimmed globule IC 1396A and photometrically classified as Class I protostars Class II T Tauri stars based on their mid-infrared colors, were observed spectroscopically at 5.5--38 $\\mu$m ({\\it Spitzer} InfraRed Spectrograph), at the 22 GHz water maser frequency (NRAO Green Bank Telescope), and in the optical (Palomar Hale 5-m), to confirm their nature and further elucidate their properties. The sources photometrically identified as Class I, including IC1396A:$\\alpha$, $\\gamma$, $\\delta$, $\\epsilon$, and $\\zeta$, are confirmed as objects dominated by accretion luminosity from dense envelopes, with accretion rates 1--10$\\times 10^{-6} M_{\\odot}$~yr$^{-1}$ and present stellar masses 0.1--2 $M_{\\odot}$. The Class I sources have extremely red continua, still rising at 38 $\\mu$m, with a deep silicate absorption at 9--11 $\\mu$m, weaker silicate absorption around 18 $\\mu$m, and weak ice features including CO$_2$ at 15.2 $\\mu$m and H$_2$O at 6 $\\mu$m. % The ice/silicate absorption ratio in the envelope is exceptionally low for the IC 1396A protostars, compared to those in nearby star-forming regions, suggesting the envelope chemistry is altered by the radiation field or globule pressure. % Only one 22 GHz water maser %(characteristic of a Class 0 protostar) was detected in IC 1396A; it is coincident with a faint mid-infrared source, offset from near the luminous Class I protostar IC 1396A$\\gamma$. The maser source, IC 1396A$\\gamma_{b}$, has luminosity $<0.1 L_{\\odot}$, the first H$_{2}$O maser from such a low-luminosity object. Two near-infrared H$_{2}$ knots on opposite sides of IC 1396A:$\\gamma$ reveal a jet, with axis clearly distinct from the H$_{2}$O maser of IC 1396A:$\\gamma_{b}$.  The objects photometrically classified as Class II, including IC1396A:$\\beta$, $\\theta$, 2MASSJ 21364964+5722270, 2MASSJ 21362507+5727502, LkH$\\alpha$ 349c, Tr 37 11-2146 and Tr 37 11-2037, are confirmed as stars with warm, luminous disks, with a silicate emission feature at 9--11 $\\mu$m, and bright H$\\alpha$ emission, so they are young, disk-bearing, classical T Tauri stars. % The disk properties change significantly with source luminosity: low-mass (G--K) stars have prominent 9--11 emission features due to amorphous silicates while higher-mass (A--F) stars have weaker features requiring abundant crystalline silicates. A mineralogical model that fits the wide and low-amplitude silicate feature of IC1396A:$\\theta$ requires small grains of crystalline olivine (11.3 $\\mu$m peak) and another material to to explain the its 9.1 $\\mu$m peak; reasonable fits are obtained with a phyllosilicate, quartz, or relatively large ($>10$ $\\mu$m) amorphous olivine grains.  The distribution of Class I sources is concentrated within the molecular globule, while the  Class II sources are more widely scattered. Combined with the spectral results, this suggests two phases of star formation, the first (4 Myr ago) leading to the widespread Class II sources and the central O star of IC 1396, and the second ($< 1$ Myr ago) occurring  within the globule. The recent phase was likely triggered by the wind and radiation of the central O star of the IC 1396 \\ion{H}{2} region. %with possible additional contributions from the outflows of LkH$\\alpha$ 349a,c %and some nearby B stars. ",
        "introduction": "{\\it Spitzer} has opened widely the mid-infrared (3.6--24 $\\mu$m) window for studies of star formation. In the earliest observations with the observatory, as in many subsequent ones, previously unknown, mid-infrared-bright sources have been discovered in star-forming regions. In this paper we follow up \\citep[][Paper I]{reachIC} on the mid-infrared sources in the bright-rimmed globule IC 1396A, known as the Elephant Trunk Nebula, a dense globule in the large, nearby (750 pc) \\ion{H}{2} region IC 1396, which is is excited by the 4 Myr-old O6 star HD 206267 \\citep[see][]{weikard}. Based on comparison to measured colors to those of young stellar objects in Taurus \\citep{kenyonhartmann}, we suggested 10 sources may be the very early Class I stage, which is thought to only last $10^{5}$ yr \\citep{spbook}, indicating that the globule is a site of very recent star formation.  Detecting new protostars in a wide variety of star forming regions is an important expansion of previous work that by necessity has concentrated on the very closest star forming regions such as Taurus and Orion. While these nearby regions provide excellent opportunities to observe the widest range of star formation (from brown dwarfs to massive stars), they are necessarily parochial. Contrary to a long-standing common belief that the Sun formed in a quiescent region like Taurus, meteoritic evidence indicates the Sun likely formed near a high-mass star, whose \\ion{H}{2} region, winds, and supernova explosion influence the evolution and composition of the protoplanetary disk \\citep{meteorite,tachibana}. To search for present-day star-forming regions similar to those in which the Sun formed, we must consider massive star forming regions. In such regions, stars can be formed by radiative driven implosion of moderately dense cores \\citep{lefloch}, and planet-building disks are exposed to strong ultraviolet radiation \\citep{proplyd}. A major goal of future star formation studies should be to determine the effects of massive stars both in triggering star formation and shaping disks around young stellar objects.  In this paper, we present follow-up observations to examine the nature of the sources in IC 1396A in more detail. New {\\it Spitzer} mid-infrared spectra demonstrate that the photometrically extremely red objects are indeed Class I protostars, with dense, highly-optically-thick envelopes shrouding a moderately luminous core inside a dust photosphere. We show that at least one of the extremely red objects, IC 1396 A:$\\gamma$, powers a molecular outflow with H$_2$O maser emission, characteristic of a very young ($\\sim 10^{4}$ yr) Class 0 protostars \\citep{furuya}. The photometrically-classified Class II objects are compared to classical T Tauri stars, based on H$\\alpha$ emission and mid-infrared disk spectra. ",
        "conclusions": "We studied a mid-infrared-selected sample of young stellar objects from {\\it Spitzer} imaging (Paper I) using spectroscopic observations from radio to optical wavelengths. The photometric classifications in Paper I are confirmed by the spectroscopy. The Class I sources have deep silicate absorption features and the vast majority of their luminosity arises in the infrared. None of them show H$_{2}$O masers (characteristic of Class 0) or optical (Class II) counterparts. The only H$_{2}$O maser associated with our original sample turns out to be a separate object, IC 1396A:$\\gamma$b, with much lower luminosity $\\sim 0.02 L_{\\odot}$ than the protostar we identified originally, IC 1396A$:\\gamma$. The mid-infrared-selected Class I sources can be fitted with 0.1--2 $M_{\\odot}$ stars with ages $<2\\times 10^{5}$ yr.  Compared to stars forming in the more quiescent Taurus molecular cloud, the physical conditions of the globule from which the stars are forming include a higher radiation field due to the O6 star exciting the \\ion{H}{2} region and nearby young B stars. There is also a higher pressure due to compression by the ionization front from the O6 star and winds from the many young and still-forming stars. The wind from the $\\sim 5 M_{\\odot}$ star LkH$\\alpha$ 349a has cleared a hole in the center of the dense globule, and outflows from the Spitzer-discovered Class I objects appear to have swept material into `compartments' whose walls are over-pressured. % The spectroscopy results show that the Class I objects are brighter in the far-infrared than the theoretical models predict, suggesting they have a warmer envelope. % Ice features in the Class I object envelopes differ somewhat from those seen around such objects in other star-forming regions. While the same set of features is present, their amplitude is lower, per unit silicate dust absorption, and the wavelengths are slightly shifted. Both of these effects suggest the Class I envelopes are warmer than in Class I sources in other star-forming regions, perhaps due to the warmer gas temperature in the star-forming material in the globule IC 1396A. % The somewhat different physical conditions of the Class I envelopes in this globule will affect the chemistry at the time when planet formation is just beginning. Because the Sun is thought to have formed near an O star, and most nearby star-forming regions lack an O star, astronomical searches for the physical conditions of the early solar nebula would be best conducted in regions like IC 1396.  The stars photometrically classified as Class II are confirmed as such based on the presence of optical spectra with bright H$\\alpha$ emission lines and mid-infrared spectra with silicate features in emission. Most of the disk spectra are similar to other commonly-seen T Tauri star spectra, dominated by amorphous silicates. But the star IC 1396A:$\\theta$ has a particularly structured silicate feature with two `horns' at 9.1 and 11.3 $\\mu$m. The 11.3 $\\mu$m horn can be explained by crystalline olivine. The origin of the 9.1 $\\mu$m horn cannot be uniquely identified, with reasonably good fits provided by quartz (SiO$_{2}$), a phyllosilicate (montmorillonite), or $>12$ $\\mu$m radius amorphous olivine grains. Similar spectra are evident in other wide samples of T Tauri stars. Such spectra show that the planet-building material around at least some stars is dominated by crystalline silicates and is significantly different from interstellar silicates.  The star formation history in IC 1396A cannot be described as a single ``burst'' of star formation: Class I and II sources in and near the globule span a range of ages and masses in a small region (2 pc across). It appears that the earliest star formation in IC 1396 occurred approximately 5 Myr ago, based on the ages of the Class II sources and the central O6 star. Some stars ahead of the present-day bright rim of IC 1396A have become exposed by the recession of the globule and by their own outflows. Recently, $< 1$ Myr ago, a powerful wind from the young AeBe star LkH$\\alpha$ 349a blew a bubble through the center of the globule. The star formation in IC 1396A we observe now as Class I sources occurred within the last 0.2 Myr, based on the ages of the Class I sources.  The most recent star formation occurs preferentially behind the center of the globule, just inside bright rims. The specific trigger for the most recent star formation may have included not only  pressurization of the globule by the O star exciting IC 1396, but also the pressure of outflows and jets from the earlier generation of stars and in particular those with powerful winds like LkH$\\alpha$ 349a. The young stars within IC 1396A appear to be reshaping the globule from the inside, clearing `compartments' within the globule. Gradually, the globule is being fragmented due to the exponentially increasing number of stars (if indeed they are sequentially triggered). The mass contained in walls of the compartments will eventually become small enough, and their overlapping structure chaotic enough, that further star formation in the globule will end. The limiting factor in the star formation efficiency of this one globule will be the dynamic effect of energy input from the stars formed within it."
    },
    "0808/0808.3383_arXiv.txt": {
        "abstract": "% The Ara OB1a association is one of the closest sites where triggered star formation is visible for multiple generations of massive stars. At about 1.3 kpc distance, it contains complex environments including cleared young clusters, embedded infrared clusters, CO clouds with no evidence of star formation, and clouds with evidence of ongoing star formation. In this review we discuss the research on this region spanning the last half-century.  It has been proposed that the current configuration is the result of an expanding wave of neutral gas set in motion between 10--40 million years ago in combination with photoionization from the current epoch. ",
        "introduction": "Ara OB1a, while somewhat enigmatic, may be one of the best examples of triggered star formation in the local Galaxy.  The triggering in the most active portion is easily imagined from images of the $<$ 3 Myr NGC~6193/RCW~108-IR complex such as Figure~\\ref{ESO}.  However the full association may be as much as 50 Myr in age and cover several degrees on the sky.  Ara OB1 was first studied in detail by Whiteoak (1963). He used photographic photometry and objective prism spectroscopy to identify about 35 O and B star members.  Whiteoak's work was a follow up to an H$\\alpha$ survey by Rodgers, Campbell \\& Whiteoak (RCW; 1960), which cataloged 181 H$\\alpha$ emission regions in the southern sky, including RCW 108.  NGC~6193 is an open cluster, discovered by James Dunlop in 1826, that is dominated by a pair of O stars, HD 150135 and HD 150136.  These are the brightest optically revealed O stars in the association and are thought to be responsible for ionizing the bright rim of emission to the west (NGC 6188, discovered by John Herschel in 1836) which separates NGC 6193 from RCW 108--IR. The youth of the region is clear in sky survey plates and Figure~\\ref{ESO} which show concurrent regions of ionized gas and dust lanes. Whiteoak noted two additional clusters in the region, NGC~6204 -- about 2 degrees to the northeast -- and NGC~6167 -- about 1 degree to the southwest.  Ara OB1a is a compact association covering about 1 sq. degree around a central cluster -- NGC 6193. It is generally thought to be about 1.3 kpc away and can be equated with RCW~108.  There is some confusion in the literature as to what is actually meant by RCW~108. The original definition of Rodgers, Campbell \\& Whiteoak (1960) refers to all the region where H$\\alpha$ nebulosity is detected, for which they give a size of 210\\arcmin x 120\\arcmin centered at ({\\em l,b} = 336.49, --1.48). Straw et al. (1987) used RCW~108--IR to refer to the embedded IR cluster about 15\\arcmin\\ to the west of the O stars in NGC~6193 and which is identified with IRAS~16362--4845.  The confusion arises when RCW~108-IR is abbreviated by dropping the \"--IR\".  For the remainder of this paper we will refer to the embedded cluster as RCW~108--IR or IRAS~16362--4845.   Moffat \\& Vogt (1973) measured photometry for 13 stars within 4\\arcmin\\ of the center of NGC 6193 and found E$_{B-V}$= 0.4 and a distance of 1360 pc.  Herbst \\& Havlen (1977) describe Ara OB1 as having a ``diamond ring'' appearance, with RCW 108--IR as the ``diamond'' and a thin circular dust ring making up almost half of the ``ring''.  They performed photoelectric and photographic photometry on 702 stars in the region.  Herbst (1974, 1975) identified parts of this region as an ``R association'' as there were three early type stars associated with reflection nebulosity.  These stars may mark a separate site of star formation within Ara OB1a.  We will discuss this further in Section 6.  Mid--infrared MSX \\& IRAS maps of the region show several condensations. The brightest is coincident with RCW~108--IR and two apparently related peaks  IRAS 16379--4856 and IRAS~16348--4849.  % One of the earliest radio studies of the region was in the survey by Shaver \\& Goss (1970).  They made a 5 GHz and 408 MHz survey of over 250 Galactic radio source including RCW 108. RCW 108 was remarkable for having one of the smallest emission regions -- unresolved at 3\\arcmin\\ and having a relatively high density and temperature of the electron population.   In addition to Ara OB1a, Whiteoak (1963) identified a background O star cluster coincident with NGC~6193 but with a different distance modulus.  While the foreground Ara OB1a has a distance modulus of 10.5$\\pm 0.6$ the second group of about 13 O and B stars (Ara OB1b) has a distance modulus of 12.7 $\\pm 0.5$ or about 3500 pc. The central O stars of the two associations are offset by about 2$^{\\rm o}$ along the Galactic plane (Humphreys 1978), so this is a concern for membership determination.      \\begin{figure}[!ht] \\begin{center} \\includegraphics[scale = 0.5, angle = 0]{figure1.eps} \\end{center} \\caption{Salient features of the most active portion of Ara OB1a. This image is a B,V H$\\alpha$ (blue, green and red respectively) from the wide field imager on the MPI-ESO 2.2m telescope. The field of view is 32\\arcmin x 32\\arcmin\\ or about 12 pc on a side.} \\label{ESO} \\end{figure} ",
        "conclusions": ""
    },
    "0808/0808.4089_arXiv.txt": {
        "abstract": "We have made a new compilation of observations of maximum stellar mass versus cluster membership number from the literature, which we analyse for consistency with the predictions of a simple random drawing hypothesis for stellar mass selection in clusters. Previously,  Weidner and Kroupa have suggested that the maximum stellar mass is lower, in low mass clusters, than would be expected on the basis of random drawing, and have pointed out that this could have important implications for steepening the integrated initial mass function of the Galaxy (the IGIMF) at high masses. Our compilation demonstrates how the observed distribution in the plane of maximum stellar mass versus membership number is affected by the method of target  selection; in particular, rather low $n$  clusters with large maximum stellar masses are abundant in observational datasets that   specifically seek clusters in the environs of high mass stars. Although we do not consider our compilation to be either complete or unbiased, we discuss the method by which such data should be statistically analysed. Our very provisional conclusion is that the data is {\\it not} indicating any striking deviation from the expectations of random drawing. ",
        "introduction": "It is well known (following \\citet{weidner+kroupa2004,weidner+kroupa2006} and \\citet{oey+clarke2005})  that in the case of clusters containing fewer than $\\sim 100$ OB stars (i.e. those with mass $<$ a few $\\times~10^4~\\Msun$) the maximum stellar mass increases with cluster mass. At higher cluster mass scales, the value of the maximum stellar mass saturates at around 150--200~\\Msun\\  for reasons that are not entirely clear \\citep[see e.g.][]{zinnecker+yorke2007}. In this paper, we restrict ourselves to considering the lower mass regime. In Section 2 we review  why the statistics of maximum stellar masses in clusters of various scales can place constraints on high mass star formation in a cluster context and how rather subtle differences in assumed algorithms for cluster building are imprinted on the integrated galactic IMF (the IGIMF). We emphasise that analysis of the statistics of maximum stellar mass versus cluster mass offers the best prospects for an observational determination of whether the IGIMF should be  different from that measured in individual clusters \\citep[see e.g.][]{weidner+kroupa2006,elmegreen2006}. We also stress that competing algorithms can only be distinguished through proper statistical analysis of the observed distributions and that selecting algorithms according to how they reproduce the {\\it mean} trend can be misleading. In Section 3 we present a new (but in all likelihood still incomplete and biased) compilation of observational information on this issue and highlight  the sensitivity of the distribution obtained  to the method of target selection. In Section 4 we discuss the statistical inferences that can be drawn from the current dataset and conclude (Section 5) with an appeal for further observational information to be used in future analyses. ",
        "conclusions": "We have high-lighted the difficulty in analysing the data contained in Figure \\ref{mmaxntot} owing to difficulties to assigning regions of the diagram where the data is believed to be complete.  Nevertheless, our preliminary conclusion is that we are {\\it not} seeing strong evidence for a systematic suppression in maximum stellar mass in small $n$ clusters in addition to that expected on the basis of the statistics of random drawing (see also the complementary analysis of the statistics of isolated stars by \\citet{parker+goodwin2007}, which reached similar conclusions). Indeed, if anything, the feature of Figure \\ref{mmaxntot} that seems to be most discrepant with the random drawing model is the data hole in the range $m_\\mathrm{max} \\sim 3-10 \\Msun, n \\sim 50-100$. We are however aware that this might indeed be filled in if we have under-estimated the incompleteness in smaller $n$ clusters (particularly  due to the effects of dynamical evolution).  Our conclusion (in support of the random drawing hypothesis) remains provisional. Although we have set out what we believe to be a statistically correct methodology for analysing the problem, we are highly aware of the difficulties of properly quantifying observational selection effects. We therefore seek further input from observers in compiling a good sample for this kind of analysis."
    },
    "0808/0808.4040_arXiv.txt": {
        "abstract": "HIZOA J0836-43 is an extreme gas-rich ($M_{\\rm{HI}}$=7.5$\\times10^{10}\\, M_{\\sun}$) disk galaxy which lies hidden behind the strongly obscuring Vela region of the Milky Way. Utilizing observations from the {\\it Spitzer Space Telescope}, we have found it to be a luminous infrared starburst galaxy with a star formation rate of $\\sim 21\\, M_{\\sun}\\, \\rm{yr^{-1}}$, arising from exceptionally strong molecular PAH emission ($L_{7.7\\micron} = 1.50 \\times 10^{9} L_{\\odot}$) and far-infrared emission from cold dust. The galaxy exhibits a weak mid-infrared continuum compared to other starforming galaxies and U/LIRGs. This relative lack of emission from small grains suggests atypical interstellar medium conditions compared to other starbursts. We do not detect significant $[$Ne\\,{\\sc v}$]$ or $[$O\\,{\\sc iv}$]$, which implies an absent or very weak AGN. The galaxy possesses a prominent bulge of evolved stars and a stellar mass of 4.4($\\pm$1.4)$\\times10^{10}\\, M_{\\sun}$. With its plentiful gas supply and current star formation rate, a doubling of stellar mass would occur on a timescale of $\\sim$2 Gyr. Compared to local galaxies, HIZOA J0836-43 appears to be a ``scaled-up\" spiral undergoing inside-out formation, possibly resembling stellar disk building processes at intermediate redshifts. ",
        "introduction": "Understanding the fundamental origin and formation of galaxies requires the synthesis of multi-wavelength observations and cosmological-based numerical simulations. Advances in both areas have created a compelling view of a cold dark matter dominated universe where structure forms hierarchically \\citep{Wh78}. Of fundamental importance to this formalism is insight into the formation and evolution of galaxy disks \\citep{Str07}.  Massive, gas-rich disk galaxies ($M_{\\rm{HI}} > 10^{10} M_{\\sun}$) are considered indicative of relatively unadvanced star building systems, making them ideal laboratories for testing this formalism. However, such systems are rare in the local universe and nearly all massive \\HI-rich disk galaxies are inactive or only passively forming stars \\citep{Spray95}, thus providing few clues as to their formation and evolution. For example, Malin 1, an extreme case of a giant low surface brightness galaxy \\citep{Imp89}, has a dormant star formation rate of $\\sim 0.38\\, M_{\\sun}\\, \\rm{yr^{-1}}$ \\citep{Rah07}. In this letter, however, we present evidence of an \\HI-massive disk galaxy, HIZOA J0836-43, undergoing a vigorous starburst.  This galaxy, discovered as part of a blind \\HI\\ survey of the southern Zone of Avoidance, contains 7.5$\\times10^{10}\\, M_{\\sun}$ of \\HI\\ gas, has a total dynamical mass of 1.4$\\times10^{12}\\, M_{\\odot}$ and a 20-cm derived star formation rate of $\\sim$35\\,$M_{\\sun}\\, \\rm{yr^{-1}}$ \\citep{Don06}. It also appears to have a prominent bulge in the near-infrared -- bulge-to-disk ratio of $\\sim 0.80$ in the $K_{s}$ band -- central to an enormous, rapidly rotating \\HI\\ disk.  The galaxy is located at $l$=262.48$\\degr$, $b$=-1.64$\\degr$, lying behind the Vela Supernova Remnant of the Milky Way. At optical wavelengths it is largely hidden by foreground gas and dust. We have therefore conducted a detailed infrared study of HIZOA J0836-43, using imaging and spectroscopy from the {\\it Spitzer Space Telescope}, to fully reveal its morphology and past and present evolutionary state. Here we highlight key results from this study which provide new insights for understanding massive galaxy formation and evolution. The adopted distance of HIZOA J0836-43, $D_{L}$=148 Mpc, is from \\citet{Don06}, as derived from its recessional velocity, $v_{hel}$=10689 km\\,$\\rm{s^{-1}}$. ",
        "conclusions": "For a galaxy in the local universe HIZOA J0836-43 possesses a number of unusual infrared properties; considered in combination with the massive reservoir of gas that feeds it, this could be a rare instance of a local galaxy undergoing inside-out evolution, possibly resembling galaxy formation at earlier epochs.  Here we compare properties of HIZOA J0836-43 with those of local and intermediate redshift samples.  The paucity of warm dust is evident from the weakly rising continuum seen in its spectrum (Fig. 2b), consistent with its low $L_{24\\micron}/L_{70\\micron}$ colour compared to normal and star-forming systems, e.g., as compared to both the SINGS ({\\it Spitzer} Infrared Nearby Galaxy Survey) sample, and the relatively nearby Great Observatories All-sky LIRG Survey (GOALS). In contrast, the strength of the PAH emission in HIZOA J0836-43, both in luminosity and relative strength of the bands compared to the continuum, is amongst the largest observed in any star-forming galaxy \\citep{Peet04}, implying unusually strong PDR emission. These unusual MIR properties may arise from (1) a paucity of very small grains (VSGs) and/or (2) a soft UV radiation field.  Powered by massive star formation, transiently heated VSGs are thought to be the source of MIR radiation.  A weaker radiation field would give rise to cooler VSGs \\citep{Dra07}, which would re-radiate at longer FIR wavelengths along with the large grain, cool (T $\\sim$20 K) component. Both mechanisms would be consistent with an evolutionary phase in which the activity is confined to heavily gas/dust-obscured star forming regions, thus shielding the hard UV radiation from the rest of the disk, which in turn would yet to have experienced significant grain processing.  \\begin{figure}[!t] \\begin{center} \\includegraphics[width=8.5cm]{f3.eps} \\caption[Comparison with SINGS and Malin 1] {\\small{Infrared luminosity vs. \\HI\\ mass, comparing HIZOA J0836-43, Malin 1 and the SINGS galaxy sample. For Malin 1 (yellow triangle) a 70\\micron\\ upper limit \\citep{Rah07} is shown; upper limits for \\HI\\ mass are shown as black triangles.} } \\label{fig:3} \\end{center} \\end{figure}    The growth progress of HIZOA J0836-43 is deduced from the stellar bulge population and current star formation rate. We estimate a stellar mass of 4.4($\\pm$1.4)$\\times10^{10}\\, M_{\\sun}$ for the galaxy \\citep[from the relation of][]{Bell03} and hence a specific star formation rate, SFR per stellar mass, of $\\sim$0.5 $\\rm{Gyr}^{-1}$. Compared to local LIRGS \\citep[see Fig. 5 of][]{Wan06}, this implies active stellar building as facilitated by its plentiful supply of gas, with a doubling of stellar mass in $\\sim$2 Gyr. The specific star formation rate of the galaxy in combination with its stellar mass appears typical for star-forming galaxies at $z\\sim 0.7$, where gas fractions of disks were likely higher compared to local galaxies \\citep{Bell05, Per05}. Fig. 3 shows that the MIR luminosity is strongly correlated with \\HI\\ content; even with its extreme \\HI\\ mass, HIZOA J0836-43 appears consistent with being a ``scaled-up\" disk galaxy, unlike Malin 1 which is explicitly quiescent by comparison. This suggests relatively ``normal'' evolution in HIZOA J0836-43, despite lying at the extreme high end (i.e., early evolutionary stage) of the relation.     HIZOA J0836-43 is a gas-rich spiral galaxy exhibiting a warp in its disk (Fig. 1a), likely the result of a disturbance in its recent past ($<<$ 1 Gyr). Such an event could cause the observed starburst as gas from the extended \\HI\\ disk flows into the central region of the galaxy. Hence the starburst is powered by gas consumption, as opposed to a major merger event. Gas-rich galaxies, like HIZOA J0836-43, were likely more common in the distant universe and there is evidence that gas consumption and not merger interactions were driving stellar mass growth in the distant universe \\citep{Dad08}. Recent work has suggested that many LIRGs seen at intermediate redshifts ($z\\sim0.8$) achieve heightened star formation as a result of the high gas fractions of their disks and were less dependent on major interactions, compared to local LIRGs, to induce starburst activity \\citep{Mel08, Mar06}.   Observational evidence suggests that disk galaxies evolve along the stellar mass-radius relation and have built stellar mass intensely since $z=1$, on average by means of inside-out growth \\citep{Bar05,Tru06}. Comparing the extended regions of active star formation with the more centrally concentrated stellar bulge distribution (e.g., Fig 1b radial profile) suggests that HIZOA J0836-43 is undergoing vigorous disk building, an instance of inside-out growth. This combined with its PDR-dominated emission manifested as strong PAH emission coupled with a weak MIR continuum, makes it enigmatic in the local universe. Observing a galaxy at such a key point in its evolution could have far reaching implications for theories of galaxy formation and evolution."
    },
    "0808/0808.2046_arXiv.txt": {
        "abstract": "We discuss the relationship between a symmetry in the neutrino flavour evolution equations and neutrino flavour oscillations in the collective precession mode. This collective precession mode can give rise to spectral swaps (splits) when conditions can be approximated as homogeneous and isotropic. Multi-angle numerical simulations of supernova neutrino flavour transformation show that when this approximation breaks down, non-collective neutrino oscillation modes decohere kinematically, but the collective precession mode still is expected to stand out. We provide a criterion for significant flavour transformation to occur if neutrinos participate in a collective precession mode. This criterion can be used to understand the suppression of collective neutrino oscillations in anisotropic environments in the presence of a high matter density. This criterion is also useful in understanding the breakdown of the collective precession mode when neutrino densities are small. ",
        "introduction": "Because of neutrino-neutrino forward scattering or neutrino self-interaction \\cite{Fuller:1987aa,Notzold:1988kx,Pantaleone:1992xh} neutrinos can experience collective flavour transformation in environments such as the early Universe (e.g., \\cite{Samuel:1993uw,Kostelecky:1994dt,Samuel:1995ri,Pastor:2001iu,Dolgov:2002ab,Abazajian:2002qx}) and supernovae (e.g., \\cite{Qian:1994wh,Pastor:2002we,Balantekin:2004ug,Fuller:2005ae}) where neutrino number densities can be very large. This phenomenon is different from the conventional Mikheyev-Smirnov-Wolfenstein (MSW) effect \\cite{Wolfenstein:1977ue,Mikheyev:1985aa} in that the flavour evolution histories of neutrinos in collective oscillations are coupled together and must be solved simultaneously. The possibility and consequences of collective neutrino oscillations in supernovae were not well appreciated until the discovery that the ordinary matter can be ``ignored'' in such phenomena \\cite{Duan:2005cp} and the first numerical demonstrations of ``stepwise spectral swapping'' (or ``spectral split'') \\cite{Duan:2006jv,Duan:2006an} which is the imprint left by the collective flavour transformation on neutrino energy spectra.  Significant progress has been made towards understanding collective neutrino oscillations in supernovae (see, e.g., \\cite{Duan:2009cd} for a brief review and references therein). In particular, Raffelt and Smirnov \\cite{Raffelt:2007cb} demonstrated an adiabatic (precession) solution for the homogeneous, isotropic neutrino gas. This solution has been shown to agree in part with the results of ``single-angle'' simulations of supernova neutrino oscillations \\cite{Duan:2007fw}. These single-angle simulations essentially neglect the anisotropic nature of the supernova environment by assuming that the flavour evolution histories of neutrinos along all trajectories are identical to those along a ``representative'' trajectory, usually taken to be the radial trajectory \\cite{Qian:1994wh,Dasgupta:2008cu}. The adiabatic precession solution requires that at any time all neutrinos reside in a pure collective oscillation mode, the ``precession'' mode, which, as shown by Duan \\etal \\cite{Duan:2006an,Duan:2007mv}, would explain the spectral swap phenomenon in the single-angle simulations.    However, the real supernova environment is highly inhomogeneous and anisotropic. Here by ``inhomogeneous'' and ``anisotropic'' we are referring to the neutrino fields. Of course, the matter density distributions in the supernova environment are also likely to be inhomogeneous and anisotropic. To date there are a few ``multi-angle'' simulations \\cite{Duan:2006jv,Duan:2006an,Duan:2007bt,Fogli:2007bk} which, like the single-angle calculations, also adopt spherically symmetric supernova models but do treat flavour evolution along different neutrino trajectories in a self-consistent way. These multi-angle calculations also exhibit spectral swaps. It is still not understood how the spectral swap phenomenon arises in the (anisotropic) multi-angle context. In fact, some studies seem to suggest that collective neutrino oscillations in the isotropic and anisotropic environments can be very different. For example, collective neutrino oscillations of the bipolar type can experience ``kinematic decoherence'' and be disrupted in anisotropic environments \\cite{Raffelt:2007yz,EstebanPretel:2007ec}. Additionally, a very large matter background can result in neutrino oscillation phase differences between different neutrino trajectories in an anisotropic neutrino gas \\cite{Qian:1994wh}, and this effect recently has been shown to result in suppression of collective neutrino oscillations \\cite{EstebanPretel:2008ni}.  In this paper we discuss a SU($N_\\mathrm{f}$) rotation symmetry in the neutrino flavour evolution equations, where $N_\\mathrm{f}=2$ and 3 for the two-flavour and three-flavour neutrino mixing schemes, respectively. The collective precession mode for neutrino oscillations can ensue from this symmetry, even in inhomogeneous, anisotropic environments. This result explains the puzzling observations of the spectral swapping phenomenon in both the single-angle and multi-angle simulations of supernova neutrino oscillations.  The rest of this paper is organized as follows. In \\sref{sec:framework} we lay out the general framework for neutrino flavour transformation and discuss the SU($N_\\mathrm{f}$) rotation symmetry in the flavour evolution equations for a dense neutrino gas. In \\sref{sec:ap-sol} we show how the collective precession mode for neutrino oscillations can arise from this symmetry in various environments. We also give criteria for when the collection precession mode can occur. In \\sref{sec:sn} we present a new multi-angle simulation of supernova neutrino oscillations.  We analyze the results of this calculation guided by our understanding of the collective precession mode. In \\sref{sec:conclusions} we give our conclusions. ",
        "conclusions": "}  We have demonstrated the existence of the SU($N_\\mathrm{f}$) rotation symmetry in the neutrino flavour evolution equations. This symmetry can facilitate the establishment of the collective precession mode for neutrino flavour oscillations in various environments. The stepwise-spectral-swapping phenomenon can develop from such a collective neutrino oscillation mode in, e.g., supernovae. We have also given criteria for significant neutrino flavour oscillations to occur if neutrinos are entrained in the collective precession mode. These criteria can be used to understand the suppression of collective neutrino oscillations in anisotropic environments in the presence of a large matter density \\cite{EstebanPretel:2008ni}. These criteria also illuminate the process of the breakdown of collective oscillations when neutrino densities are low. The results obtained in both our new simulation and previous multi-angle calculations for supernova neutrino oscillations can be understood in terms of the collective precession mode and the criteria we have provided.  There remains much to be learned about the collective precession mode for neutrino oscillations. This collective mode can not exist when there is no solution to \\eref{eq:prec-ansatz-var} and \\eref{eq:L-cons-var}. A related interesting observation is that collective oscillations, if any, for a symmetric system with equal numbers of neutrinos and antineutrinos are quickly disrupted in the presence of even an infinitesimal anisotropy and flavour equipartition is obtained as a result \\cite{Raffelt:2007yz}. In the flavour pendulum analogy \\cite{Hannestad:2006nj}, the asymmetry of neutrinos and antineutrinos constitutes the internal spin of the flavour pendulum. The existence of this internal spin causes the flavour pendulum to undergo the precession which represents the collective precession mode for neutrino oscillations. There exists no collective precession mode in a symmetric neutrino system --- this explains the finding in \\cite{Raffelt:2007yz}. Meanwhile, it has been shown in \\cite{EstebanPretel:2007ec} that a typical asymmetry in the neutrino and antineutrino fluxes in the supernova environment will suppress  multi-angle decoherence and, therefore, make the collective precession mode for neutrino oscillations possible.    \\ack This work was supported in part by DOE grants DE-FG02-00ER41132 at the INT, DE-FG02-87ER40328 at UMN, NSF grant PHY-06-53626 at UCSD, and an IGPP/LANL mini-grant. This research used resources of the National Energy Research Scientific Computing Center, which is supported by the Office of Science of the US Department of Energy under Contract No.\\ DE-AC02-05CH11231. We thank J Carlson, J J Cherry, A Friedland, W Haxton, D B Kaplan, C Kishimoto, C Lunardini, A Mezzacappa, G Raffelt and P Romatschke for insightful discussions."
    },
    "0808/0808.0619_arXiv.txt": {
        "abstract": "{}{We present a panchromatic study,  involving a multiple technique approach, of the circumstellar disc surrounding the T~Tauri star IM~Lupi (Sz~82).} {We have undertaken a comprehensive observational study of IM Lupi using photometry, spectroscopy, millimetre interferometry and multi-wavelength imaging. For the first time, the disc is resolved from optical and near-infrared wavelengths in scattered light, to the millimetre regime in thermal emission. Our data-set, in conjunction with existing photometric data,  provides an extensive coverage of the spectral energy distribution, including a detailed spectrum of the silicate emission bands. We have performed a simultaneous modelling of the various observations, using the radiative transfer code MCFOST, and analysed a grid of models over a large fraction of the parameter space via Bayesian inference. }{We have constructed a model that can reproduce all of the observations of the disc. Our analysis illustrates the importance of combining a wide range of observations in order to fully constrain the disc model, with each observation  providing a strong constraint only on some aspects of the disc structure and dust content. Quantitative evidence of dust evolution in the disc is obtained: grain growth up to millimetre-sized particles, vertical stratification of dust grains with micrometric grains close to the disc surface and larger grains which have settled towards the disc midplane, and possibly the formation of fluffy aggregates and/or ice mantles around grains.}{} ",
        "introduction": "During the first stages of planet formation following the core nucleated accretion scenario \\citep{Lissauer06PPV}, evolution of dust grains within the protoplanetary disc surrounding the central forming object is expected. Micrometre size dust grains will start to grow by coagulation during low relative velocity collisions, leading to the formation of larger, potentially fluffy aggregates \\citep[][and references therein]{Beckwith00,Dominik06PPV,Natta06PPV} that will ultimately give birth to kilometre size planetesimals promoting gravitational focusing and eventually leading to planets. Timescales for the formation of these aggregates strongly depend on the physics of aggregation as well as the physical conditions within the disc, such as differential velocities between grains and grain-gas interactions.  \\begin{figure*} \\includegraphics[height=0.32\\hsize]{sz82_f606nup_sqrt.eps}% \\hspace{\\stretch{0.5}}% \\includegraphics[height=0.32\\hsize]{sz82_f814nup_sqrt.eps}% \\hspace{\\stretch{0.5}}% \\includegraphics[height=0.32\\hsize]{IMLup_1.6.eps}% \\caption{ The IM Lupi circumstellar disc observed in scattered light by HST. {\\bf Left and middle panels:} respectively F606W and F814 \\emph{WFPC2} PC1 PSF-subtracted images.  The dark diagonal feature running through the star is an artifact of charge bleeding along the CCD detector columns. {\\bf Right panel:} \\emph{NICMOS} PSF-subtracted coronagraphy at 1.6 $\\mu$m. The central circle represents the 0.3\\arcsec\\ radius coronagraphic obscuration. All images are shown in square root stretch, with North up and East to the left.  The field of view is 5x5\\arcsec (dashed boxed in Fig.~\\ref{fig:schema}). \\label{fig:obs_HST}} \\end{figure*}  \\begin{figure} \\centering \\includegraphics[width=\\hsize]{IM_Lup_schema3.eps} \\caption{\\emph{NICMOS} F160W image in $\\log$ stretch. The field of view is 11$\\times$11\\arcsec\\ with North up and East to the left. The dashed box corresponds to the field of view in Fig.~\\ref{fig:obs_HST}. The central white circle represents the 0.3\\arcsec\\ radius coronagraphic obscuration. The full green lines indicate the upper and lower scattering surfaces of the disc (separated by the dark lane which corresponds to the disc midplane). The green dashed circle represents the possible envelope surrounding the disc.\\label{fig:schema}} \\end{figure}  In parallel to grain growth, dust grains are expected to settle towards the disc midplane and then to migrate inward as a result of the conjugate actions of the stellar gravity and gas drag \\citep[e.g.][]{Barriere05,Fromang06}. This settling is highly dependent on the grain size. Small grains ($<$~1~$\\mu$m) are strongly coupled to the gas, they follow its motion, and do not settle at all. Conversely, grains in the mm-cm regime are considerably slowed down by the gas drag. They completely decouple and settle very rapidly in a thin midplane.   The result is the formation of a dust sub-disc close to the equatorial plane \\citep[e.g.,][]{Safronov69,Dubrulle95} that may become unstable and produce planetesimals when the density of the dust layer exceeds the gas one \\citep[e.g.,][]{Goldreich73,Schrapler04,Johansen07}. Models suggest grain growth to large particle sizes with a population of small grains remaining close to the surface and larger grains deeper in the disc \\citep[e.g.][]{Dominik06PPV}. However, many details of process remain uncertain. For example, it is not clear how dust grains overcome the ``metre size barrier'' without being accreted onto the star \\citep[e.g.\\ ][]{Weidenschilling77,Brauer08} or destroyed by high speed collisions \\citep[e.g.][]{Jones96,Blum00}, or how Kelvin-Helmholtz instabilities prevent the dust sublayer from fragmenting \\citep[e.g.][]{Johansen06}.  \\defcitealias{Dullemond04}{D04} \\def\\D04{\\citetalias{Dullemond04}}  Detection of the large building blocks of planets (particle sizes $> 1$m) in the disc midplane is far beyond current observational capacities, but we can expect to detect signatures of their formation. The stratified structure resulting from grain growth and settling has direct consequences on disc observables like their spectral energy distributions (SEDs) and scattered light images \\citep[e.g.][hereafter D04]{Dullemond04}. A variety of observational techniques have been used to obtain insights into the disc properties and their dust content:  SEDs \\citep[e.g.][]{D'Alessio01}, scattered light images \\citep{Watson06PPV}, thermal emission maps in the millimetre regime \\citep{Dutrey06PPV}, mid-infrared spectroscopy \\citep[e.g.][]{Kessler-Silacci06}. However, each technique only provides a limited view of a disc. SED analysis leads to multiple ambiguities (e.g., \\citealp{Thamm94,Chiang01}) since the lack of spatial resolution precludes from solving the degeneracies between model parameters like geometry and opacity. Spatially resolved observations in a single spectral band (scattered light in the optical or near-infrared regime or thermal emission in the millimetre domain) also gives incomplete information about the disc. As the dust opacity is a steep function of wavelength and because high temperature  gradients are present within the disc, a single-band observation gives insight to only a limited region of the disc. Thus, scattered light images probe the surface layers of these optically thick discs at large radii whereas, millimetre observations are mainly sensitive to the bulk of the disc mass closer to the midplane. To precisely study fine physical processes like dust evolution and stratification, it is necessary to combine the aforementioned methods in a multi-wavelength, multiple technique observational and modelling approach.  In this paper, we study the circumstellar environment of the T~Tauri star IM~Lupi.   The disc surrounding IM~Lupi  (Schwartz~82, HBC~605, IRAS~15528-3747) was first detected in scattered light with the {\\it WFPC2} instrument on board the \\emph{Hubble Space Telescope} (\\emph{HST}) as part of a T~Tauri star imaging survey (Stapelfeldt et al. 2008, in prep.). IM~Lupi  is an M0 T~Tauri star located within the Lupus star forming clouds. It is one of four young stellar objects in the small $^{13}$CO(1-0) Lupus~2 core near the extreme T~Tauri star RU~Lupi \\citep{Tachihara96}. Despite the low accretion-related spectroscopic activity of IM~Lupi \\citep{Reipurth96,Whichmann99}, longer wavelength observations reveal ample evidence for circumstellar material in the system with strong mm continuum emission \\citep{Nuernberger97}.  Single dish $^{12}$CO and $^{13}$CO line observations indicate that the disc is gas rich and are consistent with a rotating disc model \\citep{vanKempen07}. The 3.3\\,mm continuum emission from  the disc was spatially resolved by \\cite{Lommen07}. Preliminary models by \\cite{Padgett06} showed that the infrared excess is well-reproduced by a disc model.   We have built a rich data set of observations of the disc: a well sampled SED, \\emph{HST} multiple wavelength scattered light images, \\emph{Spitzer} near- and mid-infrared spectroscopy and \\emph{SMA} millimetre emission maps. These observations provide different, complementary detailed views of the disc structure and dust properties. We investigate whether all of these observations can be interpreted in the framework of a single model from optical to millimetre wavelengths, and whether we can derive quantitative conclusions regarding the evolutionary stage of the disc.  In section~\\ref{sec:obs}, we describe the observations and data reduction procedure. In section~\\ref{sec:simple_estimation}, we first draw constraints on the disc and dust properties from the various observations and then, in section~\\ref{sec:grid_modelling}, we present simultaneous modelling of these observations. A detailed study of the dust properties is presented in section~\\ref{sec:dust_properties}. In section~\\ref{sec:discussion},  we discuss the implication of the results, and section~\\ref{sec:conclusion} contains the concluding remarks. ",
        "conclusions": "\\label{sec:discussion} \\subsection{A border line CTTS but with a massive disc}  IM~Lupi displays only a modest amount of emission-line activity, with a H$\\alpha$ equivalent width which is known to vary from 7.5 to 21.5\\,\\AA\\ \\citep{Batalha93}. Studies of H$\\alpha$ emission show evidence of accretion, including variability: \\cite{Reipurth96} concluded that the H$\\alpha$ feature shows an inverse P Cygni profile (classification IV-Rm) and \\cite{Whichmann99} observed a III-R profile. \\emph{International Ultraviolet Explorer}  low dispersion LWP spectra show that the Mg\\,II 2\\,798\\,\\AA\\ line also varies, by a factor of about 2 in net flux \\citep{Valenti03}. The relatively weak emission lines and lack of optical veiling caused \\cite{Finkenzeller87} and \\cite{Martin94} to classify this object as a weak-line T~Tauri star, although our results show it would be more properly categorised as a borderline classical T Tauri star. IM~Lupi is a relatively weak-lined T Tauri star with a large and massive circumstellar disc. Our modelling indicates a large dust disc mass of $M_\\mathrm{dust} = 10^{-3}\\,M_\\odot$, extending up to a radius of 400\\,AU. This large mass is puzzling given the weakness of the  H$\\alpha$ line, which suggests a low accretion rate. It is however possible that the diagnostics of accretion have only been observed during periods of low or moderate accretion.  \\subsection{Disc structure \\label{sec:disc_structure}}  Our multi-technique modelling allows us to quantitatively constrain most of the geometrical parameters of the disc. A flared geometry with a scale height of 10\\,AU at a reference radius of 100\\,AU is required. The best model has a midplane temperature of 14\\,K at 100\\,AU. If we assume the disc is vertically isothermal (with the temperature equal to the midplane temperature), the hydrostatic scale height $\\sqrt{k_B r^3 T(r) / G M_* \\mu}$ (where $\\mu$ is the mean molecular weight) is 12 and 8.5\\,AU, respectively, for a central mass object of 0.5 and 1\\,M$_\\odot$, respectively. This is in good agreement with the scale height deduced from the observations (10\\,AU), indicating that the outer parts of the disc are very likely to be in hydrostatic equilibrium. Figure~\\ref{fig:temperature} presents the calculated midplane  temperature compared to the temperature corresponding to the gas scale height of the best model. The agreement is very good at the inner edge and in the outer parts of the disc. In the central parts of the disc ($<30$\\,AU and excluding the very inner edge), the midplane temperature obtained in our passive disc model is too low to explain the scale height required by observations, indicating that an additional heating mechanism (like viscous heating) may be at play in the disc.  \\begin{figure} \\includegraphics[width=\\hsize]{temperature.eps} \\caption{Temperature structure of the disc. The full and dashed lines represent the midplane and surface temperatures, respectively. The disc surface at a given radius is defined as the altitude above the midplane where the temperature is maximal. Both temperatures become equal at the inner edge of the disc because it is directly heated by the stellar radiation. The shaded area represents the temperature corresponding to the vertical gas scale height of the best model, assuming that the disc is vertically isothermal and that the central mass is between 0.5 and 1\\,M$_\\odot$. \\label{fig:temperature}} \\end{figure}   The inner radius is constrained to be between 0.25 and 0.40\\,AU, which corresponds to a maximum dust temperature around 1\\,000\\,K. The modelling of the SED alone allows values of the inner radius down to 0.15\\,AU corresponding to temperatures close to the dust sublimation temperature (1\\,500\\,K).  The modelling of other observations gives additional constraints on the disc parameters and because of the correlations between parameters, this reduces the range of possible values for the inner radius. However, scattered images and millimetre visibilities constrain the large scale structure of the disc. Because we assume that the disc can be described in terms of power-laws from the inner edge to the outer radius, these constraints affect the derived inner radius. However, it is very likely that the description of the disc using power-laws is too simplistic, and more complex geometries may slightly shift the inner radius inwards, making it compatible with the dust sublimation radius.  The surface density exponent is found to be close to $-1$. This value corresponds to the median measurement for discs in Taurus-Aurigae obtained by \\cite{Andrews07}. This also corresponds to the theoretical value for a disc in steady-state accretion \\citep{Hartmann98}. The surface density at 5\\,AU is 70\\,g.cm$^{-2}$. This value is within the broad peak around the median value of 14\\,g.cm$^{-2}$ for Taurus \\citep{Andrews07}.   Considering a probable stellar mass of $\\approx 1\\,M_\\odot$ and a gas to dust mass ratio of 100, the disc to star mass ratio is $\\approx$ 0.1, meaning the disc may be unstable through gravitational collapse. The stability of the disc will depend on the surface density. Our derived value of the surface density exponent $\\alpha$=-1, as opposed to values of 0 or -2, provides disc stability at all radii according to the Toomre stability criterion \\citep{Toomre64}. Collapse of gravitationally unstable discs \\citep{Durisen07PPV} is one suggested mode for planet formation. The disc of IM Lupi representing about 1/10 of the star mass, local enhancement of density may be sufficient to start planet formation in the disc following this process.   \\cite{Hughes08} have shown that simultaneous studies of the dust continuum and CO emissions in several well-studied discs can be reproduced with disc models that include a tapered exponential outer edge and not a sharp outer radius as we have used here. Current observations of IM Lupi do not allow us to study in details the outer edge of the disc but the \\emph{NICMOS} image indicates that dust grains are still present at radii larger than 400\\,AU. However, as the counter nebulae of the disc is seen in scattered light images at 0.8 and 1.6\\,$\\mu$m, we can get a rough estimate of the maximum value for the surface density in front of this second nebulae. Dust present at radii larger than 400\\,AU must be optically thin, allowing us to see the counter nebulae through the disc and the tentative envelope. This can be translated into a upper limit for the surface density of the disc\\footnote{The constraint is stronger at 0.8\\,$\\mu\\mathrm{m}$ than at 1.6\\,$\\mu\\mathrm{m}$, due to the larger dust opacity.}: $\\Sigma (r\\approx500\\,AU) \\times \\kappa_{0.8\\mu\\mathrm{m}} \\lesssim 1$. If we assume that the dust composition and grain size distribution are the same as in the rest of the disc, this corresponds to $\\Sigma (r\\approx500\\,AU) \\lesssim 0.2\\,$g.cm$^{-2}$ at 500\\,AU, \\emph{i.e.} about 3.5 times smaller than the extrapolated density from our disc model, assuming it extends at radii larger than 400\\,AU. This implies that the density at radii larger than 400\\,AU must decrease significantly faster than the $1/r$ dependence we found for the disc in regions inside 400\\,AU.  \\subsection{Grain growth and dust settling\\label{sec:dust_evolution}}  As with many other classical T Tauri stars, the slope of IM Lupi's millimetre continuum suggests a dust opacity following a law close to $\\kappa_\\mathrm{abs}(\\lambda) \\propto \\lambda^{-1}$, indicating dust grains in the disc are larger than those in the interstellar medium. Millimetre photometry mostly traces the thermal emission of cold dust in the outer part ($> 15$\\,AU) of the disc midplane (see figure~\\ref{fig:flux_cumule}), while the hot grains in the central regions contribute little to this emission.  The strong silicate features in the mid-IR spectrum indicate the presence of micron-sized grains in the disc surface inside the first central AU. The silicate emission indeed comes from the central parts of the disc: 90\\,\\% of the emission at 10 and 20\\,$\\mu$m originates from radii smaller than 1 and 10\\,AU, respectively (Figure~\\ref{fig:flux_cumule}). Larger grains ($\\gtrsim 3\\,\\mu$m) appear to be almost absent in these regions (section \\ref{sec:mineralogie}).  \\begin{figure} \\centering \\includegraphics[width=0.944\\hsize]{flux_cumule.eps} \\caption{Cumulative received fluxes for the best model as a function of the distance from the star at a wavelength of 10\\,$\\mu$m (full line), 20\\,$\\mu$m (dashed line), 70\\,$\\mu$m (dotted line) and 1.3\\,mm (dot-dash line).\\label{fig:flux_cumule}} \\end{figure}  As already mentioned, the slope of the mm continuum and the mid-IR silicate features indicate a stratified structure for the disc, with large grains in the deeper parts of the disc and small grains ($\\lesssim 1\\,\\mu$m) in the surface of the disc. This surface layer of small grains remains difficult to characterise precisely. We have explored a structure where stratification of grain size is vertical and caused by  settling of larger particles towards the disc midplane. Our model with vertical settling allows us to accurately reproduce the SED, scattered light images and millimetre emission maps. However, we cannot assess the uniqueness of the solution and other explanations may be envisaged. The warm region emitting the 10 micron silicate feature is close to the star, while the mm/submm continuum comes from the whole volume of the disc. We may be observing small particles close in and big grains in the outer regions of the disc. However, most physical processes: grain growth, radial drift by gas drag in the disc midplane, radiation pressure or stellar wind, will preferentially result in larger grains in the central regions and/or remove small grains from these regions. Mid-infrared  interferometric observations (\\citealp{vanBoekel04} for HAeBe stars and \\citealp{Ratzka07} and \\citealp{Schegerer08} for the T~Tauri stars TW~Hydra and RY~Tau respectively) have confirmed these trends by showing the presence of larger and more processed grains at small spatial scales. Timescales for the production of large grains are significantly shorter in the central parts of the disc and it is difficult to imagine a physical process that will efficiently produce large grains in the outer disc without also producing them in the inner disc. Even if the current observations do not allow us to firmly conclude this point, vertical grain size stratification, probably resulting from a combination of settling and enhanced grain growth close to the midplane, appears to be a more natural explanation, and is, for now, our preferred model for the disc of IM~Lupi.   Following the previously described argument, the presence of millimetre grains at large radii strongly suggests that such large grains are also present in the central AU, where the process of grain growth should be more efficient. The 10\\,$\\mu$m features, dominated by the emission from grains of size around 1.5\\,$\\mu$m, indicate that the mixing in the disc is sufficient to maintain micrometric grains in the disc surface. Moreover, the absence of 3\\,$\\mu$m grains in the \\emph{IRS} spectrum indicates that the decoupling between gas and dust starts for a grain size between 1 and a few $\\mu$m, \\emph{i.e.} grains larger than this threshold are settled below the surface $\\tau_{10\\,\\mu\\mathrm{m}}$ = 1. Given the high densities in the inner regions of the disc, an increase of 1 in optical depth is obtained over a very small spatial scale. Even a moderate amount of dust settling is sufficient to produce the observed effect. The very low values we derive for the settling index confirm that the settling remains efficiently counterbalanced by vertical mixing, due to turbulence for instance.   It is very likely that the process of dust settling evolves with different timescales as a function of the distance from the star. The silicate emission bands provide strong indications of the efficiency of dust settling as a mean of removing grains larger than a few microns from the upper layers in the central parts of the disc. The SED also gives some insights on the presence of settling in the outer parts. The \\emph{MIPS} far-IR fluxes, which are low compared to the mid-IR emission, indicate that the disc intercepts a relatively small fraction of the stellar radiation at large radii ($> 10$\\,AU) compared to the fraction intercepted at radii $< 10$\\,AU probed in the mid-IR (see Fig.~\\ref{fig:flux_cumule}). This is expected in the case of dust settling (see for instance Fig.~7 in \\D04). The decreasing SED of IM~Lupi in the mid-infrared is very reminiscent of the models including dust settling of \\D04. Thus, we tried to fit the SED without taking into account the mid-infrared fluxes (between 5 and 35\\,$\\mu$m), and hence not considering the silicate features. The resulting Bayesian probabilities still exclude models without dust settling, which do not manage to reproduce simultaneously both the millimetre and far-infrared fluxes. We conclude that dust settling is very likely to occur even in the outer regions of the disc.    \\D04 predicts that settled discs may become undetectable in scattered light due to the formation of a self-shadowed opacity structure in the outer disc. This is clearly not the case for IM~Lupi. As noted by \\D04, the self-shadowed structure only appears for low values of the turbulence. This may indicate that the turbulence level in IM~Lupi is large enough to prevent the disc from self-shadowing or that we are observing the outer disc at a relatively early stage of the dust settling process.     \\subsection{Evolutionary state of IM~Lupi \\& comparison with other classical T~Tauri stars}  Our modelling has lead to a very detailed picture of the disc surrounding IM Lupi. We have already seen that the disc structure of IM Lupi is similar to those of other classical T~Tauri stars. In this section, we compare the signature of dust evolution in the disc with the results obtained for other T~Tauri discs, in order to determine  whether it is a singular object or whether it can be considered as representative of other T~Tauri stars.  Our conclusions on the degree of dust settling in IM~Lupi result from the strong silicate emission feature and shallow millimetre spectral slope. Other objects show similar characteristics. Thus, Fig.~\\ref{fig:beta} presents the strength of the 10\\,$\\mu$m silicate as a function of the exponent of the opacity law in the millimetre for a set of T~Tauri stars (Table~\\ref{tab:beta}). IM~Lupi is located in the bulk of T~Tauri stars and our results can probably be extrapolated to other sources: detailed analyses of the individual sources are required, but it is likely that at least some of them are undergoing some dust settling. A semblable conclusion was reached by \\cite{Furlan05,Furlan06} and \\cite{D'Alessio06} based on similar evidence \\emph{i.e.}, the presence in the SEDs of 10 and 18$\\,\\mu$m emission silicate bands and the slope and fluxes at  \\mbox{(sub-)millimetre} wavelengths.  \\begin{figure} \\includegraphics[width=\\hsize]{beta_s10.eps} \\caption{Strength of the 10$\\,\\mu$m silicate features as a function of the millimetre opacity slope for the T~Tauri stars listed in Table~\\ref{tab:beta}. IM~Lupi is represented by the red star. The effect of grain growth and dust settling are schematised by arrows.\\label{fig:beta}} \\end{figure}  \\begin{table} \\caption{Millimetre opacity slopes and strengths of the 10$\\,\\mu$m silicate features of T~Tauri stars. References: (1) = \\cite{Lommen07}, (2) = \\cite{Kessler-Silacci06}, (3) = \\cite{Przygodda03}, (4) = \\cite{Rodmann06}, (5) = \\cite{Furlan06}. The source T~Cha presented in \\cite{Lommen07} was not selected, because of the presence of PAH emission. Both \\cite{Rodmann06} and \\cite{Lommen07} take a contribution of optically thick emission into account in their derivation of $\\beta_\\mathrm{mm}$. \\label{tab:beta}} \\centering \\begin{tabular}{lccc} \\hline \\hline Name & $\\beta_\\mathrm{mm}$ & $S_\\mathrm{Peak}(10\\,\\mu\\mathrm{m})$ & Ref. \\\\ \\hline IM Lup   & 0.8  $\\pm$    0.25   &   1.60 &   this work   \\\\ \\hline HT Lup   & 0.4 $\\pm$     0.5\t&   1.39 &   (1, 2)\\\\ GW Lup\t & 0.5 $\\pm$\t 0.5\t&   1.48 &   (1, 2)\\\\ CR Cha   & 1.5 $\\pm$\t 0.6    &   2.52 &   (1, 3)\\\\ WW Cha   & 0.8 $\\pm$\t 0.8    &   1.92 &   (1, 3)\\\\ RU Lup   & 0.8 $\\pm$\t 0.5    &   1.54 &   (1, 2)\\\\ RY Tau   & 0.8 $\\pm$\t 0.1    &   2.75 &   (4, 5)\\\\ FT Tau   & 0.9 $\\pm$\t 0.3    &   1.74 &   (4, 5)\\\\ DG Tau   & 0.7 $\\pm$\t 0.1    &   1.52  &  (4, 5)\\\\ UZ Tau E & 0.8 $\\pm$\t 0.1    &   1.65 &   (4, 5)\\\\ DL Tau   & 1.0 $\\pm$\t 0.2    &   1.1 &    (4, 5)\\\\ CI Tau   & 1.3 $\\pm$\t 0.4    &   1.5 &    (4, 5)\\\\ DO Tau   & 0.5 $\\pm$\t 0.1    &   1.19 &   (4, 5)\\\\ GM Aur   & 1.6 $\\pm$\t 0.2    &   2.54 &   (4, 5)\\\\ \\hline \\end{tabular} \\end{table}  The tentative correlation between the strength of the silicate feature and millimetre spectral index suggested by \\cite{Lommen07} does not appear as clear in our larger sample of T~Tauri stars. This correlation was interpreted as an indication of fast grain growth in both central and outer regions of the disc. Indeed, as grains grow from sub-micron sizes to several microns, the $10\\,\\mu$m feature becomes weaker and less peaked. When they reach millimetre to centimetre sizes, the slope of the millimetre emission becomes shallower. But, as we have shown, the strength of the silicate features is strongly related to the degree of dust settling. The stronger the settling, the smaller the apparent (\\emph{i.e.} probed in the infrared) grain size, making it difficult to obtain a precise estimate of the actual grain sizes present in the midplane.   The detailed analysis of the silicate features indicates a small degree of crystallisation  ($< 10\\,\\%$) in spite of a high mass fraction of micrometric (hence evolved) grains. As shown by \\citet[see their Fig.~9 for instance]{Schegerer06}, this is the case for several objects and IM~Lupi is not unique on that aspect. \\cite{Schegerer06} did not find any correlation between the degree of crystallisation and amount of micron-sized grains.   \\cite{Furlan05,Furlan06} and \\cite{D'Alessio06} have estimated the degree of dust settling based on the colours and SEDs of a large population of Classical T Tauri stars in the Taurus molecular cloud. They claim that to make a synthetic SED consistent with  the median SED of class II objects in Taurus, the dust to gas mass ratio in the disc atmosphere should be around 1\\,\\% (between 0.1 and 10\\,\\%) of the ISM ratio. In our modelling, the grain size distribution and the dust to gas ratio are continuous functions of the height above the midplane. With $\\xi = 0.05$, the dust to gas ratio starts at 1.5 times the ISM value in the disc midplane and reaches 10\\,\\%, 1\\,\\% and 0.1\\,\\% of the ISM value at 2, 3 and 6 scale heights, respectively. These regions roughly correspond to what is defined as the disc surface in \\cite{Furlan05,Furlan06} and \\cite{D'Alessio06}. Although not directly comparable, our estimation of the degree of dust settling appears consistent with the calculations of these authors."
    },
    "0808/0808.3040_arXiv.txt": {
        "abstract": "Here I discuss possible relations between free precession of neutron stars, Tkachenko waves inside them and glitches. I note that the proposed precession period of the isolated neutron star RX J0720.4-3125  (Haberl et al. 2006) is consistent with the period of Tkachenko waves for the spin period $8.4$~s. Based on a possible observation  of a glitch in RX J0720.4-3125 (van Kerkwijk et al. 2007), I propose a simple model, in which long period precession is powered by Tkachenko waves generated by a glitch. The period of free precession, determined by a NS oblateness, should be equal to the standing Tkachenko wave period for effective energy transfer from the standing wave to the precession motion. A similar scenario can be applicable also in the case of the PSR B1828-11. ",
        "introduction": "Isolated neutron stars (NSs) being non-spherical bodies are expected to demonstrate free precession (for a brief review see, for example,  \\citealt{l2003}). However, examples of this phenomena are less than few, and even in rare cases when a precession-like behavior is observed different interpretations can be discussed (even not related to precession, see for example \\cite{rg2006}).  The problem of long period free precession in NSs is a long standing one. A NS can precess if it is non-spherical and rotation axis does not coincide with a principal axis. Typically, biaxial objects are discussed, so deviation from spherical symmetry can be described by one parameter -- oblateness (see \\cite{alw2006} for a discussion of triaxial model). Expected values of NS oblateness (due to rotation or influence of strong magnetic fields) can naturally lead to precession periods about one year. The precession period is equal to $P_\\mathrm{prec}=P/\\epsilon$. Here $P$ is the spin period of a NS, $\\epsilon$ -- its oblateness, and $P_\\mathrm{pres}$ -- precession period. Measured precession periods require oblateness about $10^{-8}$.  However, discussing dynamics of NSs it is necessary to take into accout the network of superfluid vortices inside them. The neutron superfluid liquid in the interior of a NS participates in rotation via formation of quantized vortex lines. The density of these lines per unit area is $n=2\\Omega / k$. Here $\\Omega=2\\pi /P$ is spin frequency, and $k=h/2m_\\mathrm{n}$, where $m_\\mathrm{n}$ is a neutron mass (see, for example, \\citealt{st1983}, Ch. 10). The vortices exist in the core of a NS, where they can interact with superfluid (superconducting) protons and normal electrons, and in the crust, where they can pin to it.  Coupling of superfluid neutron vortices with electrons in a core results in damping of free precession \\citep{ao1987}. But the time scale of this damping is long enough, according to these authors. For spin period about 1 second it is $\\sim 400$~--~$10^4$~$P_\\mathrm{prec}$ \\citep{ao1987}. Still, this time scale is much shorter than a NS age, so some excitation mechanism is necessary for precession. As it is discussed below, in the presented model excitation is due to a glitch.  A kind of pinning (``immobilization'') of vortices can also happen in the core due to interactions with magnetic flux tubes (see discussion, for example, in \\citealt{l2007}). In this case, the moment of inertia of ``pinned'' neutrons (which is about $I$) is about 10 times larger, than the moment of the remaining parts of a NS, $I_\\mathrm{c}$. So, $P_\\mathrm{prec}\\sim 0.1 P$.  A different kind of problem appears if pinning in the crust is taken into account. For absolute pinning no long period precession is possible. Instead, the period of precession becomes equal to $P (I/I_\\mathrm{p})$, here $I$ is NS moment of inertia, $I_\\mathrm{p}$ is the moment of inertia of pinned superfluid in the crust. Typically it is expected that $I_\\mathrm{p}/I \\sim 10^{-2}$ \\citep{s1977}, and the precession period is just $\\sim100$~$P$ if the absolute pinning is valid. However, \\cite{ao1987} showed that this is not the case due to finite temperature. Because of thermal effect always there is vortex creep which allow the pinned superfluid to follow precession.  The best example of a NS with precession-like behavior is PSR B1828-11. The proposed period is about 511 days with a harmonic at 256 days \\citep{sls2000}. Most of discussions related to free  precession deal with this source. In particular, the problem of non-existence of long period precession for strong pinning is typically confronted with observaions of PSR B1828-11.  Recently, appeared another possible example of long period free precession in NSs. The existence of $\\sim 7$ years precession period in one of a small group of isolated  NSs (called XDINS -- X-ray Dim Isolated NSs, or ICoNS -- Isolated Cooling NSs, or the {\\it Magnificent Seven}) -- RX J0720.4-3125 -- was suspected \\citep{hetal2006}\\footnote{A slightly different period $\\sim$ 4.3 years was proposed by van Kerkwijk and Kaplan (2007) based on timing analysis. However, these authors consider a model with a glitch to fit better due to a rapid change in spectral properties, see van Kerkwijk et al. (2007).}. So, this object was added to the list, and the paradoxical situation of long precession in presence of superfluid vortices was reconsidered by \\cite{l2006,l2007}. This author proposed that either protons in NS interiors are type I superconductors, or neutrons in the outer core are normal (i.e., not superfluid). More recently \\cite{gaj2007} demonstrated that for long spin period and small precession angles NSs can have long precession periods (note, that for PSR B1828-11 the precession angle is proposed to be small, about few degrees \\citep{sls2000}, but for RX J0720.4-3125 it can be larger, $>$~10 degrees \\citep{hetal2006}). So, according to \\cite{gaj2007}, the conclusion by \\cite{l2006} and other authors that in the strong drag regime $P_\\mathrm{prec}\\sim 0.1 P$  can be under doubt due to a short wavelength  instability.  Clearly, the problem of free precession in NSs is far from being solved completely. In this brief note, based on coincidence between Tkachenko wave period and precession period in cases of PSR B1828-11 and RX J0720.4-3125, I discuss a mechanism to support precession in isolated NSs. ",
        "conclusions": "Absence of free precession in absolute majority of isolated NSs indicates that this phenomena needs some rare coincidence in properties of a NS. Here it is proposed that it is necessary to have: \\begin{itemize} \\item $P_\\mathrm{T}\\approx P_\\mathrm{prec}$, \\item a glitch to generate Tkachenko waves. \\end{itemize}   Instead of the proposal by \\cite{l2007} -- {\\it \"A slowly-precessing neutron star cannot glitch\"} -- I propose another: {\\it slow precession is powered by glitches via Tkachenko waves}.  Observations  of RX J0720.4-3125 are roughly consistent with this scenario. On the other hand, in the case of PSR B1828-11 no glitches have been observed. However, it is necessary to study for how long precession can survive after a glitch. If an old estimate by \\cite{ao1987}, $400$~--~$10^4\\, P_\\mathrm{prec}$ is valid, then this time is long enough. If precession is periodically excited by glitches via Tkachenko waves even damping on a time scale of few precession cycles \\citep{l2006} would not contradict observations of RX J0720.4-3125.  The glitch in RX J0720.4-3125 reported by \\cite{vketal2007} by its consequencies is similar to the one observed in an anomalous X-ray pulsar (AXP) CXOU J164710.2-455216 \\citep{ietal2007,metal2007}. After the glitch the luminosity of the source was increased, and its spectrum changed. So, the jump in properties of the spectrum and luminosity of RX J0720.4-3125 proposed by \\cite{vketal2007} can be directly related to a glitch, which is weaker than in the case of CXOU J164710.2-455216 (still similarities in behavior of these sources can be considered as a kind of support to the hypothesis of a link between AXPs and XDINS). But the evolution of the NS parameters after the ``jump'' requires precession.  Note, that the timing solution before MJD 52821, when a possible \"glitch\" happened according to \\cite{vketal2007}, can be relatively well described by the so-called cubic solution (the second derivative of $\\nu$ is non-zero), see \\cite{vkk2007}. Spectral changes before this date are not very large \\citep{hetal2006, vkk2007}. After MJD 52821 the timing solution is well described by a periodic function, see \\cite{vkk2007} (these authors studied several models with and without a glitch in their two papers), and spectral changes follow this law, too \\citep{hetal2006}\\footnote{However, van Kerkwijk and Kaplan (2007) propose that the period is not close to $\\sim$7 years, but is $\\sim$4.3 years. Still, on a short time scale -- since MJD 52821 -- this is not very certain.}. Based on that, I suggest that the timing residuals might be also explained by a model without (or with small) precession before the glitch, and strong precession after. However, this particular model has never been tested quantitatively against observational data.  Glitches naturally can produce thermal afterglows \\citep{hetal1997}. About $10^{38}$~--~$10^{43}$~ergs can be released in a glitch (in the case of RX J0720.4-3125 according to estimates of the increase in spin frequency by \\citealt{vketal2007} this value is closer to $10^{38}$~erg). However, Hirano et al. showed that a thermal response of a NS to a glitch cannot produce a smooth temperature increase on the time scale of years. If surface temperature is increased just by few percent, as it is required by \\cite{vketal2007}, then the brightening lasts just for few days (this corresponds to weak energy release). If we require a temperature rise for a long time, then the effect is too strong \\citep{hetal1997}. So, I conclude that spectral changes on a long time scale should be attributed to precession of the NS.   Glitches of AXPs (and soft gamma-repeaters) can be different  in nature with respect to radio pulsar glitches, as the former can be related to crust fracture due to superstrong magnetic field. Still, the origin of a glitch is not important for our discussion here. ``Normal'' glitches are quite common for long period pulsars, for example, PSR J1814-1744 with spin period about 4 seconds demonstrated a glitch \\citep{js2006}.   So, RX J0720.4-3125 can glitch not only via the mechanism operating in magnetars, but also due to convenient mechanisms proposed for normal radio pulsars. For them one can estimate the reccurence time following \\cite{ab1994}.  If the glitch of RX J0720.4-3125 is due to unpinning, then using standard formulae \\citep{ab1994} one obtains that the reccurence time between two succesive glitches is:  \\begin{equation} t_{\\mathrm{g}}=\\frac{\\delta \\Omega}{\\dot \\Omega}. \\end{equation} Parameter $\\delta \\Omega$ is the critical value of the difference between the rotation frequencies of normal matter and the superfluid at a boundary layer. $\\delta \\Omega$ itself can be estimated as (Alpar \\& Baykal 1994):   \\begin{equation} \\delta \\Omega = \\Delta \\Omega \\frac{I}{I_{\\mathrm{p}}}, \\end{equation} here $I_{\\mathrm{p}}$ is the effective moment of inertia of the region of a pinning layer.  Combining these two formulae one obtains the relation for the time between glitches:   \\begin{equation} t_{\\mathrm{g}} = \\frac{2\\, I}{I_{\\mathrm{p}}} \\frac{\\Delta \\Omega}{\\Omega} t \\propto t, \\label{eq2} \\end{equation} where $t$ is the age of a pulsar, $t=\\Omega/2\\, \\dot \\Omega$.   Surprisingly, the time is about 10 years for RX J0720.4-3125. I.e. it is quite probable to observe one since the discovery of this object. Then, one can expect to see similar phenomenae in other Magnificent seven objects. However, they are less studied, and may be some  glitches are missed.  Still, to produce  luminosity and spectral changes, and to generate Tkachenko waves, it is probably more natural to have a glitch due to a starquake. In the case of the quake model \\citep{ab1994} the time between glitches is longer, about 300 years for RX J0720.4-3125.  \\begin{equation} t_{\\mathrm{g}} = \\frac{2(A+B)\\phi \\Delta \\Omega / \\Omega} {I_0\\Omega \\dot \\Omega}. \\end{equation} The estimate above was obtained assuming standard values \\citep{ab1994} $A=10^{52}$ erg, $\\phi=10^{-3}$, $B=10^{48}$~erg,  and $I\\sim 10^{45}$~g~cm$^2$   Then, we can be just lucky to find a glitch in $\\sim 10$ years of observations (but note, that it is not the only XDINS observed). Or, in XDINS quakes do not follow the formula for radio pulsars.  \\cite{vketal2007} proposed that an accretion episode can be responsible for spectral changes after a ``glitch'' in RX J0720.4-3125. I think that this is not a very probable reason for the origin of the glitch and corresponding changes.  It is hardly possible to imagine that if we observe such an episode just after 10 years of observations, other episodes were not frequent during the evolution of this source. With frequent episodes of accretion of light elements a NS should follow a slightly different cooling track \\citep{ketal2006}. Such stars with accreted envelopes are hotter in their youth, but colder after they mature. This is more similar to the properties of 1E1207.4-5209 and Kes 79 \\citep{gh2007}.  A remarkable difference between some NSs in supernova remnants (so-called CCOs) and XDINS in the solar vicinity can be due to the existence of accreted envelopes in the former. Note, that we do not observe descendants of 1E1207-like sources in our proximity. If they follow a standard cooling curve, like RX J0720.4-3125 and other XDINS, then they have to be observed. On the other hand, we do not see ancestors of XDINS in supernova remnants. This can be related to the fact that CCOs descendants are too cold at the age of XDINS to be easily detected by ROSAT, and vice versa ancestors of XDINS are not hot enough to be easily found in some supernova remnants.  In this note I neglect (as most, if not all, other authors who studied Tkachenko waves in NSs) the influence of interaction between neutron vortex lines and magnetic flux tubes. This interaction can significantly affect the velocity of waves, and so their period, and to damp them.\\footnote{This was noted to us by Prof. M. Ruderman.} The velocity of a Tkachenko wave can be estimated as:  \\begin{equation} V_\\mathrm{T}= 1/2 (h \\Omega/2\\pi m_\\mathrm{pair})^{1/2}= (\\Omega k/8 \\pi)^{1/2} \\sim b/P. \\end{equation} Here $m_\\mathrm{pair}=2 m_\\mathrm{n}$, and $b$ is the distance between vortex lines. So, $P_\\mathrm{T}\\sim (R/b) P$ in the simple case when there are no interactions with flux tubes or other complications. When vortices are able to ``communicate'' with the help of numerous magnetic flux tubes, the velocity of the wave can be larger, so the period of Tkachenko-like wave would be shorter. This question should be explored.  To conclude, in this short note I proposed that long period precession in RX J0720.4-3125 can be related to Tkachenko waves, generated in a recent glitch. In the proposed model before a glitch precession could be negligible. The critical condition for the free precession excitation is the  equality between Tkachenko wave period and the period of free precession."
    },
    "0808/0808.1045_arXiv.txt": {
        "abstract": "We discuss the main results that were recently published by the Auger Collaboration and their impact on our knowledge of the ultra high energy cosmic rays and neutrinos. ",
        "introduction": "We have discussed the problems related to the ultra high energy cosmic rays (UHECR) quite many times in recent years after the results of the high resolution Fly's Eye (HiRes) started trickling out several years ago. The cosmic ray energy spectrum extracted from the the HiRes observations indicated that the spectrum suffers a strong decrease at energies above 5$\\times$10$^{19}$ eV~\\cite{HiRes_sp}, consistent with the GZK cutoff~\\cite{GZK} that is due to UHECR interactions with the microwave background. There was thus a strong disagreement with the energy spectrum derived from the previous largest exposure cosmic ray air shower array, Agasa~\\cite{Agasa} that seemed to continue with the same slope above 10$^{20}$ eV.  HiRes and Agasa measure the cosmic ray events in different ways. HiRes is a fluorescent detector that follows the shower longitudinal development by detecting the fluorescent light emitted by the Nitrogen atoms in the air ionized by the shower charged particles. The primary cosmic ray energy is obtained from integration of the shower longitudinal profile normalized by the average electron energy loss and the fluorescent efficiency of the air. Agasa, on the other hand, is a classical air shower array that records the particle density at the observation level and uses Monte Carlo calculations to relate it to the primary cosmic ray energy. The dependence of this method on the hadronic interaction model used in the simulations is considered stronger in this approach.  In 2007 the first results of the Pierre Auger Observatory (PAO) in Argentina were published. This experiment became the hope for better understanding of UHECR several years ago when its construction started. It is a hybrid air shower experiment that applies both detection methods. PAO consists of an air shower array of area 3,000 km$^2$ where the individual detectors are 10 m$^2$ water tanks placed at distances of 1.5 km in a hexagonal pattern (SD). The shower array is surrounded by 24 fluorescent telescopes (FD) located in four groups of six. The idea is that during cloudless moonless nights when the fluorescent detectors can work the UHECR showers will be observed simultaneously by both types of detectors and this {\\it hybrid} observation will unveil better the shower properties. Here is a brief summary of the Pierre Auger Observatory results.  The energy spectrum derived from PAO data~\\cite{auger_sp} shows a feature consistent with the GZK cutoff similar, but not identical to that of HiRes - see Fig.~\\ref{AugerHR}. The spectrum is based on the much higher surface detector statistics normalized to the energy estimates of the fluorescent detector in hybrid events. The problem is that if the energy was estimated only from the SD data (particle density at 1,000 meters from the shower core - $S_{1000}$) the energy scale would be 25\\% higer~\\cite{auger_re}. Formally this is not a big problem because the uncertainty of the energy estimate is 22\\%.  Another problem appears when PAO attempts to estimate the type of the primary cosmic ray nuclei. When it is estimated from the depth of the shower maximum measured by the fluorescent detector $X_{max}$ the average primary mass appears to be close to He~\\cite{auger_mu} . If the surface detector density $S_{1000}$ is used to estimate the number of shower muons~\\cite{auger_re} the average primary nuclei appear to be much heavier. The preliminary conclusion from these two disagreements is that showers are absorbed in the atmosphere slower than models predict.  The Auger Collaboration succeeded, however, to eliminate a class of UHECR models that predicted that the primary particles are gamma rays. From studies of $X_{max}$ and other shower properties the fraction of gamma rays in the cosmic rays above 10$^{19}$ eV is limited to not more than 2\\%~\\cite{auger_g}. The much smaller statistics at higher energy does not allow for such strong limits, but at least the bulk of UHECR, which we consider extra galactic, consists of nuclei.  The Auger result that impressed the most not only the astrophysical community, but also the general public, was the published correlation of their highest energy events with active galactic nuclei~\\cite{Auger1,Auger2}. Twenty of the 27 events of energy above 5.7$\\times$10$^{19}$ eV (57 EeV) come within 3.1$^o$ of the position of AGN at redshift smaller than $\\leq$0.017 (distance of 71 Mpc) while one expects only 7 in case of isotropic source distribution. The 3.1$^o$ circle around the event direction (that was obtained from a scan in energy, angle and distance) is understood as scattering in the galactic magnetic field plus about 1$^o$ Auger angular resolution. Accounting for the number of scans the chance probability is well below 1\\%. If the events that come from galactic latitude less than 10$^o$ (where catalogs are not full and cosmic ray may scatter more) are excluded then 19 out of 21 events are closer than 3.1$^o$ from nearby AGN.  Since the Auger collaboration published the arrival directions and energies of these 27 events in Ref.~\\cite{Auger2} this announcement was quickly followed by a number of other analyses and questions, like:\\\\ \\hspace*{5mm} $\\bullet$ why only 71 Mpc when $>$60 EeV cosmic rays should come from  distances up to 200 Mpc ?\\\\ \\hspace*{5mm} $\\bullet$ why there are no events from the direction of the Virgo cluster, that contains powerful AGN ?\\\\ \\hspace*{5mm} $\\bullet$ what does the concentration of events from directions close to Cen A means ?\\\\ Many of these questions were also asked by the Auger Collaboration~\\cite{Auger2}. ",
        "conclusions": "The first results of the Auger Collaboration did improve our knowledge of the ultra high energy cosmic rays. They are fully consistent with a GZK feature in the cosmic ray spectrum as suggested previously by the HiRes Collaboration. Being a hybrid experiment Auger demonstrated that the energy estimate by fluorescent detectors and surface air shower array is different, as previously suggested by the Agasa and HiRes spectra. Now the UHE community is searching for the reasons for this disagreement that may hide in the currently used hadronic interaction models.  Auger started the direct search for the sources of UHECR by studies of the correlation of their highest energy events with active galactic nuclei. The published correlation inspired many other different analyses of correlation with powerful astrophysical objects. This work is in progress now and the UHECR sources will be identified eventually when the Auger statistics grows - it will double at the end of August 2008.  Since the current UHECR energy spectrum can be fit with several different models it is not obvious that we know better the expected flux of cosmogenic neutrinos. This may only happen when an unique fit of the spectrum is achieved.\\\\[3truemm] {\\bf Acknowledgments} This talk is based on collaboration with many colleagues including Peter Biermann, Daniel DeMarco, Thomas Gaisser, J\\\"{o}rg Rachen, \\& David Seckel. The author was supported in part by NASA APT grant NNG04GK86G."
    },
    "0808/0808.3781_arXiv.txt": {
        "abstract": "{Nonthermal radiation observed from astrophysical systems containing relativistic jets and shocks, e.g., gamma-ray bursts (GRBs), active galactic nuclei (AGNs), and Galactic microquasar systems usually have power-law emission spectra. Recent PIC simulations of relativistic electron-ion (electron-positron) jets injected into a stationary medium show that particle acceleration occurs within the downstream jet. In the presence of relativistic jets, instabilities such as the Buneman instability, other two-streaming instability, and the Weibel (filamentation) instability create collisionless  shocks, which are responsible for particle (electron, positron, and ion) acceleration. The simulation results show that the Weibel instability is responsible for generating and amplifying highly nonuniform, small-scale magnetic fields. These magnetic fields contribute to the electron's transverse deflection behind the jet head. The ``jitter'' radiation from deflected electrons in small-scale magnetic fields has different properties than synchrotron radiation which is calculated in a uniform magnetic field. This jitter radiation, a case of diffusive synchrotron radiation, may be important to understand the complex time evolution and/or spectral structure in gamma-ray bursts, relativistic jets, and supernova remnants. }  \\FullConference{Workshop on Blazar Variability across the Electromagnetic Spectrum\\\\ April 22-25 2008\\\\ Palaiseau, France}  \\begin{document} ",
        "introduction": "Shocks are believed to be responsible for prompt emission from gamma-ray bursts (GRBs) and their afterglows, for variable emission from blazars, and for particle acceleration processes in jets from active galactic nuclei (AGN) and supernova remnants (SNRs). The predominant contribution to the observed emission spectra is often assumed to be synchrotron- and inverse Compton radiation from these accelerated particles for gamma-ray bursts [1-7] and for AGN jets [8-13]. It is assumed that turbulent magnetic fields in the shock region lead to Fermi acceleration, producing higher energy particles \\cite{fermi49,blaneich97}. To make progress in understanding emission from these object classes, it is essential to place modeling efforts on a firm physical basis. This requires studies of the microphysics of the shock process in a self-consistent manner \\cite{p05b,wax06}. ",
        "conclusions": "We have started to calculate emission directly from our simulations using the same method described in the previous section. In order to calculate the (jitter-like) synchrotron radiation from the particles in the electromagnetic fields generated by the filamentation instability, the retarded electric field from a single particle is Fourier-transformed to give the individual particle spectrum as described in the previous section. The individual particle spectra are added together to produce a total spectrum over a particular simulation time span \\cite{hedeT05,hedeN05}. It should be noted that for this calculation very large simulations over a long time ($t_{\\rm s}$) are required using a small time step ($\\Delta t$) in order to increase the upper frequency limit to the spectrum (Nyquist frequency $\\omega_{\\rm N} = 1/2\\Delta t$). Frequency resolution is limited by the time span ($\\Delta \\omega = 1/t_{\\rm s}$) \\cite{hedeT05,hedeN05}. For a case with the time step $\\Delta t = 0.01/\\omega_{\\rm pe}$ and the time span $t_{\\rm s} = 50/\\omega_{\\rm pe}$, a calculated spectrum will have the highest frequency, $50\\omega_{\\rm pe}$ and the frequency resolution (the lowest frequency), $0.02\\omega_{\\rm pe}$.  $\\omega_{\\rm pe}$ is calculated with an appropriate plasma density.  Simulations over a long time allow us to obtain multiple spectra at sequential time spans so the spectral evolution can be calculated. Synthetic spectra obtained in the way we have described above should be compared with GRB prompt and afterglow observations.  In the case of AGN jets, diffusive synchrotron radiation has already been invoked by several works \\cite{flei06b,FT07b,mao07} to reproduce spectra of 3C273, M87 and Cen A knots from radio to X-rays. For TeV blazars, taking into account the relative importance of the energy densities contained in the small-scale and large-scale magnetic fields may be an elegant alternative to the choice of a broken power-law for the energy distribution of radiating electrons."
    },
    "0808/0808.0261_arXiv.txt": {
        "abstract": "Conditions are obtained for the existence of a warm inflationary attractor in the system of equations describing an inflaton coupled to radiation. These conditions restrict the temperature dependence of the dissipative terms and the size of thermal corrections to the inflaton potential, as well as the gradient of the inflaton potential. When these conditions are met, the evolution approaches a slow-roll limit and only curvature fluctuations survive on super-horizon scales. Formulae are given for the spectral indices of the density perturbations and the tensor/scalar density perturbation amplitude ratio in warm inflation. ",
        "introduction": "Inflationary models \\cite{guth81,linde82,albrecht82} have proved very sucessful in explaining many of the large scale features of the universe (see e.g. \\cite{Liddle:2000cg}). An essential feature of these inflationary models is their stability, meaning in particular that inflation solutions are attractors in the solution space of the relevant cosmological equations (see e.g. \\cite{Salopek:1990jq,Liddle:1994dx}), at least up until inflation ends and the universe becomes radiation dominated. Without this feature, inflation might never have begun, and certainly would not have lasted long enough to affect the large scale structure of the universe.  Warm inflation is an alternative inflationary scenario in which a small but significant amount of radiation survives during the inflationary era due to continuous particle production \\cite{Moss85,bererafang95,berera95}. The coupling between radiation and the inflaton field leads to thermal dissipation and fluctuations in the time evolution of the inflaton field. The stability of the inflationary solutions in warm inflationary models has only been addressed in a limited form previously \\cite{deOliveira:1997jt}, and we will present the full stability analysis here. We shall examine conditions under which warm inflation is an attractor and give conditions for a prolonged period of warm inflation.  We shall show that the stability of warm inflation can be related to conditions on two parameters describing the temperature dependence of terms in the inflaton equation of motion which where not taken into account in the earlier stability analysis \\cite{deOliveira:1997jt}. The first condition says that if the dissipation term in the equation of motion falls off too rapidly at low temperature then the temperature is driven to zero, and we fall into the conventional inflationarty scenario. The second condition limits the temperature dependence of the inflaton potential. In many models large dissipation implies large thermal corrections to the potential which prevent warm inflation taking place. This argument is essentially the one first presented in Ref. \\cite{yokoyama99}. Nowadays, we know that there are models where thermal corrections to the inflaton potential are suppressed by supersymmetry,  and these models may allow warm inflation as an attractor \\cite{BasteroGil:2006vr,BuenoSanchez:2008nc,review}.  The duration of the period of inflation is related to a set of slow-roll paramaters which where introduced in \\cite{Hall:2003zp}. We shall re-derive the slow-roll conditions for warm inflation as part of the stability analysis. A well understood feature of the slow-roll conditons for warm inflation is that they can be less restrictive than the slow-roll condtions for conventional inflation.  We stress that we are concerned here with the self-consistency of the warm inflationary scenario for given equations of motion. We shall not address how the equation of motion for the inflaton field is obtained from non-equilibrium thermal field theory. A discussion of the derivation of the equations of motion can be found in a recent review \\cite{review}. However, we would like to point out that some of the critisims of warm inflation have been based on models which do not satisfy the fundamental stability conditions derived here, and are therefore not inconsistent with the validity of warm inflation in general \\cite{Aarts:2007ye}.  The stability of warm inflation has consequences for the origin and evolution of cosmological density fluctuations \\cite{Hall:2003zp}. In warm inflation, density fluctuations originate from thermal fluctuations \\cite{bererafang95,berera00}. In particular, the fact that the inflationary solution depends on only one parameter means that only one perturbation mode, the curvature perturbation, survives on super-horizon scales despite the fact that there are entropy perturbations present on sub-horizon scales. We shall give formulae for the spectral indices of the scalar and tensor modes and say a little about the tensor/scalar ratio. ",
        "conclusions": "We shall recapitulate the main points of this paper. There are conditions on six of the parameters defined in Sect. \\ref{be} for the possibility of a stable period of warm inflation in the early universe: \\begin{itemize} \\item The parameter $Q$ which measures the strength of the friction term must satisfy \\begin{equation} Q>g_*{\\cal P}_S. \\end{equation} where $g_*$ is the effective particle number and ${\\cal P}_S$ is the scalar perturbation power spectrum on large scales. \\item The parameters which describe the inflaton dependence of the effective potential and friction term satisfy \\begin{equation} \\epsilon<<1+Q,\\qquad |\\eta|<<1+Q,\\qquad |\\beta|<<1+Q. \\end{equation} \\item The temperature dependence of the potential and the friction term is restricted by \\begin{equation} |b|<<{Q\\over 1+Q},\\qquad |c|<4. \\end{equation} The condition on $b$ implies that warm inflation is only possible when a mechanism, such as supersymmetry, reduces the size of the thermal corrections to the potential. \\end{itemize} Models of elementary particles exist in which these conditions can be satisfied. The most convincing of these models use a combination of supersymmetry and a two-stage decay process, where there are no direct coupling between the inflaton and the radiation and all thermal effects are supressed by factors of $T/\\Lambda_S$, where $\\Lambda_S$ is the supersymmtry breaking scale \\cite{berera02,BasteroGil:2006vr,BuenoSanchez:2008nc}.  When the conditions listed above are satisfied, then the solutions to the equations of motion approach a slow-roll approximation during inflation. As a result, large scale density perturbations have only one degree of freedom, which we identify as the curvature perturbation. (Entropy perturbations can only be introduced by adding extra degrees of freedom to the system.)"
    },
    "0808/0808.1737_arXiv.txt": {
        "abstract": "We present a new analysis of the long-period variables in the Large Magellanic Cloud (LMC) from the MACHO \\emph{Variable Star Catalog}. Three-quarters of our sample of evolved, variable stars have periodic light curves. We characterize the stars in our sample using the multiple periods found in their frequency spectra. Additionally, we use single-epoch Two Micron All Sky Survey measurements to construct the average infrared light curves for different groups of these stars. Comparison with evolutionary models shows that stars on the red giant branch (RGB) or the early asymptotic giant branch (AGB) often show non-periodic variability, but begin to pulsate with periods on the two shortest period-luminosity sequences (3 \\& 4) when they brighten to $K_s \\approx 13$. The stars on the thermally pulsing AGB are more likely to pulsate with longer periods that lie on the next two P-L sequences (1 \\& 2), including the sequence associated with the Miras in the LMC. The Petersen diagram and its variants show that multi-periodic stars on each pair of these sequences (3 \\& 4, and 1 \\& 2) typically pulsate with periods associated only with that pair. The periods in these multi-periodic stars become longer and stronger as the star evolves. We further constrain the mechanism behind the long secondary periods (LSPs) seen in half of our sample, and find that there is a close match between the luminosity functions of the LSP stars and all of the stars in our sample, and that these star's pulsation amplitudes are relatively wavelength independent. Although this is characteristic of stellar multiplicity, the large number of these variables is problematic for that explanation. ",
        "introduction": "\\label{intro} Stellar pulsation of giant stars appears to be a ubiquitous and important phenomenon---RR Lyrae and Cepheid variables form the basis for the distance scales we use. Miras and other long-period variables (LPVs) however, are not as well understood, largely because their cool and tenuous atmospheres are dynamic environments with a great diversity of molecular species forming and disassociating as the star pulsates. In recent years however, these stars have attracted increased attention as micro-lensing surveys of the Large and Small Magellanic Clouds (OGLE --- \\cite{1994ApJ...435L.113P}; OGLE II --- \\cite*{1997AcA....47..319U}; MACHO --- \\cite{macho}) have  produced large catalogs of LPVs. Well-sampled light curves and excellent photometry give us an opportunity to better understand both the mechanisms behind long-period variables and the physical processes at work in the latest stages of stellar evolution.  Before the wealth of data from micro-lensing surveys, LPVs were traditionally classified by the amplitude and stability of their variability in the $V$ band (e.g. The General Catalog of Variable Stars---\\cite{GCVS}). In this scheme, stars with well-defined pulsation are classified as Miras if the amplitude of their variability exceeds 2.5 magnitudes in $V$, and as Semi-Regular Type a (SRa) stars if not. Those with multiple periods, unstable periodicity, or poorly expressed periodicity, are classified as SRb stars. The MACHO Project's survey of the Large Magellanic Cloud (LMC) revealed five parallel sequences of LPVs in period-luminosity space \\citep{cook96}, prompting a classification scheme that uses the period of pulsation as its primary discriminator. \\cite{Wood99} identified the cause of the first three period-luminosity sequences (denoted \\emph{A}, \\emph{B}, and \\emph{C}) as pulsation, and suggested that the two longest period sequences (\\emph{E} and \\emph{D}) could be attributed to binary systems. The stars in Sequence E showed the characteristic light curves of contact binary systems, and Sequence D stars---those with the longest periods---simultaneously exhibited at least one shorter period that was coincident with Sequence B. This is likely to be the LMC equivalent of the ``long secondary periods'' described by \\cite{1963AJ.....68..253H} for Galactic LPVs, although the periods that comprise Sequence D are on average three times shorter than the LSPs listed in \\cite{1963AJ.....68..253H}. \\cite{Wood99} proposed that these stars are composed of accreting binary systems, with the long period caused by partial eclipses due to an unseen, dust-enshrouded companion.  In this work the LMC period-luminosity sequences will be named in the manner of \\cite{fraser05}, from shortest to longest period: 4, 3, 2, 1, E, and D. We retain the names D and E from \\cite{Wood99}, but rename his Sequence C to Sequence 1 and count up toward the shorter periods. This approach provides for a graceful way to accommodate additional short-period sequences. Indeed, the use of 2MASS $K_s$ magnitudes as the luminosity indicator caused a split in the original Sequence B---producing Sequences 2 and 3; \\cite{2003MNRAS.343L..79K}. A fifth sequence was identified by \\cite{2004AcA....54..129S} through the examination of all significant frequencies of these stars instead of just the strongest frequency.  In general, stars brighten and redden as they evolve along the Red Giant Branch (RGB) and the Asymptotic Giant Branch (AGB). Since the typical $J-K_s$ color of the stars in each sequence reddens as we progress from the short-period Sequence 4 to the longer period Sequence 1, this suggests that evolution proceeds from shorter periods toward longer periods, at least for Sequences 1--4 \\citep{fraser05}. In fact, the low luminosity bases of Sequences 2, 3, and 4 are heavily populated by RGB stars \\citep{2002MNRAS.337L..31I,  2003MNRAS.343L..79K, 2004MNRAS.347L..83K, 2004MNRAS.347..720I}. Above the tip of the RGB, models of AGB stars that include the effects of mass loss confirm that stars continue their evolution to higher luminosities \\citep{1993ApJ...413..641V}. In period-luminosity space, the Mira and SRa stars are not clearly separated, with SRa stars found throughout Sequences 1--4.  \\cite{2004AcA....54..129S} identified a more useful division for LPVs in the LMC and Small Magellanic Cloud (SMC) than the Mira/SRa/SRb system. In this system an LPV is classified as an OSARG (Ogle Small Amplitude Red Giant) if one of its three strongest periods falls onto Sequence 4, and as an SRV/Mira if not. This division separates LPVs into two groups with a variety of distinct properties, and it also shows that Sequence D is composed of two different populations: a dimmer population that covers a relatively broad period range, and a more luminous, redder population that shows a tighter period-luminosity relationship (the color changes described here can be seen in \\cite{fraser05}).  Stars in Sequence E show ``ellipsoidal'' light curves \\citep{2004AcA....54..347S}, where the brightness modulation is due to the gravitational distortion of one member of a close binary system. These light curves exhibit dual minima of unequal depths, but the effect is small enough that most methods (including ours) find a period for these systems that is half of the orbital period. We refer to this period as the Fourier period and use it to distinguish Sequences E and D in our plots. When stars in Sequence E are plotted on the period-luminosity diagram at their orbital period (as in \\cite{2004AcA....54..347S} and \\cite{2006ApJ...650L..55D}) they smoothly join Sequence D. This, along with a recent analysis of OGLE light curves in \\cite{2007ApJ...660.1486S}, supports the binary companion explanation for Sequence D put forth by \\cite{Wood99}, as does a radial velocity study of several of these stars \\citep{2006MmSAI..77..537A}.  In \\cite{fraser05} (hereafter Paper I) we used the MACHO and 2MASS magnitudes and colors to characterize LMC stars in each period-luminosity sequence, and classified the stars in each sequence as Miras, SRa, or SRb. At that time, only 52 percent of the stars in the color- and magnitude-defined sample had a well-determined period in the MACHO Variable Star Catalog. In this work, we expand the successfully analyzed stars to 93 percent of our sample, as well as consider their multi-periodic properties. We also use our results to describe the characteristic variability at each of the stages of RGB and AGB evolution by comparison with population synthesis models, including the stars that show very weak or non-existent periodicity. ",
        "conclusions": "\\label{discussion}  The analysis of the frequency spectra of variable stars has been used with great success for characterizing light curve morphology, identifying binary stars, and constraining observed pulsation modes in studies of Cepheids \\citep{1981ApJ...248..291S}, Beat Cepheids \\citep{1995AJ....109.1653A}, Type II Cepheids and RV Tauri stars \\citep{1998AJ....115.1921A}, RR Lyrae \\citep{2000ApJ...542..257A}, and LPVs (see \\S \\ref{intro}). Here we discuss how these techniques contribute to the understanding of the origin of Sequence D and the relationship between the different LPV stages and stellar evolution.  \\subsection{Sequence D}  One-third of the stars in our sample exhibit the ``long secondary period'' phenomenon. Although the exact mechanism for Sequence D is still unknown, there are many reasons to think that it is correlated with binary systems. In this paper, the evidence includes the similarity between the luminosity functions of Sequence D and our entire sample (Figure \\ref{LumFunctions}, right panel), and the similar amplitudes of Sequence D's average infrared light curves across the 2MASS bands. Other evidence includes the smooth connection between Sequence D and Sequence E---when E is plotted at its orbital period \\citep{2004AcA....54..347S}, the presence of ellipsoidal light curves \\citep{2007ApJ...660.1486S}, and radial velocity studies of these stars \\citep{2006MmSAI..77..537A}. However, there are also several significant features of Sequence D that are unique, or show differences from Sequence E.  Sequence D does not suffer from period halving like Sequence E does, and Sequence E stars often have the third and fourth harmonics present in their frequency spectra, while Sequence D stars show the second harmonic instead (\\S \\ref{multiP}). \\cite{2006ApJ...650L..55D} found that in an amplitude-luminosity plot, stars in Sequence D follow a different pattern than in Sequence E. We see that Sequence E appears to be a continuation of only the small-amplitude ($<0.2$ Blue peak-to-peak) group of stars in Sequence D, those stars that roughly correspond to the OSARGs of \\cite{2004AcA....54..129S}. This is consistent with the comparison of the average light curve amplitudes in Blue and $K_s$, where only the lower amplitude Sequence D stars were a good match to Sequence E. Finally, Sequence D's infrared light curves \\emph{lead} the optical light curves by 10--15 percent, which is a feature unique to this sequence. These facts do not necessarily preclude a binary star mechanism for Sequence D, but they are useful constraints for proposed mechanisms. We note, as an additional constraint, that stars associated with Sequence 1 are far less likely to have a period lying on Sequence D, and that Sequence D is not observed in the LMC below  $K_s \\approx 13.7$.  The model proposed by \\cite{2007ApJ...660.1486S}, based on the original model proposed by \\cite{Wood99}, is that of a binary system in which the mass lost from the red giant is concentrated near the companion, and regularly obscures the red giant. The wide range of observed light curve amplitudes for Sequence D stars, from 0.1 up to five magnitudes in the MACHO Blue filter \\citep{fraser05},  can be readily explained by the projection effects of different inclination angles in a binary system.  If Sequences E and D are truly composed of binary systems, then the population of binary stars in our sample includes stars with either their primary or secondary Fourier period lying within the boundaries of these sequences (e.g. the top panel of Figure \\ref{LightCurves} shows a star whose secondary Fourier period lies on Sequence D). After removing periods identified with the One-Year Artifact, we find that 48 percent of the stars in our sample have variability associated with Sequences E or D (see Table \\ref{sequences}). For comparison, \\cite{1991A&A...248..485D} found a binary fraction of approximately 40 percent for nearby solar-type stars, and \\cite{1997AJ....113.2246R} found a fraction of approximately 35 percent for low-mass stars, a trend in mass which is discussed in \\cite{2006ApJ...640L..63L}. It is reasonable to assume that we cannot see pole-on binaries, and there is no reason to think that all binaries have periods shorter than four years, both of which imply that 48 percent underestimates the total percentage of binaries seen in the LMC. This is a serious problem for any explanation of Sequence D that relies solely on binary systems. However, it is very likely that a subset of these stars do show variability due to binarity, perhaps stars in one of the populations that can be separated by color or amplitude.   \\subsection{Comparison to Evolutionary Models}  \\begin{figure} \\begin{center} \\includegraphics[width=0.9\\textwidth]{JK} \\caption[Average $J-K_s$ color for each sequence.]{Average $J-K_s$ color for each sequence, as well as stars in the One-Year Artifact and stars in the ``background'' of the Fourier period-luminosity diagram. $J-K_s$ colors are calculated for 0.5 magnitudes bins and shown for bins with more than 40 stars.} \\label{JK} \\end{center} \\end{figure}  We can begin to characterize the evolution of LPVs by comparison in color-magnitude space to models. \\cite{2003A&A...403..225M} produced a population synthesis model of the $K_s$ vs. $J-K_s$ color-magnitude diagram (CMD) of the LMC using the RGB and Early AGB (E-AGB) evolutionary models of \\cite{2000A&AS..141..371G}, and a preliminary version of their thermally pulsing AGB star models (published later in \\cite{2007A&A...469..239M}). Their Figure 12 illustrates the luminosities and colors corresponding to major phases in an LMC star's evolution from RGB, through the Early AGB, and finally to the thermally pulsing AGB (including a transformation to carbon-dominated atmospheres for some stars). For comparison, our Figure \\ref{JK} shows the average $J-K_s$ color binned in magnitude for each of the Sequences 1--4, E, D, as well as the One-Year Artifact and the background population of stars. The color of the ``background'' stars match Galactic disk turn-off stars and LMC intermediate-mass stars on the Early AGB in the synthetic CMD of \\cite{2003A&A...403..225M}.  As shown in Figure 12, the $J-K_s$ color is more effective at distinguishing evolutionary stages in the AGB than in the RGB. However, with the additional information from the luminosity functions, we can also estimate the importance of RGB stars to each sequence. A distinct peak in the luminosity function at the tip of the RGB  ($K_s=12.3$, \\cite{2000ApJ...542..804N}) is widely taken to indicate that the majority of the stars dimmer than this point are themselves on the RGB \\citep{2002MNRAS.337L..31I,  2003MNRAS.343L..79K, 2004MNRAS.347L..83K, 2004MNRAS.347..720I, fraser05}.  The stars in our sample that show very weak or nonexistent periodicity (the 24 percent identified with the One-Year Artifact) are predominantly RGB stars. The left two panels of Figure \\ref{LumFunctions} show the luminosity functions of stars in the One-Year Artifact with Sequences 3 and 4. Below the tip of the RGB the population of One-Year Artifact stars dominates, suggesting that very weak or non-periodic variability is common among RGB stars. Above the tip of the RGB there is a much closer correspondence between the One-Year Artifact and Sequences 3 and 4. Thus it appears that stars at the dimmest luminosities in our sample vary aperiodically while on the RGB, but most begin to show periodic behavior when they brighten to $K_s \\approx 13$.  After passing off of the tip of the RGB, stars may pulsate with shorter periods and lower luminosity as RR Lyrae on the horizontal branch. They become LPVs again as they ascend the Early AGB (i.e. prior to the first thermal pulse or helium shell flash). \\cite{2004AcA....54..129S} used the slight offset in period between RGB and AGB stars to show that AGB stars pulsate alongside RGB stars below the tip of the RGB.  Evolution proceeds to brighter luminosities until the onset of thermal pulses, which begin at $K_s \\approx 12$ on the synthetic CMD from \\cite{2003A&A...403..225M}. Stars primarily populate Sequences 1 and 2 above this luminosity. \\cite{1996A&A...311..509W} predict that thermal pulses will create large modulations in luminosity and the pulsation period (and mode) in AGB stars with timescales of thousands of years. The models of \\cite{2007A&A...469..239M} substantially agree with these predictions, and show large changes in pulsation period due to mode switching as a direct result of a thermal pulse. Period changes in LPVs are well known \\citep{2005AJ....130..776T} but only a small percentage of stars at any one time should be undergoing a thermal pulse due to the short timescales of thermal pulses relative to the long inter-pulse period. The observed period changes of LPVs are not well explained by the effects of thermal pulses alone.  The middle two panels of Figure \\ref{LumFunctions} compare the luminosity functions of the numbered sequences. The relative importance of the two giant branches shifts from the RGB to the AGB as we move to Sequences 1 and 2. Additionally, the peak number of stars above the tip of the RGB in each sequence is found at higher luminosity from Sequence 4 to Sequence 1. The OSARG versus SRV/Mira distinction of \\cite{2004AcA....54..129S}, by virtue of its definition, roughly corresponds to a division between two pairs of sequences: Sequences 3 and 4, and Sequences 1 and 2. This division is also clearly seen in the observed period ratios (Figure \\ref{Petersen2}). Considering the synthetic CMD from \\cite{2003A&A...403..225M}, OSARGS are closely related to RGB and E-AGB stars, while the SRV/Mira stars are more closely related to thermally pulsing AGB stars.  At $K_s \\approx 11$ the synthetic CMD predicts the formation of carbon stars, and the typical $J-K_s$ colors of each sequence diverge (Figure \\ref{JK}). Only Sequences 1, 2, and D redden to the expected $J-K_s$ color of the carbon star tail of \\cite{2003A&A...403..225M}. Their models also show that these stars do not evolve in brightness after this stage, so the observed range of carbon star luminosities in our sample may be interpreted as a range of stellar masses. \\cite{2007A&A...469..239M} show that only stars between 1.2 and 2.5 $M_\\odot$ undergo a dredge-up that can bring the results of nuclear burning to the atmosphere without quickly destroying it through hot bottom burning. Since Miras exist on Sequence 1 at luminosities both above and below $K_s=11$ \\citep{fraser05}, not all large amplitude pulsators are carbon stars. Also, not all carbon stars have such red $J-K_s$ colors: \\cite{2004A&A...425..595G} found carbon stars on the shorter period sequences---the popular color-cut of $J-K_s \\ge 1.4$ appears more effective at segregating M stars, which are rarely this red.  Using $J-K_s>1.4$ to select areas that include only carbon stars, we see that many occupy the highest luminosities of both Sequences 1 and 2, as also seen in \\cite{2007arXiv0710.0953L}. Approximately 40 percent of the carbon stars lie on Sequence 2, similar to the prediction of \\cite{2003A&A...403..225M}, who fit the observed $K_s$ vs. $J-K_s$ CMD by assuming a 50 percent mix of fundamental and first-overtone pulsation among the carbon stars in their model. Figure \\ref{JK} shows that the stars on Sequence 1 evolve to redder $J-K_s$ colors than on Sequence 2, presumably due to increased mass loss driven by fundamental-mode pulsation \\citep{2007A&A...469..239M}. Some stars on the long-period extreme of Sequence 1 are underluminous for their color, which may be self-extinction due to dusty outflows.  Beyond this point, stars begin their rapid post-AGB evolution, and they quickly move out of the color-magnitude space of our sample.  \\subsection{Brief Summary of LPV Evolution}  LMC stars on the RGB and AGB are characterized by the presence of multiple long periods that show increasing length and amplitude as these stars evolve. Apart from the presence of the long secondary period phenomenon, which appears in approximately half of the stars in our sample, stars initially vary non-periodically, and only later begin to pulsate with periods of 20--120 days on Sequences 3 and 4. The Petersen diagram and its variants show that stars with their primary period on either one of these sequences often have their secondary period on the other sequence. Furthermore, the period changes between the first half and last half of the MACHO light curves suggest that the amplitudes of these pairs of periods are very similar for stars on the inside edges of these two sequences. Similar results are obtained for stars which have evolved to the luminosity at which thermal pulses begin on the AGB ($K_s \\approx 12$); these stars are usually found on Sequences 1 and 2 with periods of 45--500 days. This supports the arguments of  \\cite{1986MNRAS.219..525W} and \\cite{agbstars2}, that a star excites different pulsation modes in turn as it evolves. At the points where the dominant modes switch, we observe pulsation in multiple periods with near equal strengths. The increase in the luminosity of the maximum of the luminosity function above the tip of the RGB also lends support to this argument. The highest luminosity stars on Sequences 1 and 2 have become carbon stars. After this point, the mass-loss rate of these stars increases drastically and they rapidly evolve out of our sample of luminous red stars."
    },
    "0808/0808.0507_arXiv.txt": {
        "abstract": "We present the results of a recent reverberation-mapping campaign undertaken to improve measurements of the radius of the broad line region and the central black hole mass of the quasar PG\\,2130+099. Cross correlation of the 5100\\,\\AA\\ continuum and H$\\beta$ emission-line light curves yields a time lag of 22.9$^{+4.4}_{-4.3}$ days, corresponding to a central black hole mass $M_{\\rm BH} = $ ($3.8 \\pm 1.5 $)$ \\times 10^{7} M_{\\odot}$. This value supports the notion that previous measurements yielded an incorrect lag. We re-analyzed previous datasets to investigate the possible sources of the discrepancy and conclude that previous measurement errors were apparently caused by a combination of undersampling of the light curves and long-term secular changes in the H$\\beta$ emission-line equivalent width. With our new measurements, PG\\,2130+099 is no longer an outlier in either the $R_{\\rm BLR}$--$L$ or the $M_{\\rm BH}$--$\\sigma_*$ relationships. ",
        "introduction": "Reverberation mapping uses observations of continuum and emission-line variability to probe the structure of the broad line region (BLR) in active galactic nuclei (Blandford \\& McKee 1982; Peterson 1993). It has been extensively used to estimate the physical size of the BLR and the mass of central black holes in active galactic nuclei (AGN). The observed continuum variability precedes the observed emission-line variability by a time related to the light travel time across the BLR; by obtaining an estimate of the time delay, or ``lag'' $\\tau$, between the change in continuum flux and the change in emission-line flux, one can estimate the size of the BLR. While this is an extremely effective method, high-quality datasets are difficult to obtain, as they require well-spaced observations over long timescales. To date, over three dozen AGN have estimated black hole masses obtained using reverberation methods.  In addition to being important physical parameters, the BLR radius and central black hole mass ($M_{\\rm BH}$) measurements are crucial in calibrating relationships between different properties of AGN. A useful relationship that has emerged is the correlation between the radius of the BLR ($R_{\\rm BLR}$) and the optical luminosity of the AGN (e.g. Kaspi et al.\\ 2000, 2005; Bentz et al.\\ 2006). Another key relationship is that between $M_{\\rm BH}$ and bulge stellar velocity dispersion ($\\sigma_*$) which is seen in both quiescent (Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000a; Tremaine et al.\\ 2002) and active galaxies (Gebhardt et al.\\ 2000b; Ferraresse et al.\\ 2001; Nelson et al.\\ 2004; Onken et al.\\ 2004; Dasyra et al.\\ 2007). These two relationships are critical because they allow us to estimate the masses of black holes in AGN from a single spectrum (see McGill et al.\\ 2008 for a recent summary) --- we estimate $R_{\\rm BLR}$ from the AGN luminosity, the velocity dispersion of the BLR is determined by the emission-line width, and the quiescient $M_{\\rm BH}$--$\\sigma_{*}$ relationship provides the calibration of the reverberation-based mass scale. Obtaining high-quality single-epoch spectra is much less observationally demanding than obtaining high-quality reverberation measurements; once relations such as these are properly calibrated, we can measure $M_{\\rm  BH}$ in many more objects than would otherwise be possible. An extensive set of objects with reliable black hole masses allows us to explore the connection between black holes and AGN evolution over cosmologically interesting timescales.  PG\\,2130+099 has been a source of curiosity because it is an outlier in both the $M_{\\rm BH}$--$\\sigma_{*}$ and $R_{\\rm BLR}$--$L$ relations. Previous measurements obtained by Kaspi et al.\\ (2000) and reanalyzed by Peterson et al.\\ (2004) found H$\\beta$ lags of about 180 days. These measurements yield a black hole mass $M_{\\rm BH}$ upwards of $10^{8}$ $M_{\\odot}$. At luminosity $\\lambda L_{\\lambda}$(5100\\,\\AA)=($2.24\\pm0.27$)$\\times10^{44}$ erg s$^{-1}$, this places PG\\,2130+099 well above the $R_{\\rm BLR}$--$L$ relationship (Bentz et al.\\ 2006). Dasyra et al.\\ (2007) also note that PG\\,2130+099 falls above the $M_{\\rm BH}$--$\\sigma_{*}$ relationship. Together these suggest that this discrepancy could be caused by measurement errors in the BLR radius. Our suspicions are also fueled by two other factors. First, the optical spectrum of PG\\,2130+099 is similar to that of narrow-line Seyfert 1 (NLS1) galaxies, which are widely supposed to be AGN with high accretion rates relative to the Eddington rate (see Komossa 2008 for a recent review). However, the accretion rate derived from this mass and luminosity (under common assumptions for all reverberation-mapped AGN as described by Collin et al.\\ 2006) is quite low compared to NLS1s, again suggesting that $M_{\\rm BH}$ and therefore $\\tau$ are overestimated. Second, as we discuss further below, a lag of approximately half a year on an equatorial source means that fine-scale structure in the two light curves will not match up in detail, and cross-correlation becomes very sensitive to long-term secular variations that may or may not be an actual reverberation signal.  For these reasons, we decided to undertake a new reverberation campaign to remeasure the H$\\beta$ lag for PG\\,2130+099. In this paper, we present a new lag determination from this campaign. We also present a re-analysis of the earlier dataset that suggests that the true lag is consistent with our new lag value and investigate possible sources of error in the previous analysis. ",
        "conclusions": "Analysis of previous datasets of PG\\,2130+099 yields lags greater than 150 days and $M_{\\rm BH}$ values above $10^{8}$ $M_{\\odot}$, in both cases approximately an order of magnitude larger than our new value. To investigate the source of these discrepancies, we closely scrutinized the previous dataset, which consists of data obtained at Steward Observatory and Wise Observatory (Kaspi et al.\\ 2000). The 5100\\,\\AA\\ continuum and emission-line light curves, along with their respective CCF and ZDCFs, are shown in Figure \\ref{fig:f4}. We first ran the full light curve through the time series analysis as described in section 2.3 to confirm the previous results, and successfully reproduced a lag measurement of $\\sim168$ days, in agreement with Kaspi et al.\\ (2000) and Peterson et al.\\ (2004). However, visual inspection of the light curves suggests that these values were in error; we were able to identify several features that are present in the continuum, H$\\beta$ and H$\\alpha$ light curves that are quite visibly lagging on timescales shorter than 50 days, as we show below.  It is clear that the Kaspi et al.\\ results are dominated by the data spanning the years 1993-1995, as this period contains the most significant variability as well as time sampling that is sufficient to resolve features in the light curves. We show the light curves for all spectral features measured by Kaspi et al.\\ during this time period in Figure \\ref{fig:f5}. Over this period, we can match behavior in the continuum with the behavior of the emission lines, most noticeably in the data from 1995. The most prominent features in the light curves are the maximum in the continuum at the end of the 1994 series and the maxima in the lines at the beginning of the 1995 series, which, based on the continuum variations, one would expect to be lower relative to the fluxes in 1994. However, inspection of the relative flux levels of the lines and continuum in the 1994 and 1995 data reveals that the equivalent width of the emission lines changes during this timespan, which results in a rise in emission-line flux apparently unrelated to reverberation. The maximum in the continuum and those in the emission lines do not correspond to the same features--- the maxima in the emission line light curves in 1995 correspond to the similar feature in the continuum at the beginning of 1995. Because there are no observations during the six months or so between these features, the CCF and ZDCF lock onto the two unrelated maxima that are separated by 196 days and yield lags of $\\sim$200 days over this three-year span. The emission lines likely continued to increase in flux during the time period for which there were no observations.  Inspection of the light curve segments in Figure \\ref{fig:f5} reveals similar structures in each of them within a given year. To quantify the time delays between the continuum and lines on short timescales, we cross-correlated each individual year, this time using only flux randomization in the Monte Carlo realizations, as there were too few points in each year to use random subset sampling, so the uncertainties are underestimated. The resulting lags are given in Table \\ref{Table:lags}. Again we must consider that the CCF cannot produce a lag that is greater than the duration of observations, so the individual CCFs for these three years are limited to short delays. However, the presence of the 1995 feature in both the continuum and emission-line light curves is unmistakable and it is extremely unlikely that it lags by a value exceeding the sensitivity limit of the CCF. From this we surmise that the discrepancies in measured lag values are mostly a result of large time gaps in the data and/or underestimation of the error bars in the Kaspi et al.\\ data.  Welsh (1999), and even earlier, P\\'{e}rez, Robinson, \\& de la Fuente (1992), pointed out that emission-line lags can be severely underestimated with light curves that are too short in duration, particularly if the BLR is extended: certainly our 98-day campaign is insensitive to lags as long as $\\sim180$\\,days. Is it possible that the original lag determination of Kaspi et al.\\ (2000) and Peterson et al.\\ (2004) is correct (or more nearly so) and we have been fooled by reliance on light curves that are too short? We think not, based on (1) the reasonable match between details in the continuum and emission-line light curves for four different observing seasons (1993, 1994, 1995, and 2007), (2) the improved agreement with the $R_{\\rm BLR}$-$L$ relationship with the smaller lag (demonstrated in Figure \\ref{fig:f6}), and (3) the improved agreement with the $M_{\\rm BH}$-$\\sigma_*$ relationship with the smaller lag (Figure \\ref{fig:f7}). It is also worth pointing out the difficulty of accurately measuring a $\\sim 180$\\,day lag, particularly in the case of equatorial sources which have a relatively short observing season. The short observing season, typically 6-7 months, means that there are very few emission-line observations that can be matched directly with continuum points: the observed emission-line fluxes represent a response to continuum variations that occurred when the AGN was too close to the Sun to observe.  Welsh (1999) has also pointed out the value in ``detrending'' the light curves--- removing long-term trends by fitting the light curves with a low-order function can reduce the bias toward underestimating lags. In this particular case, we find that detrending has almost no effect: in particular, the highest points in the continuum (in late 1994) and the highest points in the line (early in the 1995 observing season) remain so after detrending, and both the interpolation CCF and ZDCF weight these points heavily. It is also interesting to note that the peak in the ZDCF agrees with the peak of the interpolation CCF (Figure \\ref{fig:f4}), which demonstrates that the $\\sim180$\\,day lag is {\\em not} simply ascribable to interpolation across the gap between the 1994 and 1995 observing seasons.  The data from our 2007 campaign are not ideal: the amplitude of variability was low, the time sampling was adequate, but only barely, and the duration of the campaign was short enough that we lack sensitivity to lags of $\\sim 50$\\,days or longer. But the preponderance of evidence at this point argues that our smaller lag measurement is more likely to be correct than the previous determination. Certainly better-sampled light curves of longer duration would yield a more definitive result."
    },
    "0808/0808.2391_arXiv.txt": {
        "abstract": "We present an analysis of the radio properties of large samples of Lyman Break Galaxies (LBGs) at $z \\sim 3$, 4, and 5 from the COSMOS field.  The median stacking analysis yields a statistical detection of the $z \\sim 3$ LBGs (U-band drop-outs), with a 1.4 GHz flux density of $0.90 \\pm 0.21 \\mu$Jy. The stacked emission is unresolved, with a size $< 1\"$, or a physical size $< 8$kpc.  The total star formation rate implied by this radio luminosity is $31\\pm 7$ $M_\\odot$ year$^{-1}$, based on the radio-FIR correlation in low redshift star forming galaxies.  The star formation rate derived from a similar analysis of the UV luminosities is 17 $M_\\odot$ year$^{-1}$, without any correction for UV dust attenuation.  The simplest conclusion is that the dust attenuation factor is 1.8 at UV wavelengths. However, this factor is considerably smaller than the standard attenuation factor $\\sim 5$, normally assumed for LBGs. We discuss potential reasons for this discrepancy, including the possibility that the dust attenuation factor at $z \\ge 3$ is smaller than at lower redshifts.  Conversely, the radio luminosity for a given star formation rate may be systematically lower at very high redshift. Two possible causes for a suppressed radio luminosity are: (i) increased inverse Compton cooling of the relativistic electron population due to scattering off the increasing CMB at high redshift, or (ii) cosmic ray diffusion from systematically smaller galaxies.  The radio detections of individual sources are consistent with a radio-loud AGN fraction of 0.3\\%. One source is identified as a very dusty, extreme starburst galaxy (a 'submm galaxy'). ",
        "introduction": "The power of discovering high redshift galaxies via the broad-band drop-out technique, ie. Lyman Break galaxies (LBGs), is now well established (Steidel et al. 1999; 2000). Using this technique, over a thousand galaxies have now been detected at $z > 2$.  Detailed studies show that these galaxies have stellar masses between 10$^{10}$ and 10$^{11}$ $M_\\odot$, and a (comoving) volume density at $z \\sim 3$ of $\\sim 0.005$ Mpc$^{-3}$ (Giavalisco 2002).  A number of issues are still being investigated for LBGs. One important issue is deriving the total star formation rate. Besides model parameters such as the IMF and star formation history, the dust correction in the UV remains under investigation.  Steidel et al. (1999) originally estimated a typical UV dust attenuation factor of about 5, based on optical spectroscopy.  Adelberger \\& Steidel (2000) applied a similar method to a larger sample of LBGs, as well as observations at other wavebands (radio, submm), and conclude: \"...the mean extinction at 1600 \\AA ~ for LBGs, a factor of 6 in our best estimate, could lie between a factor 5 and a factor of 9.\"  More recently, Reddy \\& Steidel (2004) have derived the average dust correction factors for UV-selected galaxies at $z \\sim 2$ based on deep X-ray, radio, and optical spectroscopic studies of galaxies in the GOODS north field. They find that, for galaxies with total star formation rates $> 20$ $M_\\odot$ year$^{-1}$, the dust attenuation factor is between 4.4 and 5.1.  A second issue for LBGs is the AGN fraction. Shapley et al. (2003) found that $\\sim 3\\%$ of LBGs at z$\\sim$3 have optical emission line spectra that are consistent with AGN.  The Cosmic Evolution Survey (COSMOS), covering 2 \\sq\\deg, is a comprehensive study of the evolution of galaxies, AGN, and dark matter as a function of their cosmic environment. The COSMOS field has state-of-the-art multiwavelength observations, ranging from the radio through the X-ray (Scoville et al. 2007; Capak et al. 2007). A major part of this study is to identify the largest samples of LBGs to date. In section 2 we describe the LBG samples from the COSMOS survey based on U, B, and V-band drop-out searches.  In this paper we consider the radio properties of these large samples of LBGs, similar to the study of $z = 5.7$ Lyman-$\\alpha$ emitting galaxies by Carilli et al. (2007). Observations of the COSMOS field have been done at 1.5$\"$ resolution (FWHM) at 1.4 GHz with the Very Large Array (Schinnerer et al.  2007). We employ the new VLA image that has added integration time in the central $\\sim 1^o$ (Schinnerer in prep.), giving an rms at the field center of $\\sim 7$ $\\mu$Jy beam$^{-1}$.  In section 3.1 we search for radio counterparts to individual LBGs, while in section 3.2 we perform a stacking analysis to derive the statistical properties of the samples.  In section 4 we discuss the implications these observations have on the questions of dust extinction and the AGN fraction in LBGs. We adopt a standard concordance cosmology. ",
        "conclusions": "\\subsection{The AGN fraction and massive starbursts}  At the sensitivity limits of current deep radio surveys, such as COSMOS, detection of individual sources at $z > 2$ is limited to either extreme starbursts (star formation rates $> 1000$ $M_\\odot$ year$^{-1}$), or radio jet sources with luminosities comparable to M87.  After correcting for random detections, we obtain a detection rate of $\\sim 0.3\\%$.  Shapley et al. (2003) find an AGN fraction in LBG samples of $\\sim 3\\%$, based on optical spectroscopy.  If all of our detected sources are radio AGN, then the implied radio loud fraction is $\\sim 10\\%$. Interestingly, a value of 10\\% is the canonical value for radio loud AGN based on studies of nearby galaxies (eg. Ivezic et al. 2002; Petric et al. 2007).  We emphasize caution in interpreting our results on the radio loud fraction of AGN at high redshift, for a number reasons. First, in a recent comparison of the FIRST and SDSS surveys, Jiang et al. (2007) find that the radio loud fraction likely depends on both optical luminosity and redshift. This study is not directly comparable to our COSMOS study, since they were limited to much more luminous sources at high redshift, by about two orders of magnitude in the radio.  Second, the above analysis does not consider the possibility that some of the 0.3\\% of the radio detections are extreme, highly obscured starbursts, comparable to the submm galaxies. Indeed, in a companion paper, we present the discovery of a submm galaxy in our radio-selected LBG COSMOS sample (Capak et al. 2008). The radio-detected LBG has a spectroscopic redshift of $z =4.5$, and a radio flux density of $45\\pm 10\\mu$Jy. Thermal emission from warm dust is detected from this galaxy at 250GHz using the  MAMBO bolometer camera, with a flux density of $3.4\\pm 0.7$mJy.  This source is the first submm galaxy yet identified (spectroscopically) at $z > 4$ (Capak et al. 2008).  And third are the caveats mentioned in section 2 concerning our conservative source selection criteria, eg. excluding obvious optically bright broad line AGN.  \\subsection{The sub-$\\mu$Jy radio source population}  The unprecedented size of the COSMOS LBG sample has allowed us to reach sub-$\\mu$Jy sensitivity levels in the stacking analysis of the LBG samples. We obtain a clear statistical detection of $S_{\\rm 1.4GHz} = 0.90\\pm 0.21\\mu$Jy for the U-band ($z \\sim 3$) drop-outs. The implied rest frame 1.4 GHz luminosity density is $L_{\\rm 1.4\\rm GHz} = 5.1\\times 10^{29}$ erg s$^{-1}$ Hz$^{-1}$, assuming a spectral index of $-0.8$ (Condon 1992).  Assuming the radio emission is driven by star formation, we can derive a total star formation rate (0.1 to 100 $M_\\odot$), from the rest frame 1.4 GHz luminosity density. There are a number of recent calibrations of this relationship for nearby galaxies using, eg.  the IRAS and NVSS surveys.  These studies adopt the star formation rates derived from the far-IR emission based on the relations in Kennicutt (1998), assuming a Salpeter IMF, and then calibrate the radio conversion factor using the tight radio to far-IR correlation for star forming galaxies.  We adopt the conversion factor from the study of Yun, Reddy, and Condon (2001):  $$\\rm SFR = 5.9 \\pm 1.8 \\times 10^{-29} {\\sl L}_{\\rm 1.4GHz} ~ M_\\odot~ year^{-1}, $$  \\noindent where $L_{\\rm 1.4\\rm GHz}$ is in erg s$^{-1}$ Hz$^{-1}$. From this, we calculate a mean star formation rate of $31\\pm 7$ $M_\\odot$ year$^{-1}$. Note that the conversion factor above is within 10\\% of that derived by Bell (2003).  The mean observed UV luminosity at a rest frame wavelength of 1600 \\AA ~of the U-band drop-out sample is: $L_{2000A} = 1.2\\times 10^{29}$ erg s$^{-1}$ Hz$^{-1}$ (Capak et al. 2008 in prep).  Using equation (1) in Kennicutt (1998) for the relationship between UV luminosity and total star formation rate, we derive a total star formation rate of 17 $M_\\odot$ year$^{-1}$, uncorrected for dust attenuation.  The ratio of radio derived star formation rate to UV derived star formation rate is $1.8\\pm 0.4$.  The simplest conclusion is that the UV emission is attenuated by dust by a factor of 1.8. However, this factor is considerably smaller than the standard factor $\\sim 5$ adopted for LBGs (see Section 1). We consider some possible reasons for this difference.  One possibility is that the dust attenuation factor for the $z \\sim 3$ COSMOS LBG sample is indeed smaller than for other LBG samples. Most of the studies of dust attenuation of LBG galaxies have been at $z < 3$ (see section 1), and hence it is possible that at higher redshift the dust attenuation factor decreases. An argument against this decrease is the recent X-ray and optical study of a sample of LBGs at $z \\sim 3$ by Nandra et al. (2002), who also derive a dust attenuation factor of about 5. We also re-emphasize that the COSMOS LBG sample analyzed herein de-selected sources with poor phot-z fits, which can occur for very heavily obscured objects (Section 2).  On the other hand, Wilkins, Trentham, \\& Hopkins (2008) present a detailed comparison of the build up of stellar mass in galaxies versus the cosmic star formation rate density. They conclude that there is a discrepancy between the rate of stellar mass creation, and the star formation rates derived from the UV luminosities, at high redshift, if one assumes the standard factor $\\sim 5$ dust attenuation. The discrepancy is in the sense that the UV derived star formation rates are too high. They find that the peak in the discrepancy occurs at $z \\sim 3$, and they suggest that the UV-derived star formation rates at this redshift may be over-estimated by a factor $\\sim 4$. They also point out that '...this large deviation at high redshift offers an explanation for why the integrated star formation history implies a local stellar mass density in excess of that measured.'  A second possibility is that the conversion factor of radio luminosity to star formation rate is higher in the $z \\sim 3$ LBG COSMOS sample than has been derived for low redshift galaxies.  Radio studies of 24$\\mu$m selected galaxies by Appleton et al. (2004) imply that the radio conversion factor is constant out to $z \\sim 1$, ie. that the radio-FIR correlation is constant out to this redshift.  Reddy \\& Steidel (2004) extend this conclusion out to $z \\sim 2$ in their extensive study of UV selected galaxies. Most recently, Ibar et al. (2008) conclude that the radio conversion factor remains constant out to $z \\sim 3$, using a 24$\\mu$m source sample with redshifts from the SXDF. However, given the need for individual source detections in the radio, their $z \\sim 3$ radio sources have star formation rates about two orders of magnitude larger than the stacking results presented herein, and hence a direct comparison is problematic.  One physical reason why we might expect the radio conversion factor to diverge at the highest redshifts is increased relativistic electron cooling due to inverse Compton scattering off the cosmic microwave background (CMB). The ratio of relativistic electron energy losses due to synchrotron radiation, to energy losses due to inverse Compton radiation, equals the ratio of the energy density in the magnetic field to that in the photon field.  The energy density in the CMB increases as: $U_{CMB} = 4.2\\times 10^{-13} (1+z)^4$ ergs cm$^{-3}$. Figure 2 shows a comparison of $U_{CMB}$ with the typical energy densities in the magnetic fields in different regions in galaxies. We show the range of fields considered typical for spiral arms ($\\sim$ few $\\mu$G), and for starburst galaxy nuclei (of order 100$\\mu$G; Beck et al. 1994; see review by Beck 2005).  The important point is that IC losses off the CMB will dominate synchrotron losses in a typical ISM at $z \\ge 0.5$, and dominate in starburst nuclei at $z \\ge 4$.  We note that inverse Compton losses will not affect the thermal electrons responsible for Free-Free emission from star forming galaxies. Such Free-Free emission may dominate the total radio emission from galaxies at rest frequencies between roughly 40 GHz and 100 GHz (Condon 1992), and hence become an important factor in radio continuum studies of very high redshift galaxies.  A depressed radio luminosity for a given star formation rate could also arise if the $z \\sim 3$ LBGs are systematically smaller galaxies. The standard model that produces the radio-FIR correlation (Condon 1992) requires a cosmic ray processing box size $\\ge 1$ kpc. It has been observed that dwarf galaxies at low redshift depart from the radio-FIR correlation by about a factor of 2, in the sense of being radio under-luminous.  The hypothesis is that the cosmic rays diffuse out of the galaxy on timescales shorter than required to maintain the standard radio-FIR correlation (Yun et al. 2001). Giavalisco (2002) describes high $z$ LBGs as having typical half-light radii of 4 to 7 kpc, significantly larger than typical dwarf galaxies. However, he points out that: \"..frequently the galaxies have disturbed or fragmented morphologies, with one bright core, or multiple knots embedded in diffuse nebulosity, reminiscent of merger events.\"  In this paper, we report the first robust statistical (median) detection of sub-$\\mu$Jy radio emission from LBG galaxies at $z \\sim 3$. This detection was made possible by the very wide area, and depth, of the Cosmos field. While our physical interpretation of the result remains inconclusive, there are a number of future studies we are pursuing to address the interesting questions raised concerning the UV attenuation factor for LBGs, and the the radio luminosity to star formation rate conversion factor, at $z \\ge 3$. The most important study involves obtaining optical spectra of a large sample of LBGs from the COSMOS sample. Spectra will elucidate the nature of the sources, and allow for a study of the dust correction factor as a function of galaxy type (AGN, starburst, elliptical...). Also, the selection criteria for this LBG sample were very conservative (Section 2).  We will further refine (and increase) our LBG samples using the Cosmos multiband photometry, and perform statistical analyses of the radio properties as a function of eg. stellar mass, or $A_V$. We are also exploring the IR properties of these samples with Spitzer. Such IR studies have particular relevance in the light of the results of Ivison et al. (2007), who find that for the seven IR luminous $z \\sim 3$ LBGs in their AEGIS sample ($S_{24\\mu m} > 60\\mu$Jy), the median 1.4 GHz flux density is 44$\\mu$Jy, comparable to that expected for sub-mm galaxies. Lastly, unambiguous identification of radio AGN using eg. optical spectra, will provide input into faint radio source population models that are being generated in the context of planning for future, large area radio telescopes (Wilman et al. 2008). The results presented herein indicate that future deep (sub-$\\mu$Jy), wide field surveys with the Expanded Very Large Array, will detect routinely the radio emission from individual LBGs out to $z \\sim 3$."
    },
    "0808/0808.1939_arXiv.txt": {
        "abstract": "I decompose the unstable growing modes of stellar disks to their Fourier components and present the physical mechanism of instabilities in the context of resonances. When the equilibrium distribution function is a non-uniform function of the orbital angular momentum, the capture of stars into the corotation resonance imbalances the disk angular momentum and triggers growing bar and spiral modes. The stellar disk can then recover its angular momentum balance through the response of non-resonant stars. I carry out a complete analysis of orbital structure corresponding to each Fourier component in the radial angle, and present a mathematical condition for the occurrence of van Kampen modes, which constitute a continuous family. I discuss on the discreteness and allowable pattern speeds of unstable modes and argue that the mode growth is saturated due to the resonance overlapping mechanism. An individually growing mode can also be suppressed if the corotation and inner Lindblad resonances coexist and compete to capture a group of stars. Based on this mechanism, I show that self-consistent scale-free disks with a sufficient distribution of non-circular orbits should be stable under perturbations of angular wavenumber $m>1$. I also derive a criterion for the stability of stellar disks against non-axisymmetric excitations. ",
        "introduction": "\\label{sec:intro}  Both $N$-body simulations \\citep{H71} and analytical methods (Kalnajs 1978; Jalali \\& Hunter 2005, hereafter JH) show that global instabilities can generate barred structures in galactic disks. Apart from the bar mode, which is an isolated event in frequency space, global spiral modes seen in the eigenspectra of cored stellar disks (Jalali 2007, hereafter Paper I) constitute a discrete family that bifurcates form stationary \\citet{vK55} modes. But not all spiral structures are global modes as disturbances induced by close neighbors \\citep{BH92} and density inhomogeneities \\citep{T90} may also create the spiral patterns of the observed galaxies.  A mode of a stellar disk is a mathematical entity that comes out of an eigenvalue problem. However, its physical origin in isolated systems has not yet been understood clearly. Lynden-Bell \\& Kalnajs (1972, hereafter LBK) attempted to explain a mode through the transport of angular momentum between different parts of the disk. They suggested that the inner Lindblad resonance (ILR) releases the angular momentum of central regions and the spiral structure transports it to the outer parts through the corotation (CR) and outer Lindblad resonances (OLR). This mechanism is favored by some galactic dynamicists \\citep{ATHA03}, but it is seriously challenged by Toomre's (1981) theory that says that feedback through the galactic center is a critical ingredient for growing modes. In Toomre's theory, on the other hand, the modeling of feedback as the reflection of a leading spiral wave at the galactic center and its emergence as a trailing one, is a simple description of a very complicated dynamics that governs the motions of stars. Unresolved issues concerning the evolution of unstable modes include the following: (i) The swing amplification theory is not capable of predicting the fate of a growing mode against other stationary and unstable modes that coexist in the eigenspectrum of a given model. (ii) How does the bar mode saturate? \\citep{KJKJ07}. (iii) Why does the nonlinear bar terminate almost at the corotation radius? \\citep{S81} (iv) We should also understand the origin of different species in an eigenspectrum and interpret their continuous or discrete nature, and the distribution of their pattern speeds and growth rates.  Stellar orbits in a galactic disk begin to evolve once the surface density deviates from its equilibrium state, $\\Sigma_0(\\textbf{\\textit{x}})$, and develops a time-varying mean-field potential $V_1(\\textbf{\\textit{x}},t)$. Here $\\textbf{\\textit{x}}$ denotes the position vector of stars at the time $t$. In the linear regime, we are usually interested in density waves that grow/decay according to the exponential law $e^{st}$ and rotate with the fixed pattern speed $\\Omega_p$. For a wave of $m$-fold symmetry, the perturbed potential becomes \\begin{equation} V_1=\\epsilon e^{st} \\tilde V \\left (\\textbf{\\textit{x}},m\\Omega_p t \\right ). \\label{eq:perturbed-potnetial-introduction} \\end{equation} Since we are dealing with infinitesimal perturbations, I have introduced the small parameter $\\epsilon$ so that $\\epsilon e^{st}\\ll 1$. From (\\ref{eq:perturbed-potnetial-introduction}) one arrives at the equations of motion \\begin{equation} \\dot \\textbf{\\textit{x}} = \\textbf{\\textit{v}},~~ \\dot \\textbf{\\textit{v}} = -\\frac{\\partial V_0}{\\partial \\textbf{\\textit{x}} } - \\epsilon e^{st} \\frac{\\partial \\tilde V}{\\partial \\textbf{\\textit{x}} }, \\label{eq:equation-motion-for-x-and-v} \\end{equation} where $V_0(\\textbf{\\textit{x}})$ is the equilibrium potential field generated by galactic stars and a possible dark matter halo. In writing equation (\\ref{eq:equation-motion-for-x-and-v}), I have assumed that the motion of stars is restricted to the disk plane. When the equilibrium state is axisymmetric and the dark component is spherical, $V_0$ becomes a function of radial distance to the galactic center and the unperturbed equations (with $\\epsilon=0$) are integrable. In such a circumstance, the phase space is filled by rosette orbits denoted by $[\\textbf{\\textit{x}}_0(t),\\textbf{\\textit{v}}_0(t)]$. The growth of perturbations, whatever the magnitude of $\\epsilon e^{st}\\ll 1$ may be, deforms stellar orbits. Orbital deformations are measured by $\\tilde \\textbf{\\textit{x}}=\\textbf{\\textit{x}}-\\textbf{\\textit{x}}_0$ and $\\tilde \\textbf{\\textit{v}}=\\textbf{\\textit{v}}-\\textbf{\\textit{v}}_0$, which can be used in (\\ref{eq:equation-motion-for-x-and-v}) to obtain \\begin{equation} \\frac{d\\tilde \\textbf{\\textit{x}} }{dt}=\\tilde \\textbf{\\textit{v}},~~ \\frac{d\\tilde \\textbf{\\textit{v}} }{dt}=- \\left [ \\frac{\\partial^2 V_0}{\\partial \\textbf{\\textit{x}}^2 } \\right ]_{ \\textbf{\\textit{x}}_0 } \\!\\!\\!\\! \\cdot \\tilde \\textbf{\\textit{x}} - \\epsilon e^{st} \\left [ \\frac{\\partial \\tilde V}{\\partial \\textbf{\\textit{x}} } \\right ]_{ \\textbf{\\textit{x}}_0 }. \\label{eq:equation-motion-for-tilde-x-and-v} \\end{equation}  Although a proper equilibrium distribution function $f_0(\\textbf{\\textit{x}},\\textbf{\\textit{v}})$ can self-consistently reproduce $\\Sigma_0(\\textbf{\\textit{x}})$ using rosette orbits, the perturbed density $\\Sigma_1(\\textbf{\\textit{x}},t)$ (corresponding to $V_1$) cannot be supported by rosette orbits alone and orbital deformations are necessary for the self-consistency of density waves. According to equations (\\ref{eq:equation-motion-for-x-and-v}) and (\\ref{eq:equation-motion-for-tilde-x-and-v}), orbital deformations of ${\\cal O}\\left ( \\epsilon e^{st} \\right )$ are sufficient to support the growth of density/potential perturbations up to the same order of magnitude of such deformations over a time scale of $1/{\\cal O}\\left ( \\epsilon e^{st} \\right )$. As the time is elapsed, the amplitude of perturbations increases exponentially and the solution of the linearized collisionless Boltzmann equation (CBE) fails when $\\epsilon e^{st}\\sim 1$. During my mode calculations, I realized that the orbital axes of certain stars librate in a coordinate frame that rotates with the density pattern. This {\\it resonant capture} initially seemed to be a higher-order nonlinear effect but further experiments showed that the resonant gap is constrained by the magnitude of density perturbations. The complex behavior of stars for infinitesimally small yet non-zero $\\epsilon e^{st}\\ll 1$, and the role of resonant stars in the generation of discrete galactic modes, are investigated in this paper.  I use the results of Paper I and introduce a new dynamical mechanism that sparks unstable modes and governs the singular oscillations of \\citet{vK55} modes. Resolving the origin of instabilities and amplitude saturation precede my nonlinear calculations, which were made feasible in Paper I by the Petrov-Galerkin method and reducing the CBE to a system of nonlinear ordinary differential equations. Those reduced equations, however, are valid only when orbits are regular and averaging over angle variables is allowed. As unstable modes grow, chaotic orbits come into existence and the weighted residual form of the CBE must be modified to handle them. I quote some of the results of such modifications in this paper when I discuss the issue of mode saturation. In Paper III, I will give a full account of the mathematical and numerical modeling of stochastic layers, and will analyze modal interactions after their saturation phase.  For the cored exponential disk embedded in the field of the cored logarithmic potential, I describe the decomposed Fourier components of unstable bar and spiral modes in \\S\\ref{sec:mode-decomposition} and highlight the existence of a phase shift between different components. In \\S\\ref{sec:capture-into-resonances}, I derive a condition for the corotation of the orbital axes of an ensemble of stars and explain the role of such a synchronous motion in pattern formation. I dedicate \\S\\ref{sec:perturbed-stellar-dynamics} to exploring the orbital structure of a perturbed stellar disk and identify a resonance mechanism that can generate both stationary and growing modes. I reveal the mechanism of angular momentum transfer between Fourier components and derive analytical expressions for the growth of resonance zone. I address the origin of instabilities in \\S\\ref{sec:origin-of-instabilities}, present a saturation mechanism for unstable modes in \\S\\ref{sec:mode-saturation}, and discuss about the global stability of soft-centered and scale-free disks. I explain the restrictions of LBK's mode mechanism in \\S\\ref{sec:discussion-and-conclusions} and end up the paper with concluding remarks. ",
        "conclusions": "\\label{sec:discussion-and-conclusions}  I had already given an evidence in \\S\\ref{sec:modes-in-frequency-space} that LBK's theory cannot be extended to large growth rates. Here I provide a more detailed analysis of LBK's mode mechanism. Taking the $s\\rightarrow 0$ limit of (\\ref{eq:dL-dt-from-JH}) leads to equation (30) of LBK. The resulting integrand of the $l$th component involves the term $\\delta(l\\Omega_R+m\\Omega_{\\phi}-\\omega)$ but its argument never becomes zero for $l\\le -1$ as long as $\\Omega_p$ exceeds $\\Omega_{\\rm ILR}(m)$. Once the delta function vanishes for $\\Omega_p>\\Omega_{\\rm ILR}(m)$, the angular momentum content of $l\\le -1$ components remains zero. Therefore, angular momentum transfer between inner and outer resonances as suggested by LBK, is feasible only for $\\Omega_p < \\Omega_{\\rm ILR}(m)$. This is a constraint on the applicability of LBK's theory.  According to the arguments of \\S\\ref{sec:role-of-ILR} and the WKB theory, the ILR is not transparent to short wavelength disturbances with $X\\gg 1$ \\citep{BT08,Mark74} and any developing spiral structure must be damped unless the ILR remains invisible to the CR. Therefore, an inconsistent point in LBK's paper is the imagination of an angular-momentum-transferring spiral structure while the condition $\\Omega_p <\\Omega_{\\rm ILR}(m)$ has already abandoned the existence of such structures. An issue yet needs to be explained: If there is no spiral structure in stable disks, which mechanism does transfer the angular momentum between inner and outer resonances? In fact, in the limit of $s\\rightarrow 0$ the density components $\\Sigma^{l-}_1$ and $\\Sigma^{l+}_1$ become stationary waves that are azimuthally separated by a phase shift of $90^{\\circ}$ (like the components of mode B1 in Figure \\ref{pic:components-modes-B1-S2}). Thus, they exert opposite gravitational torques on each other without any need for a communicating spiral structure.  In this paper I decomposed unstable modes of a model disk galaxy to its constituent Fourier components and showed how different components experience a gravitational torque. My results clearly showed that only the Fourier component associated with the CR generates a resonant zone in the phase space, and other components exchange angular momentum far from resonances. This result led me to introduce a new dynamical mechanism that triggers unstable modes. According to my calculations, an irreversible resonant capture of stars into the CR causes a synchronous precession of their orbital axes, which in turn, support a rotating density pattern. The emerged pattern grows because the resonant zone expands in the frequency space.  The irreversibility of resonant trapping is the most destructive event that happens in a stellar disk when a group of resonant stars with $d_i<I_1<d_o$ gain angular momentum forever. An immediate consequence of this phenomenon is the angular momentum imbalance in the disk, which requires proper reaction of other stars in order to conserve ${\\cal L}$. This is how the Fourier components of $l\\not =0$ are generated and the angular momentum transfer between them is initiated. Reacting stars do not have a librating $\\theta$ and they cooperate in a way that the quantity $\\vert \\mu^0_m({\\bf \\Omega})\\vert$ remains small for all components.  Using the maps of resonance zones in the phase and frequency spaces, I argued that resonance overlapping should stop the growth of unstable modes and stabilize stellar disks in the presence of an ILR. I showed that for $m>1$ the emergence of an ILR is inevitable in scale-free models. The competition between the CR and ILR can therefore stabilize sufficiently warm scale-free disks against $m>1$ excitations. \\citet{ER98b} suggested in \\S6.1 of their paper that self-consistent scale-free disks (without inner cutouts) do not admit growing non-axisymmetric modes at all, and they ruled out the possibility of a critical temperature. Their prediction is in agreement with my results except in two aspects: (i) Since the line $\\mu^{-1}_{1}({\\bf \\Omega})=0$ does not intersect the frequency space of scale-free models, there is no ILR to compete with a growing CR and density waves will be amplified for $m=1$. (ii) The ILR in cold disks is very special and it cannot develop a resonant zone of finite size. It is therefore hard to believe a serious influence by the ILR on the phase space flows up to a critical temperature.  Note that each unstable mode dominates an isolated region of the frequency (action) space until a chaotic layer emerges due to resonance overlapping. The rate by which the chaotic layer diffuses itself in the phase space, is determined by the resistance of KAM tori against competing resonant zones. The limited space that the fastest growing bar mode occupies after its saturation (Khoperskov et al. 2007), is an $N$-body evidence for such a resistance of spiral modes that fill outer regions of the cored exponential disk.  The other implication of resonance overlapping, which has observational support too, is that the central bar in grand-design barred spiral galaxies must join the spiral arms at the tips of the bar, which are the overlapping regions of two neighboring B and S modes in the frequency space (see Figure \\ref{pic:zones-overlap}). I note that the pattern speeds of the spiral and bar components are not the same because these structures are associated with different modes. The overlapping region is filled by chaotic orbits that yield the turbulent mixing of density waves. This process can enhance star formation near the tips of the bar. If we accept this scenario, barred spiral galaxies should have been formed from hot stellar disks whose eigenfrequency spectra include few spiral modes. On the other hand, flocculent spirals with many wave packets will be the natural destiny of initially cold disks that give birth to a rich family of spiral modes.  The mathematical background for the nonlinear evolution of modes was developed in Paper I but that formulation is valid as long as averaging is allowed over angle variables. In the presence of chaotic orbits it is impossible to average out resonant angles and some modifications are required to deal with the CBE in its full nonlinear form. I will present such modifications in Paper III with the aim of discovering the dynamical processes that generate different classes of barred and spiral structures."
    },
    "0808/0808.3611_arXiv.txt": {
        "abstract": "Using the 100-m radio telescope at Effelsberg, we mapped a large area around the Andromeda Galaxy in the 21-cm line emission of neutral hydrogen to search for high-velocity clouds (HVCs) out to large projected distances in excess of $100~\\mathrm{kpc}$. Our $3 \\, \\sigma$ \\ion{H}{i} mass sensitivity for the warm neutral medium is $8 \\times 10^{4}~{\\rm M}_{\\odot}$. We can confirm the existence of a population of HVCs with typical \\ion{H}{i} masses of a few times $10^{5}~{\\rm M}_{\\odot}$ near the disc of M31. However, we did not detect any HVCs beyond a projected distance of about $50~\\mathrm{kpc}$ from M31, suggesting that HVCs are generally found in proximity of large spiral galaxies at typical distances of a few $10~\\mathrm{kpc}$.  Comparison with CDM-based models and simulations suggests that only a few of the detected HVCs could be associated with primordial dark-matter satellites, whereas others are most likely the result of tidal stripping. The lack of clouds beyond a projected distance of $50~\\mathrm{kpc}$ from M31 is also in conflict with the predictions of recent CDM structure formation simulations. A possible solution to this problem could be ionisation of the HVCs as a result of decreasing pressure of the ambient coronal gas at larger distances from M31. A consequence of this scenario would be the presence of hundreds of mainly ionised or pure dark-matter satellites around large spiral galaxies like the Milky Way and M31. ",
        "introduction": "One of the currently most favoured cosmological models is the so-called Lambda Cold Dark Matter ($\\Lambda$CDM) model. It assumes that the evolution of the universe is dominated by dark energy and dark matter which in sum account for 96~per~cent of the total energy density of the universe \\citep{Spergel2003}. An important prediction of CDM models is the hierarchical formation of gravitationally bound structures. The smallest dark-matter haloes are expected to form first, whereas larger structures, ranging from spiral galaxies to galaxy clusters, are formed at a later stage through merging and accretion of smaller dark-matter haloes (bottom-up scenario). Numerical simulations of structure formation in CDM cosmologies have successfully reproduced the mass function and radial distribution of galaxies on the scales of galaxy clusters. On smaller scales, however, simulations predict significantly more dark-matter haloes than being observed \\citep{Klypin1999,Moore1999}. This discrepancy has been named the `missing satellites' problem.  To overcome this problem, \\citet{Blitz1999} suggested that high-velocity clouds (HVCs) might be the gaseous counterparts of the `missing' dark-matter haloes around the Milky Way. HVCs are gas clouds observed all over the sky in the 21-cm line emission of neutral atomic hydrogen. They were discovered with the Dwingeloo radio telescope by \\citet{Muller1963}, and they are characterised by high radial velocities of typically $|v_{\\rm LSR}| \\gtrsim 100~\\mathrm{km \\, s}^{-1}$ (see \\citealt{Wakker1991a} for details). If HVCs were the `missing' satellites, they could not have experienced significant star formation during their evolution. Therefore, they would appear in the form of pure gas clouds without any noticeable stellar population. In addition, HVCs would be spread all over the Local Group with typical distances of hundreds of kpc and high \\ion{H}{i} masses of about $10^7~{\\rm M}_{\\odot}$.  At the same time, the expected large distances of HVCs from the Milky Way would result in fairly decent angular diameters of the clouds. Therefore, \\citet{Braun1999} defined a sub-sample of compact and isolated HVCs (the so-called CHVCs) which are characterised by angular sizes of less than $2^{\\circ}$~FWHM as well as isolation and separation from neighbouring \\ion{H}{i} emission. The overall kinematics of the CHVC population is consistent with a distribution throughout the Local Group \\citep{deHeij2002b}, making them promising candidates for the `missing' dark-matter satellites around the Milky Way and the Andromeda Galaxy. \\citet{Putman2002} extended the CHVC catalogue into the southern hemisphere, using the \\ion{H}{i} Parkes All-Sky Survey \\citep{Barnes2001}. The data for the northern and southern hemispheres were later combined by \\citet{deHeij2002b} into an all-sky catalogue of 216~CHVCs.  Several arguments have been raised against the idea of HVCs and CHVCs being the `missing' dark-matter haloes predicted by CDM cosmologies. First of all, attempts to identify a population of HVCs in other nearby galaxy groups \\citep{Zwaan2001,Braun2001,Pisano2004} have failed, resulting in upper distance limits for HVCs and CHVCs from the Milky Way of the order of $150~\\mathrm{kpc}$. Additional evidence for HVCs being nearby at distances of the order of only $10~\\mathrm{kpc}$ has been provided by the detection of H$\\alpha$ emission from both HVCs and CHVCs (e.g., \\citealt{Kutyrev1986,Weiner2001,Tufte2002,Putman2003}) and from the determination of distance brackets for several HVC complexes (e.g., \\citealt{Danly1993,Wakker1996,Wakker2007a,Wakker2007b,Thom2006,Thom2008}). Small distances from the Milky Way are also consistent with the head-tail structures found in numerous HVCs and CHVCs (e.g., \\citealt{Bruens2000,Bruens2001,Westmeier2005b}) which are thought to result from ram-pressure interaction of the clouds with the ambient gas of the Galactic corona \\citep{Quilis2001,Konz2002}.  \\begin{table} \\caption{Observational parameters of the northern and southern part of our Effelsberg \\ion{H}{i} blind survey of M31. The final baseline RMS is given for a system temperature of $25~\\mathrm{K}$ at the given velocity resolution. The specified sensitivity is the baseline RMS times the spectral bin width at the original velocity resolution, converted to \\ion{H}{i} column density. Sensitivities for the WNM were calculated for a brightness temperature of $T_{\\rm B} > 3 \\, \\sigma$, a velocity resolution of $20 \\; \\mathrm{km \\, s}^{-1}$, and a line width of $25 \\; \\mathrm{km \\, s}^{-1}$~FWHM.} \\label{tab_obspar} \\begin{center} \\begin{tabular}{lrrl} \\hline parameter                 &  south &   north & unit                          \\\\ \\hline autocorrelator            &    old & \t new &  \t\t\t     \\\\ bandwidth                 &    6.3 & \t  10 & MHz\t\t\t     \\\\ polarisations             & \t 2 & \t   2 &  \t\t\t     \\\\ channels per polarisation &    512 & \t4096 &  \t\t\t     \\\\ velocity resolution       &    2.6 & \t 0.5 & $\\mathrm{km \\, s}^{-1}$       \\\\ frequency switching       & normal & in-band &  \t\t\t     \\\\ total integration time    &    180 & \t 180 & s\t\t\t     \\\\ final baseline RMS        & \t45 & \t  60 & mK\t\t\t     \\\\ sensitivity               & \t21 & \t 5.5 & $10^{16} \\; \\mathrm{cm}^{-2}$ \\\\ WNM sensitivity           &    2.2 & \t 1.3 & $10^{18} \\; \\mathrm{cm}^{-2}$ \\\\ WNM mass sensitivity      & \t 8 & \t   5 & $10^4 \\; {\\rm M}_{\\odot}$\t     \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  The problem of determining the spatial distribution of HVCs can ultimately be solved by searching for the expected HVC population around the nearest large spiral galaxy, the Andromeda Galaxy (M31). First, the distance of M31 is well known so that important physical parameters of HVCs, such as their \\ion{H}{i} mass or their diameter, can directly be determined from the observations. Second, we will look at the HVC population of M31 from the outside, allowing us to determine the (projected) radial distribution of HVCs. Finally, M31 is relatively close to the Milky Way. Therefore, the sensitivity and angular resolution of large single-dish radio telescopes will easily allow us to detect HVCs of the mass and size of the large HVC complexes observed near the Milky Way. The first comprehensive search for HVCs around M31 was carried out by \\citet{Thilker2004} with the 100-m Green Bank Telescope. They mapped an area of $7^{\\circ} \\times 7^{\\circ}$ in the 21-cm line of \\ion{H}{i} with high sensitivity and discovered a population of about 20~HVCs out to the edge of their map at about $50~\\mathrm{kpc}$ projected distance from M31. Their discovery marked the first detection of an extensive HVC population around a galaxy other than our own. Several of these HVCs were studied in detail with the Westerbork Synthesis Radio Telescope (WSRT) by \\citet{Westmeier2005a}.  Thus, the radial extent of the population of HVCs and CHVCs around galaxies like the Milky Way and M31 is confined by two limits. An upper limit of about $150~\\mathrm{kpc}$ can be derived from the non-detection of HVCs in other galaxy groups, and a lower limit of about $50~\\mathrm{kpc}$ is marked by the edge of the \\ion{H}{i} survey of \\citet{Thilker2004} out to which HVCs were found to be present. Therefore, we decided to carry out a complementary \\ion{H}{i} survey with the 100-m radio telescope at Effelsberg to search for HVCs and CHVCs around M31 out to much larger projected distances in excess of $100~\\mathrm{kpc}$. This would allow us to trace the distribution of the HVC population of M31 over its entire radial extent and compare our results with the predictions of CDM structure formation scenarios.  Our paper is organized as follows. In Sect.~\\ref{sect_observations} we explain the technical aspects of our observations. In Sect.~\\ref{sect_reduction} our data reduction, calibration, and analysis strategy is described. In Sect.~\\ref{sect_completeness} we discuss the various completeness issues of our survey and their implications for our results. In Sect.~\\ref{sect_results} we describe the derived observational and physical properties of the high-velocity clouds detected in our survey. In Sect.~\\ref{sect_comparison} we discuss the results of our comparison of the observational parameters of the HVCs near M31 with different CDM-based models and with the distribution of satellite galaxies around M31. Sect.~\\ref{sect_origin} discusses the evidence for different hypotheses on the origin of HVCs. Finally, Sect.~\\ref{sect_summary} summarises our results and conclusions.   \\begin{figure*} \\includegraphics[width=0.74\\linewidth]{westmeier2.eps} \\caption{The plot shows the geometry of the field (thick solid line) mapped around M31. The dashed circles separate the five different regions (labelled with numbers) for which individual spatial completeness calculations have to be performed.} \\label{fig_spatcomp} \\end{figure*} ",
        "conclusions": "\\label{sect_summary}  We used the 100-m radio telescope at Effelsberg to map a large area around the Andromeda Galaxy, M31, in the 21-cm line emission of neutral atomic hydrogen. Our survey extends out to a projected distance of about $140~\\mathrm{kpc}$ in the south-eastern direction (equivalent to about two thirds of the projected distance towards M33) and about $70~\\mathrm{kpc}$ in the north-western direction. With this map layout we are able to fill the previously existing gap between the outer boundary of the GBT survey of \\citet{Thilker2004} at about $50~\\mathrm{kpc}$ projected distance from M31 and the upper limit of about $150~\\mathrm{kpc}$ for the distance of HVCs as derived from the non-detections in nearby galaxy groups \\citep{Zwaan2001,Braun2001,Pisano2004}. The achieved spectral baseline RMS is $45~\\mathrm{mK}$ at $2.6~\\mathrm{km \\, s}^{-2}$ velocity resolution, corresponding to a $3 \\, \\sigma$ \\ion{H}{i} column density detection limit of $2.2 \\times 10^{18}~\\mathrm{cm}^{-2}$ for the warm neutral medium ($\\Delta v = 25~\\mathrm{km \\, s}^{-1}$ FWHM). This translates into an \\ion{H}{i} mass sensitivity of $8 \\times 10^{4}~{\\rm M}_{\\odot}$.  In total, we detected 17~individual HVCs and several regions of more extended extra-planar gas all around M31. The discrete clouds are predominantly unresolved by the HPBW of $9~\\mathrm{arcmin}$ and characterised by typical \\ion{H}{i} masses of a few times $10^{5}~{\\rm M}_{\\odot}$. We did not detect any clouds beyond a projected distance of about $50~\\mathrm{kpc}$, suggesting that HVCs are generally found in proximity of their host galaxies. In particular, we did not find an extended populations of hundreds of CHVCs as observed around the Milky Way, suggesting that the Galactic CHVCs are intrinsically small clouds in the immediate vicinity of the Milky Way.  A comparison with the Local Group population model of CHVCs, as proposed by \\citet{deHeij2002b}, reveals that their best-fitting model {\\#}9 is discarded by our data with high confidence. Neither the observed projected radial distribution nor the \\ion{H}{i} mass function of HVCs around M31 can be explained by the model. Instead, we find that a Gaussian radial scale length of the order of $50~\\mathrm{kpc}$ can best explain the observed projected distribution of HVCs around M31. In addition, we find that the HVCs are also distinct from the M31 satellite galaxies through typically lower \\ion{H}{i} masses and smaller projected distances from M31.  CDM-based structure formation simulations by \\citet{Kravtsov2004} suggest that about 50 to 100~dark-matter haloes with total gas masses of greater than $10^{6}~{\\rm M}_{\\odot}$ should exist within about $300~\\mathrm{kpc}$ of M31. Only~2 to~5 of these haloes should be located within $50~\\mathrm{kpc}$ of M31, suggesting that some of the HVCs found near M31 could instead be tidally stripped gas from present or former satellite galaxies of M31. This idea is supported by our high-resolution follow-up observations of several HVCs \\citep{Westmeier2005a}.  The lack of detections beyond $50~\\mathrm{kpc}$ projected radius, however, is in conflict with the predictions made by \\citet{Kravtsov2004}. A possible explanation of this discrepancy could be ionisation as a result of decreasing pressure of the ambient coronal medium at larger distances from M31, as suggested by hydrostatic simulations of \\citet{Sternberg2002}. An important consequence of this scenario would be the presence of hundreds of mainly ionised or pure dark-matter satellites near large spiral galaxies, such as the Milky Way and M31, which would be undetectable in the 21-cm line of neutral hydrogen. Finding these almost invisible satellites would be an important observational result in support of CDM cosmologies.  Another promising scenario recently discussed by \\citet{Bekki2008} is a possible tidal interaction between M31 and M33 during their previous encounter. This would have stripped gas from M33 which could have formed some of the HVCs observed near M31. More detailed and extensive observations and simulations will be required to further investigate this interesting scenario.  With our Effelsberg \\ion{H}{i} survey of M31 we have shown that HVCs are most likely concentrated around the large spiral galaxies with typical distances of no more than a few $10~\\mathrm{kpc}$. It is likely that different physical processes, such as tidal stripping, accretion of primordial dark-matter haloes, or galactic outflows, have contributed to the HVC populations of M31 and the Milky Way. The general volatility of these processes could naturally explain some of the differences between the HVC populations of M31 and the Milky Way, for example the extended and complex gaseous streams of the Magellanic Clouds which have no counterpart in M31."
    },
    "0808/0808.3561_arXiv.txt": {
        "abstract": "The stellar asymmetry of faint thick disk/inner halo stars in the first quadrant ({\\it l} = 20 --45$\\arcdeg$) first reported by \\cite{lar96} and investigated further by  \\cite{par03,par04} has recently been confirmed by the SDSS \\citep{jur08}. Their interpretation of the excess in the star counts as a ringlike structure, however, is not supported by critical complementary data in the fourth quadrant, not covered by the SDSS.  We present stellar density maps from the Minnesota Automated Plate Scanner (MAPS) Catalog of the POSS I showing that the overdensity does not extend into the fourth quadrant. The overdensity is most probably not a ring. It could be due to interaction with the disk bar, evidence for a triaxial thick disk, or a merger remnant/stream. We call this feature the Hercules Thick Disk Cloud. ",
        "introduction": "\\cite{lar96} initially reported a substantial asymmetry of faint blue stars in the first quadrant (Q1) of the inner Galaxy,  $l = 20\\arcdeg - 45\\arcdeg$ compared with complementary fields in the fourth quadrant (Q4) based on star counts from MAPS\\footnote{The Minnesota Automated Plate Scanner Catalog of the POSS I is online at:  http://aps.umn.edu} \\citep{cab03}. \\cite{par03} made a more in-depth survey to map the extent of the asymmetry using 40 contiguous fields in each of three regions:  Q1 above and below the plane and Q4 above the plane.   Q4 below the plane is not covered in the POSS I.  They found a 25\\% excess in the number of  probable thick disk stars in Q1 above \\it{and }\\rm below the plane when compared to the complementary Q4 fields.  The region was irregular in shape and covered several hundred square degrees, but with a completeness limit at $\\approx$ 18--18.5  mag, the stars showing the excess in Q1 were relatively nearby, $\\sim$ 1 -- 2 kpc from the Sun.  \\cite{par04} also found an associated kinematic signature, a significant lag  of  ~80 to 90 km/sec in the direction of Galactic rotation for the associated thick disk stars in Q1.  The recent release of the SDSS Data Release 5 (DR5) photometry in the direction of the observed asymmetry in Q1 led to the discovery of a feature at much fainter magnitudes, the distant Hercules-Aquila cloud \\citep{bel07} and the photometric parallax study by \\cite{jur08} confirmed our nearer asymmetry in the inner Galaxy as an overdensity at a galactocentric radius of 6.5 kpc situated 1.5 kpc above the plane. The SDSS survey however is not well designed for a good study of the thick disk inside the Solar orbit. It extends below $b = 30\\arcdeg$ in only a few directions in Q1 and has only limited coverage in Q4.  We are continuing our program of photometric and spectroscopic observations to map the size and extent of the asymmetry along our line of sight and to determine the degree of spatial and kinematic asymmetry above and below the plane.  We are using the SMARTS Consortium CTIO 1-meter Y4KCam and the Steward Observatory 90\" Bok telescope 90Prime Mosaic imager to obtain wide-field multi-color CCD imaging to fainter completeness limits than the POSS I.  Spectra of the candidate thick disk stars for radial velocities and metallicity estimates have been observed using the Hydra multi-object spectrometer on the CTIO Blanco 4-meter telescopes and the Hectospec on the MMTO 6.5 meter.  In this Letter we present a stellar number density map we created from the MAPS POSS I scans to provide a more global reference for our current deeper but more spatially restricted photometric and spectroscopic study of the thick disk in the inner Galaxy.  Other works \\citep{xu07} have used plate data to supplement gaps in the SDSS at more southern declinations with good success. Our map of the stellar density in Q1 and Q4  covers much of the sky unavailable to SDSS and  demonstrates that the nearby asymmetry in Q1 does not represent a ring above the Galactic plane, but instead is a significant substructure or cloud extending over many square degrees in galactic longitude and latitude. We call this feature the Hercules Thick Disk cloud. ",
        "conclusions": "After the reduction steps outlined above, we then binned the stars $0.25 \\arcdeg \\times 0.25 \\arcdeg$ in {\\it l} and {\\it b} to create maps of the stellar density distribution of the faint blue and intermediate color stars shown in Figures~\\ref{fig5} and ~\\ref{fig6}. The figures are color-coded with respect to number density per 0.0625 square degree.  \\cite{jur08} described the excess in Q1 as due to a ``ringlike\" structure because the overdensity appeared to be radially constant and circular in cross section in their Figure 27. The center of the overdensity region ranges from $(X,Y,Z)=(6.5 kpc, -2.2 kpc, 1.5 kpc)$ to $(6.5 kpc, 0.3 kpc, 1.5 kpc)$ and can easily be converted into galactic coordinates. This is shown as the purple line on Figure~\\ref{fig5}. If the feature were symmetric about $l=0\\arcdeg$ it would project into Q4 to $(6.5 kpc, 2.2 kpc, 1.5 kpc)$.  This projection is also shown as a purple line on Figure~\\ref{fig6}. Comparison of Figures~\\ref{fig5} and ~\\ref{fig6}, however show that the density of these stars is not symmetric with respect to the Sun-Center line. There is a clear excess of stars in Q1 over Q4 in the range {\\it l} $= 25$ -- $45\\arcdeg$ and {\\it b} $= 30$ -- $40\\arcdeg$. Furthermore, we emphasize that the excess would not have been initially discovered if it had been a symmetric ring since we \\citep{lar96} were comparing star counts for complementary fields in Q1 and Q4.  Could the ``ringlike\" structure be inclined to the plane and therefore not visible in Q4?  Most probably not.  The full height of the overdensity in Z over an azimuthal distance of 2500 parsecs is only 500 pc (Juric{\\c'} et al.). Even for the pathological case of a paper thin inclined distribution in Z the maximum inclination could only be 11 degrees and given its width in X it should have been visible in Figure~\\ref{fig6}. Additionally, there is strong evidence from Juri{\\'c} et al.'s Figure 27 (left panel) that for X = 7250-7750 parsecs the overdensity is exclusively above $Z = 1500$ pc.  Examination of the same figure's middle panel shows that all significant contributors to the overdensity in this same X range have $Y < 1000$ pc.  This would be the opposite of what should be happening if the overdensity were falling into the plane in Q4.  Finally, \\cite{par04} studied the velocities of samples of stars taken from overdensity regions in Q1 compared with a control sample Q4 and found in a somewhat weak result that the Z component of velocity was less negative for Q1 than it was for Q4. If a coherent population of stars were moving together towards the disk from above the plane as they entered Q4, the opposite should be true.  The cloud of stars detected by Juri{\\'c} et al. is not symmetric about the $l=0\\arcdeg$ line and almost certainly is not a ring. This does not change their other possible explanation for the feature, however.  Given the broad extent of the cloud (Figures~\\ref{fig5} and ~\\ref{fig6}) together with its apparent small ranges in radial distance from the Sun and its distance above the galactic plane (their Figure 27), the Hercules Thick Disk cloud may be a debris stream consistent with the disk formation scenario described by \\cite{abd03}. While the Hercules Thick Disk Cloud is relatively nearby on the sky it is not related to the more distant (10 - 20 kpc) Hercules-Aquila cloud of \\citep{bel07}. The northern extent of the Hercules-Aquila cloud (Figure 2 in \\citet{bel07}) is confined to galactic latitudes less than $30\\arcdeg$ and even in those regions the bulk of the stars are much fainter than our magnitude limit.  The contamination of our sample by Hercules-Aquila cloud stars would be less than 5 objects per square degree (0.5 stars/bin) in any case given our relatively bright magnitude limits.  Other explanations are still possible such as a triaxial thick disk and an overdensity due to an interaction with the disk bar. Analysis of our CCD photometry and spectroscopy for fainter stars in Q1 and Q4 will be used to address this question."
    },
    "0808/0808.0081_arXiv.txt": {
        "abstract": "In this work we estimate the radius and the mass of a self-gravitating system made of axions. The quantum axion field satisfies the Klein-Gordon equation in a curved space-time and the metric components of this space-time are solutions to the Einstein equations with a source term given by the vacuum expectation value of the energy-momentum operator constructed from the axion field. As a first step towards an axion star we consider the up to the $\\phi^6$ term in the axion potential expansion. We found that axion stars would have masses of the order of asteroids ($\\sim 10^{-10}M_\\odot$) and radius of the order $\\sim \\mbox{few} ~$centimeters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.1805_arXiv.txt": {
        "abstract": "The brick wall method in calculations of the entropy of black holes can be applied to the FRW cosmology in order to study the statistical entropy. An appropriate cutoff satisfying the covariant entropy bound can be chosen so that the entropy has a definite bound. Among the entropy for each of cosmological eras, the vacuum energy-dominated era turns out to give the maximal entropy which is in fact compatible with assumptions from the brick wall method. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.3207_arXiv.txt": {
        "abstract": "% The dark clouds in the constellation of Chamaeleon have distances of 160-180~pc from the Sun and a total mass of $\\sim$5000~$M_\\odot$. The three main clouds, Cha~I, II, and III, have angular sizes of a few square degrees and maximum extinctions of $A_V\\sim5$-10. Most of the star formation in these clouds is occurring in Cha~I, with the remainder in Cha~II. The current census of Cha~I contains 237 known members, 33 of which have spectral types indicative of brown dwarfs ($>$M6). Approximately 50 members of Cha~II have been identified, including a few brown dwarfs. When interpreted with the evolutionary models of Chabrier and Baraffe, the H-R diagram for Cha~I exhibits a median age of $\\sim$2~Myr, making it coeval with IC~348 and slightly older than Taurus ($\\sim$1~Myr). The IMF of Cha~I reaches a maximum at a mass of 0.1-0.15~$M_\\odot$, and thus closely resembles the IMFs in IC~348 and the Orion Nebula Cluster. The disk fraction in Cha~I is roughly constant at $\\sim50$\\% from 0.01 to 0.3~$M_\\odot$ and increases to $\\sim65$\\% at higher masses. In comparison, IC~348 has a similar disk fraction at low masses but a much lower disk fraction at $M\\ga1$~$M_\\odot$, indicating that solar-type stars have longer disk lifetimes in Cha~I. ",
        "introduction": "The southern constellation of Chamaeleon contains one of the nearest groups of dark clouds to the Sun ($d\\sim160$-180~pc). An extinction map of these clouds is shown in Figure~\\ref{fig:map1} \\citep{dob05}. The main clouds in Chamaeleon have angular sizes of a few square degrees and are referred to as Chamaeleon I, II, and III \\citep{hof62}.\\footnote{In most publications, including this review, the designations Cha~II and Cha~III have been reversed from the original ones assigned by \\citet{hof62}.} Digitized Sky Survey (DSS) images of Cha~I and II are shown in Figure~\\ref{fig:map2} and a color optical image of Cha~I obtained by G.\\ Rhemann is shown in Figure~\\ref{fig:map3}. These clouds contain signposts of recent star formation in the form of several reflection nebulae, including Ced~110, 111, \\citep{ced46} and the Infrared Nebula \\citep[IRN,][]{sh83}. An optical image of the area surrounding Ced~111 is shown in Figure~\\ref{fig:ced111}. Newborn stars were first directly discovered in Chamaeleon through their variability and H$\\alpha$ emission \\citep{hof62,hen63,men72}. The masses of the Chamaeleon clouds and the stellar densities of young stars within them are low compared to many other star-forming regions. Because Chamaeleon is nearby and well-isolated from other young stellar populations, it has been a popular target for studies of low-mass star formation.  \\begin{figure} \\includegraphics[width=\\textwidth]{f1.eps} \\caption{Extinction map of the Chamaeleon dark clouds \\citep{dob05}. The maximum extinction in this map is $A_V\\sim$10.} \\label{fig:map1} \\end{figure}  \\begin{figure} \\includegraphics[width=\\textwidth]{f2.eps} \\caption{ DSS images of Cha~I ($2\\hbox{$^\\circ$}\\times2\\hbox{$^\\circ$}$) and Cha~II ($1\\fdg5\\times1\\fdg5$). The reflection nebulae Ced~111 and Ced~112 are associated with the B stars HD~97048 and HD~97300, respectively. } \\label{fig:map2} \\end{figure}  \\begin{figure} \\includegraphics[width=\\textwidth]{f3.eps} \\caption{ A wide-field optical color-composite image of the Cha~I cloud ($1\\fdg4\\times2\\hbox{$^\\circ$}$). North is up and east is left. Courtesy G. Rhemann. } \\label{fig:map3} \\end{figure}  \\begin{figure} \\includegraphics[scale=1.08]{f4.eps} \\caption{ An optical color-composite image of the Ced~111 reflection nebula in Cha~I obtained with VLT ($6\\farcm8\\times11\\farcm2$). Four of the brightest young stars within this area are labeled. North is up and east is left. Courtesy ESO. } \\label{fig:ced111} \\end{figure} ",
        "conclusions": ""
    },
    "0808/0808.0132.txt": {
        "abstract": "The GHASP survey (Gassendi HAlpha survey of SPirals) consists of 3D \\ha~data cubes for 203 spiral and irregular galaxies, covering a large range in morphological types and absolute magnitudes, for kinematics analysis. It is the largest sample of Fabry-Perot data published up to now. In order to provide an homogenous sample, reduced and analyzed using the same procedure, we present in this paper the new reduction and analysis for a set of 97 galaxies already published in previous papers but now using the new data reduction procedure adopted for the whole sample. The GHASP survey is now achieved and the whole sample is reduced using adaptive binning techniques based on Voronoi tessellations. We have derived \\ha~data cubes from which are computed \\ha~maps, radial \\VFs~as well as residual \\VFs, \\PVMs, \\RCs~and kinematical parameters for almost all galaxies. The \\RCs, the kinematical parameters and their uncertainties are computed homogeneously using the new method based on the power spectrum of the residual \\VF. This paper provides the kinematical parameters for the whole sample. For the first time, the integrated \\Ha~profiles have been computed and are presented for the whole sample. The total \\Ha~fluxes deduced from these profiles have been used in order to provide a flux calibration for the 203 GHASP galaxies. This paper confirms the conclusions already drawn from half the sample concerning (i) the increased accuracy of \\PAs~measurements using kinematical data, (ii) the difficulty to have robust determinations of both morphological and kinematical inclinations in particular for low inclination galaxies and (iii) the very good agreement between the \\TF~relationship derived from our data and previous determinations found in the literature. ",
        "introduction": "\\label{intro}  The GHASP survey consists of a large sample of spiral and irregular galaxies observed with a scanning Fabry-Perot for studying their kinematical and dynamical properties through the ionized hydrogen component. The goals of this study have been described in \\citet{Epinat:2008}. % The GHASP sample is by now the largest homogenous sample of Fabry-Perot data ever published, comprising 3D data for 203 galaxies. % This paper is the seventh of a series called hereafter Paper I to VI \\citep{Garrido:2002,Garrido:2003,Garrido:2004,Garrido:2005,Spano:2008, Epinat:2008} presenting the data obtained in the frame of the GHASP survey.  The observations were lead during fourteen observing runs at the ``Observatoire de Haute Provence (OHP)'', France, from 1998 to 2004. The survey is now achieved. % The observing runs number 8 to 14 have been presented in Paper VI, which relies on a set of 108 galaxies, providing 106 \\VFs~and 93 \\RCs.  Those data have been reduced with the new data reduction procedure (see Paper VI and references therein). % The data presented in this paper are those of the seven first observing runs (already presented in Papers I to IV) that have been re-reduced using the same method as in Paper VI. It also contains data for UGC 3382 and UGC 11300 that have been improved by adding new data (runs 3, 5 and 6) to the ones already published in Paper VI (runs 10 and 13). Thus, the data presented here consist of a set of 97 galaxies, providing 96 \\VFs~and 82 \\RCs.  To be clear on the goals and limits of this present work, we summarize hereafter what we present and what we do not in this paper. We present: \\begin{itemize} \\item the new individual maps and \\PVMs~in Appendix \\ref{maps} (on line data only), \\item the new \\rcs, in Appendix \\ref{rc} and the new corresponding tables in Appendix \\ref{rc_tables}. \\end{itemize} In this paper, we lead the same kind of analysis as the one presented in Paper VI concerning: \\begin{itemize} \\item the study of the parameters of the kinematical models, \\item the study of the residual \\VFs, \\item the \\TF~relation. \\end{itemize} Because it is useful to display and analyze all the data together, we put here in the same tables (in Appendix \\ref{tables}) the new parameters and the results already published in Paper VI so that the reader does not have to compile tables coming from different publications. With respect to Papers I to IV, some distances and absolute magnitudes have been recomputed (using better estimations). %Furthermore, all the parameters for the whole GHASP sample are in %a single table and the reader does not have to compile tables %coming from different publications. %Thus we provide: %\\begin{itemize} %    \\item all the parameters of the kinematical models (in section \\ref{xxx} and \\ref{tables}). %    \\item all the residual \\VFs~(in section \\ref{xxx}), %    \\item the \\TF~relation. %\\end{itemize} For the whole GHASP sample we make a new analysis on an absolute flux calibration made using the data calibrated by \\citet{James:2004} and the integrated \\ha~profiles deduced from our data cubes.  Because this has been discussed in previous papers, we do not present any more: \\begin{itemize} \\item the morphological types and luminosity distributions of the whole GHASP sample (see Paper VI), \\item the data reduction procedure used here, including the computation of the \\RCs, the determination of the kinematical parameters and the determination of the uncertainties (see Paper VI), \\item the instrumental set-up of the instrument for the data re-reduced in this paper (see Papers I to IV), %    \\item the log of the observation (see Papers I to IV) xxxpourtant on remet des info recapitulatives en table B1: que dire donc?xxx, \\item the individual comments for each galaxy (see Papers I to IV), except when the new reduction procedure leads to new comments or to conclusions noticeably different from the previous ones (see Appendix \\ref{notes}). \\end{itemize}  In section \\ref{calibration} we make an indirect flux calibration of the \\ha~profiles. In section \\ref{analysis} we present the data and in section \\ref{tullyfisher} we compute the \\TF~relation. In section \\ref{conclusion} we give the summary and conclusions. When the distances of the galaxies are not known, a Hubble constant H$_{0}$=75\\kms~Mpc$^{-1}$ is used throughout this paper. ",
        "conclusions": "\\label{conclusion}  The knowledge of the links between the kinematical and dynamical state of galaxies helps us to increase our understanding of the physics and evolution of galaxies. The GHASP sample, which consists of 203 spiral and irregular galaxies, covering a wide range of morphological types and absolute magnitudes, has been constituted in order to provide a kinematical reference sample of nearby galaxies. The galaxies have been observed in the \\ha~line using Fabry-Perot techniques, leading to the construction of data cubes.  This sample is by now the largest set of galaxies ever homogeneously observed with Fabry-Perot techniques.  Major improvements in the reduction (adaptive binning techniques, ghost suppression, treatment of faint outskirts regions, etc) and in the analysis (determination of the \\RC~and of the kinematical parameters and their uncertainties, etc) have been developed and implemented in the data reduction procedure and homogeneously applied to the whole GHASP sample (see Paper VI for additional details).  In this paper, 97 galaxies have been re-reduced using adaptive binning techniques in order to provide homogeneous data for the whole sample. For each galaxy, we have presented the \\ha~\\VF, the \\ha~monochromatic image and eventually the \\ha~residual \\VF, the \\PVM~along the major axis and the \\RC, when available, leading for the whole sample to 200 \\VFs~and 177 \\RCs.  From the data cubes, integrated \\Ha~profiles have also been produced. A post calibration has allowed to compute indirect absolute \\Ha~flux for all the galaxies belonging to the GHASP sample.  This post calibration has been done using fluxes for 69 galaxies found in the literature \\citep{James:2004}.  We confirm and strengthen most of the results already obtained from half the sample:  \\begin{itemize}  \\item A high quality model has been achieved to represent the axi-symmetric rotational component of the galaxies since no typical signatures for biases are observed in the residual \\VFs. This means that the residuals observed in the residual \\VF~are due to actual non circular motions and not to an uncorrect determination of the kinematical parameters (position of the center, \\PA, inclination and rotation velocity). In addition, the \\PVMs~confirm the validity of the \\RCs.  \\item The mean residual velocity dispersion is strongly correlated with the maximum amplitude of the \\VF. For a given velocity range, this correlation does not clearly depend on the morphological type. However strongly barred galaxies have a higher residual velocity dispersion than mild-barred or unbarred galaxies. Peculiar galaxies also show a high residual velocity dispersion.  \\item The determinations of kinematical \\pas~are robust whatever the inclination of the galaxy whereas morphological \\pas~are poorly determined for low inclination systems. Moreover, morphological \\PAs~have systematically higher uncertainties than kinematical ones. This is a major argument for deriving \\rcs~from integral field spectroscopy rather than long slit spectroscopy instruments that could lead to incorrect positioning of the slit (a difference between the morphological and kinematical \\PAs~larger than 30\\degr~is found for $\\sim$15$\\%$ of the GHASP sample). This may strongly bias mass distribution models and \\TF~studies. In order to build a mass model, the stellar mass distribution derived from the surface brightness profile is combined with the \\rc~deduced from the \\vf. The \\pas~of the major axis deduced from the surface brightness image and from the \\vf~should be identical. Important inconsistencies may appear if these \\PAs~are misaligned.  \\item Galaxies with poor determination of their morphological \\pas~have usually unreliable and overestimated morphological inclinations. The agreement between kinematical and morphological inclinations is better when assuming a thin disk in particular for high inclination galaxies. For galaxies with intermediate disk inclinations (higher than 25\\degr~and lower than 75\\degr), to improve the quality of the \\RC, it is possible to reduce the degrees of freedom in kinematical models by fixing the inclination to the morphological value. This is specially true when only low quality kinematical data are available as it is the case for high redshift galaxies.  \\item The use of the whole GHASP sample leads to a \\TF~relationship in perfect agreement with \\citet{Tully:2000}, despite important differences in the selection of both samples. With respect to the result presented in Paper VI, the use of the whole sample increases the agreement with \\citet{Tully:2000}. Three comments should be underlined: (i) galaxies with inclination lower than 25\\degr~are inappropriate for \\TF~relation determination since their estimated velocities are easily overestimated; (ii) fast rotators (V$_{max}>$300\\kms) are maybe less luminous (than expected from the \\TF~relation); (iii) for fast rotators and high luminosity galaxies, the agreement with the \\TF~relation is better when the morphological inclination of the galaxy is computed without taking into account the increasing thickness of the disk when the morphological type of the galaxies moves from early to late types.  \\end{itemize}  From these data and analysis, it is now possible to adress the scientific drivers on the whole GHASP sample in forthcoming papers."
    },
    "0808/0808.1728_arXiv.txt": {
        "abstract": "\\noindent   We present comprehensive models for the Herbig Ae stars MWC275 and AB~Aur that aim to explain their spectral energy distribution (from UV to millimeter) and long baseline interferometry (from near-infrared to millimeter) simultaneously. Data from the literature, combined with new mid-infrared (MIR) interferometry from the Keck Segment Tilting Experiment, are modeled using an axisymmetric Monte Carlo radiative transfer code. Models in which most of the near-infrared (NIR) emission arises from a dust rim  fail to fit the NIR spectral energy distribution (SED) and  sub-milli-arcsecond NIR CHARA interferometry. Following recent work, we include an additional gas emission component with  similar size scale to the dust rim,  inside the sublimation radius, to fit the NIR SED and long-baseline  NIR interferometry on MWC275 and AB~Aur. In the absence of shielding of star light by gas, we  show that the gas-dust transition region in these YSOs will have to contain highly refractory dust, sublimating at   $\\sim$1850K. Despite having nearly identical structure in the thermal NIR, the outer disks of MWC275 and AB~Aur differ substantially. In contrast to the AB~Aur disk, MWC275 lacks small grains in the disk atmosphere capable of producing significant 10-20 $\\mu$m emission beyond $\\sim$7AU, forcing the outer regions into the ``shadow'' of the inner disk. ",
        "introduction": "\\label{intro} Herbig Ae (HAe) stars are pre-main-sequence stars of intermediate mass (1.5-3 solar masses). They exhibit a robust excess in emission over stellar photospheric values from near-infrared (NIR) to the millimeter (mm) wavelengths. This excess is now attributed to the passive reprocessing of stellar light by dust in the circumstellar environment \\citep{lkha2001, natta2001, dullemond2001}. The geometry of the circumstellar environment of HAe stars has been actively debated in the astronomy community over the last two decades. Some of the early workers  in this field \\citep{Hillen92} showed that spectral energy distribution (SED) of  HAe stars could be explained by emission from circumstellar matter in disk-like geometry. Others \\citep{Mirosh97} argued that the emission could also arise from dust in a spherical geometry around the star, proving the inadequacy of SED modeling alone in uniquely fixing the geometry of the circumstellar matter. The first observational evidence in favor of a disk geometry came from millimeter (mm) interferometry in the form of asymmetries detected \\citep{Man97} in the mm images. Asymmetries in the NIR emission were also detected by the Palomar Test-Bed Interferometer \\citep{eisner2003, eisner2004}, settling the debate in support of a disk geometry for circumstellar material in Herbig Ae stars.  Most interferometric studies of HAe stars have relied on simple geometric models \\citep{Man97, rmg1999, rmg2001, eisner2003, eisner2004, monnier2005}  that explain the emission geometry of the system in only narrow wavelength ranges. This method, albeit extremely useful in elucidating some of the morphology details, is not adequate for exploring the interdependency in structure of the inner and outer parts of the disk. A number of studies \\citep{dullemond2001, dullemond2004, boekel2005b} have shown that the structure of the inner disk at fractions of an AU scale clearly affects the structure of the outer disk. A complete understanding of the circumstellar disk structure in HAe stars therefore requires models that simultaneously explain the SED and interferometry over a large wavelength range. Such models have begun to appear in the literature only recently \\citep{ponto2007, kraus2007}.   In this paper, we develop comprehensive disk models to  explain the SED and interferometry of the HAe stars MWC275 and AB Aur.   MWC275 and AB~Aur are prototype pre-main-sequence stars of similar ages and  spectral type  with extensive circumstellar disks. Due to the availability of photometric and interferometric data over a large wavelength range, MWC275 and AB~Aur are ideal candidates for testing disk models  for YSOs. The extent  of their circumstellar-dust disks was first measured by \\citet{Man97} to be several 100AU using the Owens Valley Radio Observatory (OWRO). \\citet{Natta2004} resolved the MWC275 disk in the mm and reported a de-convolved,  projected dust-disk size of 300AU$\\times$180AU.  More recently, \\citet{isella2007} analyzed IRAM, SMA and VLA continuum and $^{12}$CO, $^{13}$COand $^{18}$CO line data  constraining the gas-disk radius to be 540AU with the gas in Keplerian rotation around the central star. Scattered light studies of MWC275 \\citep{grady2000} and AB~Aur \\citep{grady1999, oppenheimer2008} show the presence of arcs and rings in the circumstellar disk.  \\citet{corder2005} resolved the AB~Aur CO disk radius to be $\\sim$600AU, finding strong evidence for Keplerian rotation for the bulk of the disk.  \\citet{corder2005} and \\citet{lin2006} detected spiral arms in CO emission with radii of  $\\sim$150AU, while \\citet{fukagawa2004} detected similar structure in Subaru H-band scattered light images. AB~Aur also has substantial envelope material on scales larger then 600 AU \\citep{grady1999, semenov2005, corder2005, lin2006}.  MIR emission probes the giant planet  formation region in circumstellar disks \\citep{calvet1992, cg97, dullemond2001} with the emission arising from warm dust (T $>$ 150K). \\citet{Meeus2001} and \\citet{boekel2005} used the 10$\\mu$m MIR silicate emission feature from MWC275 and AB~Aur to show that dust grains in these systems had  grown larger than the typical interstellar medium grain sizes. \\citet{marinas2006} imaged AB~Aur  at 11.7$\\mu$m and found  the  emission FWHM size to be 17$\\pm$4 AU consistent with the flared disk models of \\citet{dullemond2004}.  In this paper,  we present  new 10$\\mu$m  measurements of AB~Aur  and MWC275 with the Keck Segment Tilting Experiment \\citep[described in \\S\\ref{data}]{Monnier2004}.   In contrast to AB~Aur, the MWC275 disk is unresolved by the Segment Tilting Experiment (maximum baseline of 10m),  requiring the VLT Interferometer (100m baseline) to probe to its MIR  structure \\citep{Leinert2004}. These observations suggest that  MWC275 disk differs considerably from AB~Aur and we present  a detailed comparison of the two disk structures in the discussion  (\\S\\ref{discuss}).  Thermal NIR emission probes hot regions (typically the inner AU) of the disk with temperatures greater than 700K. The NIR disks of MWC275 and AB~Aur were first resolved with IOTA by \\cite{rmg1999, rmg2001} and subsequently observed at higher resolution with PTI \\citep{eisner2004}, Keck Interferometer \\citep{monnier2005} and the CHARA  interferometer array \\citep{Tannirk2008}. In \\citet[hereafter T08]{Tannirk2008} we showed that inner-disk models in which majority of the K-band emission arises in a dust rim \\citep{dullemond2001, isella2005, Tannirk2007} fail to fit the CHARA data at milli-arcsecond resolution. We also demonstrated that the presence of additional NIR emission (presumably from hot gas) inside the dust destruction radius can help explain the CHARA data and the NIR SED. First calculations for the effects of gas on  rim  structure \\citep{muzerolle2004}  showed that for plausible disk parameters,  presence of gas does not modify dust-rim geometry significantly. Besides a poorly understood interferometric visibility profile, MWC275 also displays as yet ill-understood NIR and MIR SED time variability \\citep{sitko2007} which has been interpreted as variations of the inner disk structure. In  \\S\\ref{MWC275_model} and \\S\\ref{ABAur_model} we present a detailed analysis of  the NIR visibility and SED for MWC275 and AB~Aur, placing constraints on the wavelength dependence of the opacity source inside the dust destruction radius.    In this study, we focus on (i) explaining the inner-disk structure and discuss important open problems and  (ii) modeling the MIR emission morphology of the disks and the shape of the MIR spectrum. The paper is organized into 7 sections with \\S\\ref{data} detailing the observations. \\S\\ref{model} explains the disk model and the modeling strategy. \\S\\ref{MWC275_model} and \\S\\ref{ABAur_model} analyze MWC275 and AB~Aur SED and visibilities in relation to the  disk models. We present a discussion on our results  and our conclusions  in \\S\\ref{discuss} and \\S\\ref{conclusion} respectively. ",
        "conclusions": "\\label{conclusion} We have presented the first set of comprehensive disk models for the SED and interferometry of Herbig Ae stars MWC275 and AB~Aur. We have shown that  `standard'  models  for the dust evaporation front   where the bulk of the near-infrared emission arises from a dust wall, fail to explain the near-infrared spectral energy distribution and interferometry. Standard  models produce large bounces in visibility at high spatial frequency, which is not observed in the data.  We have conclusively demonstrated that the presence of an additional smooth emission component (presumably hot gas) inside the dust destruction radius and on a similar size scale to the dust rim  can ameliorate the situation. In the absence of shielding of star light by gas, we have established that dust grains in the gas-dust transition region will have to be highly refractory, sublimating at ~1850K. The small mid-infrared size of MWC275 relative to AB~Aur, shows that the dust grains in the outer disk MWC275 are  significantly more evolved/settled than the  grains  in the AB~Aur disk. We suggest that dynamical processes (like planetesimal collisions) that maintain the  population of micron-sized grains producing the 10$\\mu$m feature in the spectrum, are operational only in the inner ~7AU of MWC275. However, in AB-Aur the small-dust producing mechanisms exist at least out to ~20 AU and maybe even beyond."
    },
    "0808/0808.1172_arXiv.txt": {
        "abstract": "{}{The high-redshift ($z=4.048)$ gamma-ray burst GRB 060206 showed unusual behavior, with a significant rebrightening by a factor of $\\sim4$ at about 3000 s after the burst. We argue that this rebrightening implies that the central engine became active again after the main burst produced by the first ejecta, then drove another more collimated jet-like ejecta with a larger viewing angle. The two ejecta both interacted with the ambient medium, giving rise to forward shocks that propagated into the ambient medium and reverse shocks that penetrated into the ejecta. The total emission was a combination of the emissions from the reverse- and forward- shocked regions. We discuss how this combined emission accounts for the observed rebrightening.} {We apply numerical models to calculate the light curves from the shocked regions, which include a forward shock originating in the first ejecta and a forward-reverse shock for the second ejecta.} {We find evidence that the central engine became active again 2000 s after the main burst. The combined emission produced by interactions of these two ejecta with the ambient medium can describe the properties of the afterglow of this burst. We argue that the rapid rise in brightness at $\\sim3000$ s in the afterglow is due to the off-axis emission from the second ejecta. The precession of the torus or accretion disk of the central engine is a natural explanation for the departure of the second ejecta from the line of sight.}{} ",
        "introduction": "The gamma-ray burst (GRB) 060206 at Galactic Coordinates $l=78.07$ deg, $b=78.28$ deg triggered $Swift$-\\textbf{BAT} on February 6th, 04:46:53 UT (trigger time $t=0$) (Morris et al. 2006). It exhibited a single peak, with a duration of ${T_{90}=7\\pm 2}$ s and a total fluence of $8.4\\pm0.4\\times10^{-7}$ $\\rm{erg/cm^2}$ in the 15-350 keV band (Palmer et al. 2006). The spectroscopic redshift $z$ is 4.048 (Fynbo et al. 2006). Applying the peak energy ${E_{\\rm{peak}}}={75.4\\pm19.5}$ keV, the best-fit low energy photon index ${\\Gamma_1=1.06\\pm 0.34}$ and a fixed high energy photon index $\\Gamma_2=2.5$, the isotropic-equivalent energy integrated from 1 to $10^4$ keV in the explosion rest frame is ${E_{\\gamma,\\rm{iso}}=5.8\\times10^{52}}$ erg (Palmer et al. 2006).  $Swift$-$\\rm{XRT}$ began to observe this burst 58 s after the \\textbf{BAT} trigger time. At the same time, $Swift$-UVOT started the on-target monitoring and detected the optical afterglow (Boyd et al. 2006). A number of ground-based telescopes performed follow up observations. The 2-m robotic Liverpool Telescope began to observe it at $t=309$ s and carried out multicolor $r'i'z'$ photometry. In the R-band the light-curve exhibited three obvious bumps in the first 75 minutes including a steep rise ($\\Delta r'\\approx -1.6$ at $t\\approx 3000$ s) (Monfardini et al. 2006). About 48.1 minutes later after the trigger time, the Rapid Telescopes for Optical Response (RAPTOR) system at Los Alamos National Laboratory began to take optical images. The obtained light curve confirmed the rebrightening from ${r'\\sim 17.3}$ to a peak value $r'\\sim 16.4$. The subsequent decline to $r'{\\sim 16.75}$ at $t=80$ min was followed by a secondary rebrightening by $\\Delta r'\\sim -0.1$ around $t=90$ min ($\\rm{Wo\\acute{z}niak}$ et al. 2006). The MDM telescope observed a smooth break at $t_b=0.6$ days with another bump at $t\\approx16000$ s. The overall X-ray light curve has a similar shape as the optical light curve (Stanek et al. 2007).  One of the most remarkable features of this burst is that the optical light curve had a significant rebrightening and exhibited small ``bumps'' and ``wiggles''. Similar bumps and wiggles have also been seen in a number of optical afterglows (Stanek et al. 2007). GRB 970508 had an optical afterglow light curve rather similar to that of GRB 060206 (Galama et al. 1998). The optical light curve of another recent burst, GRB 060210, also displayed a rebrightening at time $t\\sim500$ s and a shallow decay in the early epoch. The above ``unusual'' behavior , which is not predicted by the standard fireball model, may be more the norm than the exception (Stanek et al. 2007).  Possible scenarios for the remarkable rebrightening in GRB 060206 at $\\sim$3000 s are a renewed energy injection (Rees \\& \\Mesz 1998; Kumar \\& Piran 2000; Sari \\& \\Mesz 2000) and a density-jump in the circum-burst medium (Dai \\& Lu 2002). However, as discussed by Monfardini et al. (2006), the X-ray band frequency is above the cooling frequency at ${t\\sim 3000}$ s, so the flux does not depend on the ambient density (Freedman \\& Waxman 2001). Nakar \\& Granot (2007) showed that even a sharp and large increase in the ambient medium density cannot produce a significant rebrightening as seen in the afterglow. So the rebrightening cannot be due to a density jump in the ambient medium. If the rebrightening is caused by energy injection, a huge impulsive energy injection ${\\Delta E\\sim 1.8 E_0}$ at $\\sim$3000 s was required, where $E_0$ is the blast wave energy before the rebrightening.  In this paper, we present an alternative scenario. The central engine of this burst become active again after the initial burst and ejects another more collimated jet with a larger viewing angle. This jet and the initial jet sweep up the interstellar medium (ISM). The multi-wavelength emission predicted by this model can reproduce both the observed X-ray and optical data. The observational results are presented in Sect 2. We describe the scenario in Sect 3 and fit the remarkable optical rebrightening of GRB 060206 in Sect 4. Finally, we summarize our results and discuss their implications in Sect 5. ",
        "conclusions": "We have presented a solution for the remarkable rebrightening observed in the afterglow of GRB 060206, which is attributed to emission from an off-axis beamed jet originating from the late activity of the central engine after the prompt gamma-ray emission phase. We have argued that the precession of the torus or accretion disk of the central engine has caused the two jets to move in different directions. Although we only attempt to describe the large rebrightening, it is reasonable to speculate that the five bumps in the R-band afterglow light curve of similar profiles may have the similar origins: emission from different delayed off-axis jets ejected by the GRB central engine. The difference between the temporal indices of the five post-bump light curves indicates that the electrons in different jets have different spectral indices $p$ for their electron distributions. As a result, the spectra exhibit an SED evolution, especially when the separate jets have comparable contributions to the total emission. This can be naturally explained by the significant SED evolution in GRB 060206 detected by the Liverpool Telescope (Monfardini et al. 2006).  The isotropic energy of the second jet is more than one order of magnitude higher than the first jet in our scenario. The collimation-corrected energy of Jet $B$ is $\\sim 4.3$ times larger than Jet $A$, which is larger than the energy required to explain the big rebrightening in the energy injection model as mentioned in Sect 1. The precise mechanism that triggers the central engine again remains unknown. One possibility is that a mass of debris falls back onto the central compact object, generating another more energetic jet. Since our scenario can reproduce well the observations of the large bump in GRB 060206, as a conservative extrapolation, we propose that GRB 970508 and GRB 060210, which display remarkable rebrightening, may have in addition a precessing torus or accretion disk.  A further consequences of our scenario is that the jet break is determined by the off-axis second jet rather than the first one, which produces the main burst. In this case, we cannot measure the isotropic energy of the second jet directly and must fit the afterglow data to obtain an estimation. When a large bump appears in GRB afterglows, we are therefore unlikely to be able to derive the jet opening angle, using the break time and isotropic gamma-ray energy release, because these two quantities originate in two different jets (Stanek et al. 2007).  Finally, several models could be applied to explain afterglow light curves exhibiting rebrightening. These include variable external density profiles (Lazzati et al. 2002), refreshed shocks (Granot et al. 2003; \\Bj et al. 2004), and angular dependence of the energy profile on the jet structure (Nakar et al. 2003), each of which can play a role. Peculiar behavior in light curves caused by the precession of the central engine was discussed by Reynoso et al. (2006), although more observations are required to identify its true nature."
    },
    "0808/0808.2492_arXiv.txt": {
        "abstract": "{ The deceleration mechanisms for relativistic jets in active galactic nuclei remain an open question, and in this paper we propose a model which could explain sudden jet deceleration, invoking density discontinuities. This is particularly motivated by recent indications from HYbrid MOrphology Radio Sources, suggesting that in some case Fanaroff-Riley classification is induced by variations in density of the external medium. } {Exploiting high resolution, numerical simulations, we demonstrate that for both high and low energy jets, always at high Lorentz factor, a transition to a higher density environment can cause a significant fraction of the directed jet energy to be lost on reflection. This can explain how one-sided jet deceleration and a transition to FR I type can occur in HYbrid MOrphology Radio Sources, which start as FR II (and remain so on the other side).} {For that purpose, we implemented in the relativistic hydrodynamic grid-adaptive AMRVAC code, the Synge-type equation of state introduced in the general polytropic case by Meliani et al. (2004). To demonstrate its accuracy, we  set up various test problems in appendix, which we compare to exact solutions that we calculate as well. We use the code to analyse the deceleration of jets in FR II/FR I radio galaxies, following them at high resolution across several hundreds of jet beam radii. } { We present results for 10 model computations, varying the inlet Lorentz factor from 10 to 20, including uniform or decreasing density profiles, and allowing for cylindrical versus conical jet models. As long as the jet propagates through uniform media, we find that the density contrast sets most of the propagation characteristics, fully consistent with previous modeling efforts. When the jet runs into a denser medium, we find a clear distinction in the decelaration of high energy jets depending on the encountered density jump. For fairly high density contrast, the jet becomes destabilised and compressed, decelerates strongly (up to subrelativistic speeds) and can form knots. If the density contrast is too weak, the high energy jets continue with FR II characteristics. The trend is similar for the low energy jet models, which start as underdense jets from the outset, and decelerate by entrainment in the lower region as well. We point out differences that are found between cylindrical and conical jet models, together with dynamical details like the Richtmyer-Meshkov instabilities developing at the original contact interface. } {} ",
        "introduction": "Accretion disks surrounding black holes, jets found in association with compact objects, and Gamma Ray Bursts (GRBs) all represent violent astrophysical phenomena. They are associated with relativistic flows, both with respect to the occuring velocities and to their prevailing equation of state. The jets from Active Galactic Nuclei (AGN) and in GRBs are accelerated in a short distance to reach a high Lorentz factor: typical values are $\\gamma\\sim 5-30$ \\citep{Kellermannetal04} for AGNs and $\\gamma>100$ for GRBs \\citep{Woods&Loeb95, Sari&Piran95}. In this acceleration phase, situated at the base of the jet, it is believed that jet energy is dominated by thermal energy and Poynting flux, and that a fraction of these energies contributes to the efficient acceleration of the flow. Thereby, the relativistic fluid changes its state from relativistic, corresponding to an effective polytropic index $\\Gamma=4/3$, to classical (polytropic index $\\Gamma=5/3$) when the thermal energy is converted to kinetic energy. Also further in the jet paths, where the jets interact with the surrounding medium, a fraction of the directed kinetic energy is converted to thermal energy at the shock fronts formed. According to the prevailing Lorentz factor of the jet, these shocks could be relativistic and therefore very efficient to convert kinetic to thermal energy, or could be Newtonian and only weakly efficient in compression. To investigate the occuring relativistic flows, a realistic equation of state must therefore be adopted, to handle both classical as well as relativistic temperature variations in space and time. For that purpose, we implemented in the relativistic (magneto)hydrodynamic grid-adaptive AMRVAC code~\\citep{Melianietal07, Keppens03, vanderHolst07}, the Synge-type equation of state introduced for a general polytropic case as in Meliani et al. (2004) and as applied in the adiabatic case by Mignone \\& McKinney (2007). To demonstrate its accuracy, we set up various stringent test problems in an appendix, which we compare to exact solutions for Riemann problems, that we calculate as well. The latter include Riemann problems at Lorentz factors of order $\\gamma\\approx 100$, which represent extreme values relevant for Gamma Ray Burst flows.   The development of relativistic numerical hydrodynamic codes can help us understand the physics of astrophysical jet propagation. In the last decade, significant progress was made in numerical special relativistic hydrodynamic (HD) and magnetohydrodynamic codes. Various authors worked on the development of conservative shock-capturing schemes which use either exact or approximate Riemann solver based methods for relativistic hydrodynamics \\citep{Eulderink&Melemma94, Fontetal94, Aloyetal99, DelZanna&Bucciantini02, Mignone&Bodo05} (for a review see \\cite{Marti&Muller03}). The study of relativistic hydrodynamic fluids benefits also from using spatial and temporal adaptive techniques, or Adaptive Mesh Refinement \\citep{Duncan&Hughes94, Zhang&MacFadyen06, Melianietal07, Wangetal07}. The various numerical simulations usually involve a simplified equation of state (EOS) with a constant polytropic index. Notable exceptions with more realistic EOS treatments are found in \\cite{Mignone05, Mignone&McKinney07, Perucho&Marti07}. We present in the appendix to this paper the required extension of the AMRVAC code \\citep{Melianietal07, Keppens03, vanderHolst07} with the more realistic polytropic EOS introduced by~\\cite{Melianietal04} which is based on the Synge gas equation, the relativistically correct perfect gas law~\\citep{Synge57, Mathews71}. The equations which we solve, and the schemes used, are also mentioned there.  Many previous investigations of relativistic jet propagation through the interstellar medium (ISM) in radio galaxies concentrate on uniform ISM conditions \\citep{Duncan&Hughes94, Martietal97, KomissarovetFall98, Aloyetal99}. These investigations contributed to our understanding of the jet deceleration, where one then distinguishes dynamics for Faranoff-Riley type I (FR I) and FR II radio galaxies, according to the power of the jet and hence the accretion rate in their galactic center. In the FR I radio galaxies, the associated jets are relativistic on parsec scale and sub-relativistic on kiloparsec-scales, so that jet deceleration, and thus energy redistributions, must happen on kiloparsec scales \\citep{Hardcastleet05}. Various studies have looked into possible deceleration processes with both non-relativistic and relativistic HD codes. \\cite{Hooda&wiita96} and \\cite{Hooda&wiita98} investigated with a classical HD code the propagation of a 3D jet through an interface between dense ISM and less dense intercluster medium. They found that for such interface marking a density decrease, the jet does not decelerate when crossing it, and that the deceleration mostly happens in the lower region (dense ISM). Moreover, in their study the jets are not deflected at the interface ISM/ICM. \\cite{Normanetal88} also use a classical code to study the sudden deceleration of jets, when the jet crosses a shock wave in the external medium. Other works investigate the interaction of jets with dense clouds in both non-relativistic hydrodynamic simulations \\citep{Saxtonetal05} and relativistic simulations \\citep{Choietal07}. \\citet{Duncan&Hughes96} study relativistic jet propagation in uniform overdense media and use an equation of state for a pair-plasma. \\citet{Schecketal02} study the influence of the matter composition of a high jet energy jet on its interaction with the external medium. They investigate the two extreme cases of pure leptonic and baryonic plasmas. \\citet{Perucho&Marti07} examined the propagation of the low energy jet of the specific FR I radio source 3C 31 using an elaborate hydrodynamics model, with an equation of state that distinguishes the contribution of leptons and baryons to density and pressure. Most recently, \\citet{Rossietal08} investigate the propagation of jets in uniform overdense media in 3D, using a realistic synge EOS.  \\begin{table*} \\caption{The most relevant characteristics and parameters for various selected relativistic hydrodynamic simulations in the literature: DH94 \\citep{Duncan&Hughes94}, M97 \\citep{Martietal97}, H98 \\citep{Hardeeetal98}, R99~\\citep{Rosenetal99} S02~\\citep{Schecketal02},  PM07 \\citep{Perucho&Marti07}, R08 \\citep{Rossietal08}. Our simulations are indicated by MKG. We mention the type of the simulation 2D or 3D and use of AMR, jet beam kinetic luminosity $L_{\\rm b}$ (when available), $\\gamma_{\\rm b}$ Lorentz factor of the beam, $M_{\\rm b}$ relativistic Mach number, $\\eta$ density ratio (``d\" means density variation), $\\theta$ open angle of the jet, $R_{\\rm b}/r_{\\rm cell}$ grid cells through jet beam, type of EOS (\"Lep\" means Leptonic, \"Bar\" Baryonic), endtime of the simulation in units of light crossing time of the jet beam radius.} \\label{table:1}      % \\centering                                      % \\begin{tabular}{c c c c c c c c c c} \\hline\\hline                        % Paper & case & $L_{\\rm b}$ & $\\gamma_{\\rm b}$ & $M_{\\rm b}$ & $\\rho_{\\rm medium}/\\rho_{\\rm jet}$& $\\theta_{\\rm b} $& $R_{\\rm b}/r_{\\rm cell} $ &EOS& Size in $R_{\\rm b}$ (and time)\\\\    % \\hline                                   % \\hline                                   % DH94/R99  & A (2D, AMR)  & - &1.048 & 6 & $10.0$&0& 24& $5/3$&  $10 \\times 41.67 $ (-)\\\\ DH94/R99  & B (2D, AMR)  & - &5.0 & 8 & $10.0$&0& 24& $5/3$&  $16.67 \\times  41.67$ (-)\\\\ DH94/R99  & C (2D, AMR)  & - &10.0 & 8 & $10.0$ &0& 24& $5/3$& $16.67\\times 41.67 $ (-)\\\\ DH94/R99  & D (2D, AMR)  & - &10.0 & 15 & $10.0$ &0& 24& $4/3$&  $16.67 \\times 41.67$ (-)\\\\ M97  & A1 (2D)  & - &7.1 & 9.97 & $10^2$ &0& 20& $4/3$&  $10.5 \\times 50$ (50)\\\\ M97  & A2 (2D)  & - &22.37 & 31.86 & $10^2$ &0& 20& $4/3$&  $10.5 \\times 50$ (50)\\\\ M97  & a1 (2D)  & - &7.1 & 9.97 & $1.0$ &0& 20& $4/3$& $10.5 \\times 25$ (25)\\\\ M97  & a2 (2D)  & - &7.1 & 9.97 & $10.0$ &0& 20& $4/3$& $10.5 \\times 25$ (25)\\\\ M97  & B1 (2D)  & - &7.1 & 41.95 & $10^2$ &0& 20& $4/3$&   $10.5 \\times 50$ (50)\\\\ M97  & B2 (2D)  & - &22.37 & 132.32 & $10^2$ &0& 20& $4/3$&   $10.5 \\times 50$ (50)\\\\ M97  & b1 (2D)  & - &2.29 & 13.61 & $10.0$ &0& 20& $4/3$&  $10.5 \\times 25$ (25)\\\\ M97  & b2 (2D)  & - &7.1 & 41.95 & $10.0$ &0& 20& $4/3$& $10.5 \\times 25$ (25)\\\\ M97  & C1 (2D)  & - &2.29 & 13.61 & $10^2$ &0& 20& $5/3$&  $10.5 \\times 50$ (50)\\\\ M97  & C2 (2D)  & - &7.1 & 41.95 & $10^2$ &0& 20& $5/3$& $10.5 \\times 50$ (50)\\\\ M97  & C3 (2D)  & - &22.37 & 132.32 & $10^2$ &0& 20& $5/3$&   $10.5 \\times 50$ (50)\\\\ M97  & c1 (2D)  & - &2.29 & 13.61 & $10.0$ &0& 20& $5/3$&  $10.5 \\times 25$ (25)\\\\ M97  & c2 (2D)  & - &7.1 & 41.95 & $10.0$ &0& 20& $5/3$&  $10.5 \\times 25$ (25)\\\\ H98 & B (2D, AMR)   & - &5.52 & 8.87 & $10.0$ &0& 24& $5/3$& $16.67 \\times  41.67$ (140.0)\\\\ H98 & C (2D, AMR)   & - &14.35 & 11.52 & $10.0$ &0& 24& $5/3$& $16.67 \\times  41.67$ (21.43)\\\\ H98 & D (2D, AMR)   & - &10.0 & 15.0 & $10.0$ &0& 24& $4/3$&  $16.67 \\times  41.67$ (35.273)\\\\ H98 & E (2D, AMR)   & - &2.55 & 8.53 & $10.0$ &0& 24& $5/3$& $16.67 \\times  41.67$ (140.0)\\\\ S02 & A (2D)  & $10^{46}$ &6.62 & 9.24 & $10^5$ &0& 6& Lep& $200\\times 500$ (5950)\\\\ S02 & B (2D)  & $10^{46}$ &6.62 & 14.33 & $10^3$ &0& 6& Lep& $200\\times 500$ (5400)\\\\ S02 & C (2D)  & $10^{46}$ &7.95 & 130.14 & $10^3$ &0& 6& Bar& $200\\times 500$ (5300)\\\\ PM07 & A (2D)  & $10^{44}$ &2.0 & 4.75 & $10^5$ ``d\" &0& 16& Lep/Bar& $200\\times 450$ (37231)\\\\ R08& A (3D)  & - &10.0 & 28.3 & $10^2$ &0& 20& Synge&  $50 \\times 150 \\times 50$ (240)\\\\ R08& B (3D)   & - &10.0 & 28.3 & $10^4$ &0& 20& Synge& $60 \\times 75 \\times 60$  (760)\\\\ R08& C (3D)   & - &10.0 & 28.3 & $10^4$ &0& 12& Synge&  $50 \\times 75\\times 50$ (760)\\\\ R08& D (3D)   & - &10.0 & 300 & $10^4$ &0& 20& Synge&  $50 \\times 150\\times 50$ (760)\\\\ R08& E (3D)   & - &10.0 & 300 & $10^2$ &0& 12& Synge&  $24 \\times 200\\times 24$ (150)\\\\ MKG & A (2D, AMR)   & $10^{46}$ &20.0 & 1200 & $0.1496$-$4.687$  &0& 64& Synge& $10 \\times 400$ (380)\\\\ MKG & B (2D, AMR)  &  $10^{46}$ &20.0 & 1200 & $0.1496$-$671.22 $ &0& 128& Synge& $10 \\times 400$ (900)\\\\ MKG & C (2D, AMR)   & $10^{43}$ &10.0 & 39 & $36.52$-$203.66$  ``d\" &0& 76& Synge& $10 \\times 400$ (820)\\\\ MKG & D (2D, AMR)   & $10^{43}$ &10.0 & 39 & $36.52$-$148.15$  ``d\"  &1& 76& Synge& $10 \\times 400$ (820)\\\\ MKG & E (2D, AMR)   & $10^{43}$ &20.0 & 35 & $146$-$847.46$ ``d\" &0& 76& Synge& $40 \\times 400$ (820)\\\\ MKG & F (2D, AMR)   & $10^{43}$ &20.0 & 35 & $146$-$645.16$   ``d\" &1& 76& Synge& $40 \\times 400$ (820)\\\\ MKG & G (2D, AMR)   & $10^{46}$ &10.0 & 1300 & $0.033$-$0.495$   ``d\" &0& 144& Synge& $40 \\times 400$ (380)\\\\ MKG & H (2D, AMR)   & $10^{46}$ &10.0 & 1300 & $0.033$-$30.6748$  ``d\"  &1& 76& Synge& $40 \\times 400$ (800)\\\\ MKG & I (2D, AMR)   & $10^{46}$ &10.0 & 1300 & $0.033$-$0.306748$  ``d\"  &1& 76& Synge& $40 \\times 400$ (480)\\\\ MKG & J (2D, AMR)   & $10^{46}$ &20.0 & 1200 & $0.1496$-$12.95$  ``d\" &1& 76& Synge& $40 \\times 400$ (480)\\\\ \\hline                                             % \\end{tabular} \\end{table*} In this paper, we investigate a scenario of relativistic jet propagation through an ISM with a sudden jump in density. Since jets propagate over enormous distances, it is inevitable that they encounter density jumps in the external medium. As we assume that matter in the inner part of the radio galaxy has been cleared away during the evolution of the radio galaxy, we typically follow jet dynamics through a lighter medium at first, which then suddenly changes far away to a denser medium. In a sequence of model runs, we vary the jet and medium properties to cover both high and low jet beam kinetic luminosities, straight as well as conical jets (varying the jet opening angle), and allow for either uniform or decreasing density profiles. To put our work into perspective and to appreciate its relation to previous works better, we compiled in Table~\\ref{table:1} the most relevant parameters and simulation specifics for selected relativistic hydro models. The most novel assumption is the transition from low to high density across a contact interface (across which the pressure is necessarily constant), which is different from the work done by \\citet{Hooda&wiita98} where they use a classical code with a density decrease across the contact, and also different from the study by \\citet{Lokenetal93}, where they use a classical code to look into jet propagation through a pressure wall. We explore the influence of a sudden density jump in ISM on jet propagation, stability, and formation of the bow shock. We aim to better understand the jet efficiency in transporting energy and mass from the central AGN regions to the denser ISM at large distances, where the jet cocoon gets formed. Note also from Table~\\ref{table:1} that our models extend the previous works most notably by typically covering much larger distances than previously studied (as we need to obtain representative endstate morphologies in both lower and upper media), and are invariably in the high Lorentz factor regime, combined with a high resolution through the jet beam. The high energy jets we consider start out as denser than their immediate surroundings (a likely property deduced from all common jet launch scenarios, but previously ignored by all jet simulations to date). Moreover, the density contrasts we investigate between jet and outer medium are much more reasonable than the 5 orders of magnitude density differences previously used by~\\citet{Schecketal02}. In what follows, we first motivate our model assumptions and the simulation setup in Section~\\ref{secmodel}. We discuss our main findings in Section~\\ref{secres}. ",
        "conclusions": "In this paper, we extended, presented and applied the relativistic hydrodynamics AMRVAC code \\citep{Melianietal07,vanderHolst07} with a generalized variable polytropic index equation of state for the purpose of modeling the relativistic shocks in GRBs and AGN jets. As follows in the appendix, we used various shock-capturing schemes on a variety of test cases for adiabatic and non-adiabatic cases for code validation. We demonstrated the code ability with stringent new test problems, developed according to the astrophysical context. We tested the code also for a case with a heated flow, using the Synge-like equation of state from \\cite{Melianietal04}. The Riemann problems were also solved exactly, and the AMR simulations handled cases with Lorentz factors of order 100 accurately.  We explored a new scenario for sudden deceleration of relativistic jets in the FR I radio galaxies. This model for jet deceleration is based on a density jump in the external medium, with density suddenly increasing to the upper medium. We include models of conical jets with an initial opening angle, and allow for decreasing denisty profiles within the upper medium. We investigated the propagation and the dynamics of relativistic AGN jets over a long distance, resolving small scale instabilities which develop in the cocoon and which could be responsible for particle acceleration. This study was only possible with adaptive mesh refinement. We quantified and discussed the deceleration of the jet resulting from its interaction with a layered medium. We point out that jet deceleration can be significantly aided by an internal oblique shock that forms in jets with conical expansion. Under fairly extreme (but not as extreme as pursued in other simulations to date) density conditions in the dense upper region, it can reproduce the strong deceleration observed in FR I jets.  In the early phase, the head of the cold fast (beam Lorentz factor 10 or 20) jet propagates in the low density medium at a high Lorentz factor $\\gamma=5$. In cases with high energy, almost no cocoon or backflows develop in this phase, and as the jet is heavier than the external region, it behaves more ballistically. When observed in this region, the relativistic beaming is high, so that one would observe a one-sided jet. In low energy jets, a cocoon always forms and slow backflows develop. They in turn induce multiple internal shocks making the jet decelerate and re-accelerate. In jets with low energy and with an initial conical flow, the flow remains ballistic until it reaches a self-consistently forming oblique shock. This compresses and decelerates the jet and two regions appear in the jet beam, the first with high relativistic beaming, and the second with low Lorentz factor, showing clear disturbances by the surrounding cocoon.  In the outer region where the density increases suddenly, we find increased entrainment of ambient material through velocity shear instabilities at the jet boundary and working surfaces. Note that we here for the first time address how the prior interaction of the jet with the low density ambient medium in the inner region (which makes the head of the jet consist of swept-up ambient medium and a shocked beam) affects jet propagation in denser media. The pre-structured jet head has lower Mach number, which gives rise to a strong interaction with the denser ambient medium.  For models with a weak increase in density in the external medium, the cocoon shows little turbulent structure and no backflow appears. The jet in this case remains stable and relativistic in the denser upper region. Also a fraction of its energy is deposited in the external medium, but only through the frontal shock. This model could be relevant for jets in FR II. With this scenario, one can explain how most of the energy of the jet is deposited in the outer region, since relatively little jet energy gets transferred to the medium in the inner region.  Those models with more extreme increases in density in the external medium, showed the formation of an overpressured and turbulent cocoon. A strong backflow develops in the outer region, which disturbs the jet structure. The backflow compresses the jet and induces internal shocks which give rise to knot formation. The jet  becomes subrelativistic within the simulated domain. The models with initial conical jets show the development of an oblique shock that decelerates the jet. In future work, we will follow these jets over increasingly larger distances in 3D scenarios. We also intend to provide synthetic observational radio maps from our computed jet models.   \\appendix\\label{appendixA}"
    },
    "0808/0808.2812_arXiv.txt": {
        "abstract": "We report the discovery by the Robotic Optical Transient Search Experiment (ROTSE-IIIb) telescope of SN 2008es, an overluminous supernova (SN) at $z=0.205$ with a peak visual magnitude of $-$22.2.  We present multiwavelength follow-up observations with the \\textsl{Swift} satellite and several ground-based optical telescopes. The ROTSE-IIIb observations constrain the time of explosion to be 23$\\pm$1 rest-frame days before maximum.  The linear decay of the optical light curve, and the combination of a symmetric, broad H$\\alpha$ emission line profile with broad P Cygni H$\\beta$ and Na I~$\\lambda5892$ profiles, are properties reminiscent of the bright Type II-L SNe 1979C and 1980K, although SN 2008es is greater than 10 times more luminous.  The host galaxy is undetected in pre-supernova Sloan Digital Sky Survey images, and similar to Type II-L SN 2005ap (the most luminous SN ever observed), the host is most likely a dwarf galaxy with $M_{r} > -17$.  \\textsl{Swift} Ultraviolet/Optical Telescope observations in combination with Palomar 60 inch photometry measure the spectral energy distribution of the SN from 200 to 800 nm to be a blackbody that cools from 14,000 K at the time of the optical peak to 6400 K 65 days later.  The inferred blackbody radius is in good agreement with the radius expected for the expansion speed measured from the broad lines (10,000 km s$^{-1}$). The bolometric luminosity at the optical peak is $2.8 \\times 10^{44}$ erg s$^{-1}$, with a total energy radiated over the next 65 days of $5.6 \\times 10^{50}$ erg.  The exceptional luminosity of SN 2008es requires an efficient conversion of kinetic energy produced from the core-collapse explosion into radiation.   We favor a model in which the large peak luminosity is a consequence of the core-collapse of a progenitor star with a low-mass extended hydrogen envelope and a stellar wind with a density close to the upper limit on the mass-loss rate measured from the lack of an X-ray detection by the \\textsl{Swift} X-Ray Telescope. ",
        "introduction": "} Hydrogen-rich supernovae (SNe II) produce their radiative energy in several phases; with a diversity of luminosities, light curves, and spectroscopic features that divide them into subclasses.  An initial burst of UV/X-ray radiation is observed at the time of shock breakout (Schawinski \\etal 2008; Gezari \\etal 2008a), followed by declining UV/optical emission from the adiabatic expansion and cooling of the SN ejecta (e.g., SN II-pec 1987A, Hamuy \\etal 1988; SN IIb 1993J, Schmidt \\etal 1993; Richmond \\etal 1994, SN II-P 2006bp, Immler \\etal 2007).  Type II-plateau (II-P) SNe are characterized by their subsequent plateau in optical brightness, which is understood as the result of a progenitor with a substantial hydrogen envelope, for which a cooling wave of hydrogen recombination recedes through the inner layers of the ejecta (Falk \\& Arnett 1977; Litvinova \\& Nadyozhin 1983).  Type II-linear (II-L) SNe do not have a plateau phase, but rather a steep drop in luminosity that is thought to be the result of the ejection of large amounts of radioactive material (Young \\& Branch 1989) or a small hydrogen envelope mass (Barbon \\etal 1979; Blinnikov \\& Bartunov 1993).  The subclass of bright SNe II-L ($M_{B} < -18$; Patat \\etal 1994) has been proposed to be the result of an extended low-mass envelope (Swartz \\etal 1991) and reprocessing of UV photons in the superwind of the progenitor star (Blinnikov \\& Bartunov 1993) or a gamma-ray burst (GRB) explosion buried in a hydrogen-rich envelope (Young \\etal 2005).  Type II-narrow (IIn) SNe show signs of interaction of the ejecta with  circumstellar material (CSM) in the form of narrow emission lines, X-ray emission, and an excess in optical luminosity (e.g., SN 1988Z: Turatto \\etal 1993).  After the photospheric phase in SNe II, a final exponential decline is powered by heating due to the radioactive decay of $^{56}$Co.  In this paper, we present the discovery of an overluminous SN II that has the properties of an SN II-L, but with an exceptional luminosity that we argue is best explained by the core-collapse of a progenitor star with a low-mass, extended hydrogen envelope and a dense stellar wind.  We disfavor more exotic scenarios such as an interaction with a massive shell of CSM expelled in an episodic eruption, a pair-instability explosion, or a buried GRB. In \\S \\ref{sect_obs}, we present the Robotic Optical Transient Experiment (ROTSE-IIIb) discovery data, and follow-up observations with the \\textsl{Swift} satellite and several ground-based telescopes; in \\S \\ref{sec_disc} we compare the properties of SN 2008es to other known Type II SNe, and use our observations to constrain the mechanism powering the extraordinary radiative output; and in \\S \\ref{sec_conc} we make our conclusions. ",
        "conclusions": "\\label{sec_conc}  The overluminous SN 2008es ($M_{V}=-22.2)$ is the second most luminous SN observed, and in the same league as the extreme SNe 2005ap, 2006gy, and 2006tf, which had peak visual magnitudes of $-22.7$, $-22.0$, and $-20.7$, respectively.  In Figure \\ref{fig_sne}, we compare the absolute $R$-band light curves of the 4 SNe (all discovered by the ROTSE-IIIb telescope) in rest-frame days since the time of explosion.  SN 2006gy and 2006tf are classified spectroscopically as Type IIn SNe because of their narrow peaked P Cygni Balmer lines (Ofek \\etal 2007; Smith \\etal 2008), although with deviations from the smooth profiles typical of SNe IIn in SN 2006gy (Smith \\etal 2007).  The time of explosion for SN 2005ap is estimated by Quimby \\etal (2007a) to be $\\sim 3$ weeks before maximum.  The time of explosion for SN 2006tf is poorly constrained (Smith \\etal 2008), so for comparison we assume that it takes the same time to reach the peak as SN 2008es.  Of the SNe in the same luminosity class, the light curve of SN 2008es is most like SN 2005ap, the most luminous SN ever observed \\citep{2007ApJ...668L..99Q}. Interestingly, SN 2005ap also has a dwarf galaxy host, with $M_{R} = -16.8$.    However, SN 2005ap does differ from SN 2008es in the fact that it demonstrated absorption features at early times, as well as P Cygni absorption in its broad H$\\alpha$ profile.  This is not consistent with the spectra of bright II-L SNe which have decreasing P Cygni absorption with increasing luminosity (Patat \\etal 1994).  The extended envelope/wind model is problematic for SN 2005ap, since it would be hard to produce early absorption features without corresponding emission. The extreme luminosity and linear decay of SN 2005ap were attributed instead to a collision of the SN ejecta with a dense circumstellar shell, a GRB explosion buried in a H envelope, or a pair-instability eruption powered by radioactive decay \\citep{2007ApJ...668L..99Q}.  The peak luminosity, linear decay, and spectroscopic features of SN 2008es place it in a subclass of ultra-bright Type II-L SNe. The light curve and SED of SN 2008es, measured in detail from the UV through the optical, point toward a core-collapse explosion of a nonstandard progenitor star with a super wind and extended envelope.  Wide-field optical synoptic surveys such as the ROTSE-III Supernova Verification Project, Pan-STARRS, and LSST will continue to explore the large parameter space of core-collapse explosions, and increase our understanding of the effects of variations in progenitor star properties and environments.  {\\bf Note Added in Proof:} More recent $V$-band observations from \\textsl{Swift}, as well as $V, R, I$ measurements from the MDM 2.4m, reveal that the slope of the light curve is steeper than the decay rate of $^{56}$Co up to $\\sim$145 rest-frame days after maximum.  Therefore, radioactive decay must not yet be the dominant source of the luminosity from the SN, and the inferred upper limit on the initial $^{56}$Ni mass is securely within the range of normal Type II SNe (Hamuy 2003)."
    },
    "0808/0808.0483_arXiv.txt": {
        "abstract": "We report on a problem found in {\\sc Mercury}, a hybrid symplectic integrator used for dynamical problems in Astronomy. The variable that keeps track of bodies' statuses is uninitialised, which can result in bodies disappearing from simulations in a non-physical manner. Some {\\sc fortran} compilers implicitly initialise variables, preventing simulations from having this problem. With others compilers, simulations with a suitably large maximum number of bodies parameter value are also unaffected. Otherwise, the problem manifests at the first event after the integrator is started, whether from scratch or continuing a previously stopped simulation. Although the problem does not manifest in some conditions, explicitly initialising the variable solves the problem in a permanent and unconditional manner. ",
        "introduction": "\\label{intro} {\\sc Mercury} \\citep{Chambers1999} is a general-purpose software package, written in {\\sc fortran77}, for doing N-body integrations to investigate dynamical problems in Astronomy. In a set of simulations to study accretion dynamics using {\\sc Mercury} \\citep{Torres2008}, it was found that the results contained fewer planets than expected. Analysis of output files and creation of scripts and movies to carefully follow the processes showed discontinuities in the number of embryos during the integrations. No events for these bodies were being registered in the output files and their disappearance was non-physical. ",
        "conclusions": "\\label{conclusions}  We repeated all the tests showed in Section \\ref{charac} with a corrected version of {\\sc Mercury}, initialising the variable array {\\sc stat} at the end of the subroutine MIO\\_IN in the source code. No discontinuities were seen in any results, independently of conditions. The new version is now being used in 40 new simulations of accretion dynamics. We believe this small change can improve the program making it more reliable for any type of N-body problem simulations. The corrected version can be downloaded from \\url{http://www.astro.keele.ac.uk/~dra/mercury/}."
    },
    "0808/0808.3449_arXiv.txt": {
        "abstract": "We test the proposal by El-Zant \\etal\\ that the dark matter density of halos could be reduced through dynamical friction acting on heavy baryonic clumps in the early stages of galaxy formation.  Using $N$-body simulations, we confirm that the inner halo density cusp is flattened to $0.2$ of the halo break radius by the settling of a single clump of mass $\\ga 0.5\\%$ of the halo mass.  We also find that an ensemble of 50 clumps each having masses $\\ga 0.2\\%$ can flatten the cusp to almost the halo break radius on a time scale of $\\sim9\\;$Gyr, for an NFW halo of concentration 15.  We summarize some of the difficulties that need to be overcome if this mechanism is to resolve the apparent conflict between the observed inner densities of galaxy halos and the predictions of LCDM. ",
        "introduction": "\\label{intro} The LCDM model for structure formation in the universe has been very successful on large scales (\\eg\\ Springel, Frenk \\& White 2006), but has run into a number of difficulties on galaxy scales.  One difficulty is that the predicted dark matter density in the halos of bright galaxies appears to be greater than that inferred from the best observational data (\\eg\\ Weiner \\etal\\ 2001; Dutton \\etal\\ 2007) or from dynamical friction constraints (Debattista \\& Sellwood 2000); see Sellwood (2008b) for a review.  A number of possible solutions have been suggested.  Binney, Gerhard \\& Silk (2001) and others suggest that the halo density can be reduced by feed-back from star formation, but Gnedin \\& Zhao (2002) show that the density cannot be reduced by more than a factor two, even for the most extreme feedback, unless the baryons are unreasonably concentrated.  Weinberg \\& Katz (2002) propose that dynamical friction between a bar and the halo can reduce the central density, but Sellwood (2008a) shows that significant reductions require extreme bars and remove a large fraction of the angular momentum from the baryons.  A third possible solution was proposed by El-Zant, Shlosman \\& Hoffman (2001, hereafter EZ01), who suggested that dynamical friction between heavy gas clumps and dark matter would transfer energy to the dark matter, thereby reducing its inner density.  Essentially EZ01 propose a process of mass segregation that allows the baryonic matter to displace the dark matter.  If it works, it would amount to baryon settling with halo {\\it de}-compression, as Dutton \\etal\\ (2007) argued would be required to improve agreement of LCDM predictions with scaling relations of $\\sim L_*$ galaxies.  Kassin, de Jong \\& Weiner (2006) also remark that rotation curves could be more easily reconciled with LCDM predictions if halo compression could somehow be avoided.  Tonini, Lapi \\& Salucci (2006) use the idea to predict rotation curves for galaxies.  Mo \\& Mao (2004) apply the mechanism to pre-process small halos in the very early stages of structure formation.  A similar, but significantly distinct, idea was tested by Gao \\etal\\ (2004), who argue that some unspecified attractor mechanism causes violent relaxation processes to reset the collisionless mass profile of the combined dark and baryonic matter.  Since violent mergers disrupt disks, Gao \\etal\\ clearly invoke a different process from the slow frictional in-spiral of baryonic clumps proposed by EZ01.  EZ01 suppose that the baryons collect into dense mass clumps which then sink to the center by dynamical friction.  They assume: (1) the clumps are small enough they do not collide and (2) are sufficiently tightly bound that they are not tidally disrupted.  They also implicitly assume that the dense clumps (3) maintain their coherence independent of any internal physical processes (such as possible star formation) and (4) they require the mass clumps to contain little dark matter.  The in-spiral of massive clumps has been studied extensively and we do not attempt a complete list of references.  Van den Bosch \\etal\\ (1999) estimate the infall rate of dark matter sub-clumps through dynamical friction.  Ma \\& Boylan-Kolchin (2004) simulate a mass spectrum of satellites each composed of particles moving within a main halo; they find that the combined density profile of the inner halo and disrupted satellites can either steepen or flatten, depending on the properties of the clumps.  Arena \\& Bertin (2007) study both the strength of the friction force and the effect on the density profile for both cuspy and cored initial galaxy profiles.  Here we undertake a direct test of the proposal by EZ01, using self-consistent $N$-body simulations, in order to determine the clump mass required to effect a significant density reduction.  We first derive the settling time-scale for single massive clumps in cuspy halos, and then present simulations containing an ensemble of heavy clumps in a background of light halo particles.  We discuss the physical plausibility of their assumptions in the concluding section.  El-Zant \\etal\\ (2004, hereafter EZ04) apply the same idea to account for the flattened density profile reported by Sand \\etal\\ (2002) in the galaxy cluster MS 2137-23.  In this context, they identify the galaxies themselves as the heavy baryonic clumps that settle to the cluster center.  Both they, and Nipoti \\etal\\ (2004), present simulations to show that the in-spiral of galaxies can cause the background dark matter density profile to flatten, although both studies neglect the DM halos attached to the infalling galaxies.  In an appendix, we show that our simulations are in good agreement with those of EZ04 for the same problem. ",
        "conclusions": "El-Zant \\etal\\ (2001) proposed that dynamicl friction on massive gas clumps could lead to the outward displacement of dark matter. They estimated the magnitude of the effect on a galaxy halo using the Chandrasekhar dynamical friction formula.  They followed up (El-Zant \\etal\\ 2004) with simulations of a galaxy cluster, in which the heavy particles were the individual galaxies, finding some density decrease.  We present simulations of softened heavy particles moving under the influence of gravitational forces within a background NFW halo composed of a large number of light particles.  We confirm that the inner halo density is lowered by dynamical friction on heavy particles, as has been found previously (\\eg\\ Ma \\& Boylan-Kolchin 2004; Arena \\& Bertin 2007), and show that larger reductions result from more massive particles that also settle more rapidly.  We specifically test the model proposed by EZ01, who divided all the baryons into a number of equal mass clumps that were already somewhat concentrated towards the halo center.  We find (Fig.~\\ref{collect}) that the $\\sim 10\\%$ baryonic mass fraction must be entirely made up of $\\la 150$ equal clumps if the cusp in the dark matter is to be flattened to $r \\la r_s$ within 300 dynamical times, which scales to $\\sim 9\\;$Gyr for a $c=15$ halo.  The principal reason that we observe a more mild density reduction than predicted by EZ01 is that friction is significantly weaker than they assumed.  They, and Tonini \\etal\\ (2006), adopted $\\ln\\Lambda \\simeq \\ln\\eta \\simeq 8.5$ whereas we find $\\ln\\Lambda \\simeq 2.5$ (\\S\\ref{tscale}), which causes the mass clumps to settle 2-3 times more slowly than they assumed, with a corresponding reduction in the rate energy is given up to the halo.  More massive clumps both spiral in more rapidly and displace more dark matter (Fig.~\\ref{massch}).  Ma \\& Boylan-Kolchin (2004) presented a pair of simulations that revealed a much smaller density change when the three most massive clumps were eliminated compared with the result when they were retained.  Thus substantial reductions of the inner halo density require just a few extremely massive clumps, assuming they are not disrupted as they settle.  The time scale, in years, varies as $\\rho_s^{-1/2}$, and therefore would be longer in lower concentration halos.  At higher redshifts, the orbital period is a fixed fraction of the age of the universe at a constant overdensity (since both scale as the square root of the density in an Einstein-de Sitter universe).  While halo densities rise with redshift, their relative over-densities are lower (\\eg\\ Zhao \\etal\\ 2003), and therefore more massive clumps are required to have any effect in the time available; Mo \\& Mao (2004) invoke baryon clumps having 5\\% of the halo mass when $2 \\la z \\la 5$.  EZ04 argue that the same physical process applies to galaxies in clusters, and present supporting simulations.  Again adopting their numerical model and assumptions, we confirm (Appendix) their result that the density cusp flattens for $r \\la 0.2r_s$, where we also demonstrate that the result is independent of the numerical method.  The physical assumptions underlying these successes are more questionable, however.  Purely baryonic gas clumps are indeed expected to form through the Jeans instability as gas cools and settles in the main halo (\\eg\\ Maller \\& Bullock 2004), but the high-resolution experiments of Kaufmann \\etal\\ (2006) show that gas fragments formed in this way have masses ranging up to $\\sim 10^6\\;$M$_\\odot$ only, some two orders of magnitude smaller than required to have interestingly short settling times.  Dark matter sub-halos (\\eg\\ Diemand, Kuhlen \\& Madau 2007) may contain gas and would spiral in more rapidly.  However, settling of such sub-halos would bring both baryons and DM into the center, diluting the desired separation of baryons from dark matter.  The resulting changes to the overall dark matter density profile will clearly be less substantial.  When Ma \\& Boylan-Kolchin (2004) treat sub-clumps as collections of particles, they find that the stripped mass is added to the main DM halo in such a way as approximately to replace the mass scattered to larger radii by dynamical friction.  (The net effect can be either a small increase or decrease in the inner dark matter density, depending on the clump mass spectrum.)  If dynamical friction is to separate the baryons from dark matter with the desired efficiency, baryonic mass clumps in sub-halos must somehow be stripped of their dark matter with high efficiency, without dissolving the gas clumps themselves.  In the case of the cluster simulation, EZ04 describe the heavy particles as purely baryonic galaxies.  Here again, the change in the central dark matter density will not be as substantial when the dark halos of the galaxies are taken into account, as also noted by Nipoti \\etal\\ (2004).  Additional density reduction would be possible if massive baryonic clumps in the center of the halo could be re-accelerated.  Models of this kind have been proposed by Mashchenko, Couchman \\& Wadsley (2006, 2007), who invoke star formation activity, and by Peirani, Kay \\& Silk (2008), who use AGN jets.  It should be noted, however, that massive clumps of dense gas are not easily accelerated by these astrophysical processes (\\eg\\ MacLow \\& Ferrara 1999).  The rapid density reduction reported by Mashchenko \\etal\\ (2006) occurs because they employed a gas clump of $10^8\\;{\\rm M}_\\odot$ within the cusp driven sinusoidally with a bulk speed that peaked at 18~km~s$^{-1}$.  Most halo density reduction occurs as clumps plunge rapidly within the cusp (Fig.~\\ref{massfr}), as noted by EZ01.  Therefore, clumps that start out in the cusp fall in more rapidly and are most efficient at displacing the dark matter.  However, only a small fraction ($\\la 5\\%$) of the baryons start out in the cusp, if they are distributed as the dark matter.  The desired halo density reduction would be more rapid if the baryonic fraction in the cusp could be increased, but it is hard to see how such an initial mass segregation could have arisen without compressing the halo.  Thus the challenge, if the proposal by EZ01 is to help solve the central density problem of dark matter halos, is to find ways to collect gas into clumps exceeding $\\sim 0.5\\%$ of the halo mass that can settle as coherent entities into the center."
    },
    "0808/0808.3480_arXiv.txt": {
        "abstract": "The dynamics of small global perturbations in the form of linear combination of a finite number of non-axisymmetric eigenmodes is studied in two-dimensional approximation. The background flow is assumed to be an axisymmetric perfect fluid with the adiabatic index $\\gamma=5/3$ rotating with power law angular velocity distribution $\\Omega \\propto r^{-q}$, $1.5<q<2.0$, confined by free boundaries in the radial direction. The substantial transient growth of acoustic energy of optimized perturbations is discovered. An optimal energy growth $G$ is calculated numerically for a variety of parameters. Its value depends essentially on the perturbation azimuthal wavenumber $m$ and increases for higher values of $m$. The closer the rotation profile to the Keplerian law, the larger growth factors can be obtained but over a longer time. The highest acoustic energy increase found numerically is of order $\\sim 10^2$ over $\\sim 6$ typical Keplerian periods. Slow neutral eigenmodes with corotation radius beyond the outer boundary mostly contribute to the transient growth. The revealed linear temporal behaviour of perturbations may play an important role in angular momentum transfer in toroidal flows near compact relativistic objects. ",
        "introduction": "Hydrodynamical equations allow the existence of laminar axisymmetric toroidal flows with free boundaries in the vicinity of a gravitating body. The fundamental question is whether the accretion is possible if small perturbations of some sort are present in the medium, i.e whether the perturbations will grow and affect the background to allow the angular momentum transfer. Astrophysically, it is interesting to know if such a transition exists in flows with an almost Keplerian rotation. In the last case setting free boundaries leads to the hypersonic character of the basic flow with Mach number $M>>1$.  The stability of toroidal flows was studied extensively in the 1980s and 1990s by means of the traditional eigenvalue analysis which has been widely used since the well-known investigations of Lord Rayleigh. The growing non-axisymmetric {\\it definite frequency modes} were discovered in the barotropic tori with a positive radial gradient of the specific angular momentum. This phenomenon was cold the Papaloizou-Pringle instability in honour of the authors who published their work in 1984. In numerous subsequent papers this sort of instability was studied in two- and three-dimensional approximations by both analytical and numerical methods \\cite{b24,b27,b1,b25,b3,b26,b28,b2,b22}. It was found that there are growing surface gravity and sound waves. The general picture of the instability is constructed by these modes which grow due to the mutual coupling as well as by the Landau mechanism, i.e. due to interaction with the background flow \\cite{b8,b5}. However, as the rotation profile approaches the Keplerian law the gravity surface waves become damping while the sonic instability ceases because of very small increments \\cite{b3,b4}. This fact actually dropped the subject since the near-Keplerian rotation is of the most interest in connection with the problem of angular momentum transport in the accretion disk theory. Besides, the growing linear modes can exist owing to additional degrees of freedom in the flow. For instance, \\cite{b39} and \\cite{b40} claimed that stratorotational instability emerges in vertically stratified rotationally stable flows in the shearing sheet approximation.  However, instability of axisymmetric flows can be considered from a quite different viewpoint. It is closely associated with the formulation of the {\\it initial value problem} for perturbations. Indeed, to examine the stability thoroughly one must show that {\\it any} initial perturbation will not ever affect the background flow. However, the absence of growing eigenmodes guarantees only the asymptotic stability and tells us nothing about the behaviour of perturbations at finite time intervals. The generalized approach to the hydrodynamical instability was realized and declared in the most distinctive form among the fluid physicists in the middle 1990s \\cite{b9,b10}. It is often called the nonmodal stability theory and concentrates on the initial value problem, partially on searching for optimal initial perturbations that can exhibit large transient growth even in asymptotically stable flows. Yet the role of this linear mechanism is usually emphasized in connection with the transition to turbulence \\cite{b32}. According to the so called bypass concept the linear dynamics is responsible for extraction of the perturbation energy from the background. It is supposed that nonlinear processes only redistribute energy between the modes. The general description and some references to hydrodynamical literature on this subject can also be found in \\cite{b13} and in application to astrophysical problems in the more recent paper \\cite{b14}.  Mathematically, the evolution of small perturbations is governed by a linear dynamical operator which in general can be normal or non-normal. In the first case eigenmodes are orthogonal and the eigenanalysis gives the complete picture of the temporal evolution of perturbations. Strictly speaking, the largest possible growth of perturbations is determined at any moment by the highest eigenvalue \\cite{b21}. In the case of non-normal dynamical operator eigenmodes become non-orthogonal and the substantial temporal growth of perturbations is possible even if there are no growing eigenmodes \\cite{b9}.  The nonmodal approach to {\\it local} perturbations was first pointed out also in astrophysical literature in \\cite{b11} and \\cite{b12}. The local framework allows us to perform the spatial Fourier transform and thus to study directly the temporal behaviour of perturbations. This direction was developed then in a number of papers \\cite{b18,b31,b30,b29,b15,b16,b17}.  Another branch of astrophysical investigations that use nonmodal approach is devoted to global dynamics in turbulent accretion disks. For instance, \\cite{b19} considered the dynamics of small stochastic perturbations in an incompressible Keplerian disk and found the resulting steady-state coherent structures enhancing the angular momentum transfer outward. It is also worth mentioning the work \\cite{b20} which studied the transient growth of axisymmetric perturbations in a thin accretion disk with account of compressibility and vertical structure of the flow.  Here we concentrate on the global two-dimensional dynamics of the bounded laminar axisymmetric flows which were studied actively 15-20 years ago. We do not introduce viscosity since its action is negligible over the first $\\sim 10-50$ rotation periods but we include compressibility which is necessary to do since the sound speed is always comparable or less than the shear velocity difference when we are dealing with free boundaries. First we briefly describe the solution of the boundary problem, i.e. the calculation of eigenmodes. This will be followed by the discussion of the optimization algorithm which allows us to determine the specific group of eigenmodes with the strongest transient growth. This is actually the standard mathematical procedure that has been used by many authors. Then we present the results and make the conclusions. ",
        "conclusions": "Our study shows the significance of the transient growth concept in application to astrophysical problems. We considered here the dynamics of {\\it global} perturbations in a sense that free boundary conditions which are natural in astrophysical practice were involved. Previous studies of a pure hydrodynamical instability in axisymetric flows with free boundaries have been carried out by methods of the traditional eigenanalysis, i.e. by searching for growing modes with exponential dependence on time. This way was proved to be unable to solve the problem of angular momentum transfer in almost Keplerian flows with finite compressibility, since the surface gravity eigenmodes are damping \\cite{b3} and the growing sonic eigenmodes have negligible increments. However, the present two-dimensional analysis demonstrates that situation noticeably changes if one turns to the dynamics of a linear {\\it combination} of eigenmodes. In fact, we studied mainly a finite number of slow neutral eigenmodes from acoustic spectrum of the basic flow with $\\omega<m\\Omega(r_2)$. It turns out that such a combination exhibits a substantial transient acoustic energy growth $g$ on the dynamical timescale. Its magnitude strongly depends on parameters of the problem, i.e. the azimuthal wavenumber of the modes, the radial size and the rotation profile of the background flow. The numerical tests produce $g\\propto 10^2$ in a few typical Keplerian periods. Possibly the most notable fact is that $g$ grows as we approach the Keplerian rotation. This result is opposite to the conclusions of eigenanalysis according to which the sonic instability ceases to affect the basic flow just because of the small sound speed. If one introduces the Mach number $M$ as the ratio of the shear velocity difference between the boundaries $r_1$ and $r_2$ to the maximum sound speed in the flow, then the maximum optimal growth and the corresponding time interval are approximately proportional to $M$. So the results of the present two-dimensional analysis indicate that the flat hypersonic tori in the vicinity of gravitating objects can exhibit a large linear transient growth of global perturbations which may be related to the problem of angular momentum transfer. However, the conclusions we made here certainly must be checked by complete three-dimensional calculations.  \\vspace{1cm}  {\\noindent \\bf Acknowledgements}  \\vspace{0.2cm}  We would like to thank K.A. Postnov for careful reading of the manuscript. This paper was supported by grant RFFI 06-02-16025."
    },
    "0808/0808.3213_arXiv.txt": {
        "abstract": "{Various studies have established that the dynamical M/L ratios of ultra-compact dwarf galaxies (UCDs) tend to be at the limit or beyond the range explicable by standard stellar populations with canonical IMF. We discuss how IMF variations may account for these high M/L ratios and how observational approaches may in the future allow to discriminate between those possibilities. We also briefly discuss the possibility of dark matter in UCDs. } ",
        "introduction": "In recent years, significant effort has been put into studying the internal dynamics of extragalactic compact stellar systems in the mass regime of massive globular clusters and ultra-compact dwarf galaxies ($10^6<M/M_{\\sun}<10^8$) (Drinkwater et al. 2003, Martini \\& Ho 2004, Ha\\c{s}egan et al. 2005, Maraston et al. 2004, Rejkuba et al. 2007, Evstigneeva et al. 2007, Hilker et al. 2007, Mieske et al. 2008; see also Dabringhausen et al. 2008 and Forbes et al. 2008). A striking outcome of these studies is that the dynamical M/L ratios of massive compact stellar systems are on average about two times larger than for normal globular clusters of comparable metallicity. This is illustrated in Fig.~\\ref{fig1}. The rise of M/L ratios starts at about 2$\\times 10^6 M_{\\odot}$, separating ordinary globular clusters at lower masses, from the so-called ultra-compact dwarf galaxies, at higher masses (see also Mieske et al. 2008). Possible reasons for these high M/L ratios include extreme stellar mass functions (Mieske \\& Kroupa 2008; Dabringhausen, Hilker \\& Kroupa 2008) or densely packed dark matter (Goerdt et al. 2008).     \\begin{figure} \\includegraphics[width=83mm]{UCD_mass_ML_alltalk.eps} \\caption{Mass vs. M/L for compact stellar systems. Literature sources are Mieske et al. (2008, green dots), Rejkuba et al. (2007, cyan), Ha\\c{s}egan et al. (2005, red) Evstigneeva et al. (2007, magenta), Hilker et al. (2007, green asterisk), and Milky Way globular clusters (McLaughlin \\& van der Marel 2005, black dots).  The vertical dashed line at 2$\\times 10^6 M_{\\odot}$ indicates the approximate mass where the relaxation time is equal to one Hubble time (Mieske et al. 2008; Dabringhausen et al. 2008). All M/L ratio estimates have been normalised to solar metallicity (Mieske et al. 2008).  The horizontal dotted lines indicate the M/L ratios expected for single stellar populations of age 13, 9, and 5 Gyrs (from top to bottom, based on Bruzual \\& Charlot 2003 and Maraston et al. 2005).} \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0808/0808.3025_arXiv.txt": {
        "abstract": "We re-examine the correlation between the colors and the inclinations of the Classical Kuiper Belt Objects (CKBOs) with an enlarged sample of optical measurements. The correlation is strong ($\\rho=-0.7$) and highly significant ($>8\\,\\sigma$) in the range $0^{\\circ}-34^{\\circ}$. Nonetheless, the optical colors are independent of inclination below $\\approx 12^{\\circ}$, showing no evidence for a break at the reported boundary between the so-called dynamically ``hot'' and ``cold'' populations near $\\approx 5^{\\circ}$. The commonly accepted parity between the dynamically cold CKBOs and the red CKBOs is observationally unsubstantiated, since the group of red CKBOs extends to higher inclinations. Our data suggest, however, the existence of a different color break. We find that the functional form of the color-inclination relation is most satisfactorily described by a non-linear and stepwise behavior with a color break at $\\approx 12^{\\circ}$. Objects with inclinations $\\geqslant 12^{\\circ}$ show bluish colors which are either weakly correlated with inclination or are simply homogeneously blue, whereas objects with inclinations $<12^{\\circ}$ are homogeneously red. ",
        "introduction": "The Kuiper Belt is a disk of icy bodies having semi-major axes larger than that of Neptune. Its members are usually known as Kuiper Belt Objects (KBOs) or Trans-Neptunian Objects (TNOs). The distribution of their orbits is structured leading to the identification of several dynamical families.  Resonant KBOs are those which are trapped in mean-motion resonances with Neptune (those trapped in the 3:2 resonance are also known as Plutinos).  Scattered KBOs, also known as Scattered Disk Objects, are essentially highly eccentric KBOs under strong gravitational influence of Neptune. Classical KBOs (CKBOs) possess relatively circular orbits that are neither located in any strong mean-motion resonance with Neptune nor strongly subject to its gravitational influence.  Since their discovery in 1992 more than 1200 KBOs have been identified. Due to their faintness only about 50 can be spectroscopically studied with the currently available instruments.  Multicolor photometry provides, however, a first-order approximation of their spectra, hence of their surface composition.  Most KBOs can be studied photometrically and about 230 objects have at least one measured color. Their surface colors have shown to be most diverse, ranging from neutral and even slightly blue (relative to the Sun) to extremely red, suggesting a large compositional diversity \\citep[see review by][]{Dor+08}.  The origin of the color diversity remains unclear. Various suggestions have been made in the context of collisional resurfacing \\citep{LuuJew96, GilH02, Del+04}. Nevertheless, none of the proposed models has been able to consistently explain the colors \\citep{JewLuu01, TheDor03, Del+04}.  Another possibility is that the observed color differences reflect primordial compositional variations \\citep[e.g.][]{TRC03}. Such compositional differences would be hard to explain if KBOs formed {\\it in situ}, since the temperature difference between 30 and 50 AU is a very modest $\\approx 10$ K. However, larger temperature and compositional differences might be possible if the KBO population, or part of it, did not form in place. Some dynamical models, in fact, suggest outward migration of KBOs \\citep[e.g.][]{Malho95, Gomes03}. Although there are no detailed chemical studies to address how varied such compositions would be and how these would reflect in the surface colors, these dynamical models imply a link between the current orbital inclinations of classical KBOs and their presumed location of origin.  As a whole, KBOs do not show significant correlations between their colors and orbital parameters such as semi-major axis or perihelion distance. On the other hand, the CKBOs do show a correlation between orbital inclination and optical color \\citep{TegRom00, TruBro02} and a correlation between perihelion and color \\citep{TegRom00, Peix+04}. In parallel, several works have pointed out the existence of two groups of CKBOs with a separation at $\\approx 5^{\\circ}$ in inclination. The two groups, usually referred to as ``cold'' ($i\\lesssim 5^{\\circ}$) and ``hot'' ($i\\gtrsim 5^{\\circ}$), have been identified using not only orbital properties \\citep{Brown01, Ell+05} but also physical properties such as size and binarity \\citep{LS01,Peix+04,Gul+06, Noll+08}. Given its potential importance, we re-examine the color-inclination relation using a new data set and appropriate statistical tests. ",
        "conclusions": "The main goal of this work has been to investigate the color-inclination trend seen for Classical Kuiper Belt Objects (CKBOs). We have analyzed a sample of $B-R$ colors of 69 objects, excluding 2 apparent outliers as well as colors with error bars larger than $0.21$. Objects were classified as CKBOs according to a definition by \\cite{LykMuk07}. Orbital inclinations, denoted $i_k$, were calculated relative to the Kuiper Belt Plane following \\cite{Ell+05}. Our results may be summarized as:  \\begin{enumerate}  \\item The linear $B-R$ color-inclination correlation of CKBOs measured over the full range of inclinations from $0\\degr$ to $34\\degr$ is $\\rho=-0.70^{+0.09}_{-0.07}$, corresponding to a significance level larger than $8\\,\\sigma$. This is a strong and highly significant correlation, consistent with previously published values.  \\item In contrast, the $B-R$ colors of CKBOs with inclinations $i_k\\leqslant 12.0^{\\circ}$$^{+0.5}_{-1.5}$ are statistically uncorrelated with inclination and are well described by $B-R=1.71\\pm0.11$.  \\item CKBOs with $i_k >12.0^{\\circ}$$^{+0.5}_{-1.5}$ show a slight color vs. inclination dependence following $(B-R)=-0.0159\\,i_k+1.703$. The data are also formally consistent with a constant but bluer color, $B-R = 1.33\\pm0.20$, for  $i_k \\geqslant 12.5\\degr$, and a constant red color $B-R=1.70\\pm0.11$ for $i_k <12.5\\degr$.  \\item The data provide no evidence for a break or change in the $B-R$ color distribution at the boundary between the dynamically hot and cold populations, purportedly near $i_k \\approx 5^{\\circ}$. In this sense, we find no observational support for the frequently-cited parity between red CKBOs and the dynamically cold population. The CKBOs are red up to $i_k \\approx 12^{\\circ}$ and, therefore equally red into the dynamically hot population.  \\end{enumerate}"
    },
    "0808/0808.1811_arXiv.txt": {
        "abstract": "~~~\\\\ \\begin{center} {\\bf Abstract} \\end{center}  The detection of a ``Cold Spot\" in the CMB sky could be explained by the presence of an anomalously large spherical underdense region (with radius of a few hundreds ${\\rm Mpc}/h$) located between us and the Last Scattering Surface. Modeling such an underdensity with an LTB metric, we investigate whether it could produce significant signals on the CMB power spectrum and bispectrum, via the Rees-Sciama effect. We find that this leads to a bump on the power spectrum, that corresponds to an ${\\cal O}(5\\%-25\\%)$ correction at multipoles $5 \\leq \\ell \\leq 50$; in the cosmological fits, this would modify the $\\chi^2$ by an amount of order unity. We also find that the signal should be visible in the bispectrum coefficients with a signal-to-noise $S/N\\simeq {\\cal O} (1-10)$, localized at $10 \\leq \\ell \\leq 40$. Such a signal would lead to an overestimation of the primordial $f_{NL}$ by an amount $\\Delta f_{NL}\\simeq 1$  for WMAP and by $\\Delta f_{NL}\\simeq 0.1$ for Planck. ",
        "introduction": "The recent WMAP~\\cite{WMAP} experiment has measured with great accuracy the anisotropies of the Cosmic Microwave Background (CMB), whose features are in good agreement with the expectations from inflation of a Gaussian spectrum of adiabatic fluctuations, fully described by a nearly scale-invariant power spectrum. However, it has been pointed out by several authors that the data contain various unexpected features. Some of them are localized at very large angular scales, such as the low quadrupole, the alignment of the low multipoles (for a review see~\\cite{reviewlowl} and references therein) and the power asymmetry between the northern and southern hemispheres~\\cite{hemispherical}. Another anomaly is the presence of the so-called Cold Spot~\\cite{ColdSpot1,ColdSpot2,ColdSpot3}, which is a large circular region on an angular scale of about $10^\\circ$ that appears to be anomalously cold: the probability that such a pattern would appear from Gaussian primordial fluctuations is estimated to be about $1.8\\%$~\\cite{ColdSpot1,ColdSpot2,ColdSpot3}. So, while this could still be a statistical fluke, some authors have put forward the idea that it could instead be due to an anomalously large spherical underdense region of some unknown origin, on the line of sight between us and the Last Scattering Surface (LSS)~\\cite{Tomita,InoueSilk}. We may also remind that~\\cite{rudnick} has claimed that, looking at the direction of the Cold Spot in the Extragalactic Radio Sources (the NVSS survey), an underdense region is visible at redshift $z\\sim 1$ (see, however~\\cite{huterer} for a paper that challenges this claim). Another motivation for studying such objects is that an underdense region of two or three hundreds of ${\\rm Mpc}/h$ could be enough to give an acceptable fit to the Supernova data and the CMB without Dark Energy~\\cite{ABNV}, if we happen to live near its centre.  In this paper we also take the point of view that the Cold Spot could be due to such an underdense region, customarily denoted as a ``Void''. By modeling it through an inhomogeneous Lema\\^itre-Tolman-Bondi (LTB) metric (which also requires an overdense compensating shell), we compute the impact of such a Void on some observational quantities, focusing on the statistical properties of the CMB: in particular on the two-point correlation function (power spectrum) and the three-point correlation function (the bispectrum). This is interesting for the following reasons. First, if there is such a Void, does the power spectrum get a sizable correction? And, if yes, to what extent the estimation of the cosmological parameters is affected? Second, if there is such a Void, could it be detectable also in the bispectrum? Third, does the presence of such an object interfere with the measurement of a primordial non-gaussianity, which constitutes a very important piece of information in order to distinguish between models of inflation?   It is well known that, passing through a Void, photons suffer some blue-shift due to the fact that the gravitational potential is not exactly constant in time -- the so-called Rees-Sciama (RS) effect~\\cite{ReesSciama}. By RS effect we mean the one associated to the variation of the gravitational potential from non-linear effects. There can be an effect already at the linear level, usually referred to as the Integrated Sachs-Wolfe (ISW) effect, which however would vanish in a matter dominated flat Universe. The ISW effect would be significant only when a Dark Energy component becomes dominant with respect to matter. Hence, here we focus our attention only on the RS effect, which is always present, and briefly comment on the extension to a $\\Lambda$CDM Universe later on. Note that there is a second physical effect due to a Void on the line of sight, in addition to the RS blue-shift of the photons: the lensing of the primordial perturbations, which will be analyzed in detail in a companion paper~\\cite{LENS}.  It is also known that the RS effect scales as the third power of the comoving radius of the Void $L$ times the present value of the Hubble parameter $H_0$, $\\Delta T/T \\propto - (L H_0)^3$. For Void sizes compatible with the expectations from the usual structure formation scenarios, the RS effect happens to be suppressed with respect to the primordial temperature fluctuations ~\\cite{Fullana}. In order to produce a signal comparable to the measured temperature anisotropy of $|\\Delta T/T| \\sim 10^{-5}$, the Void should be of very large size, {\\it i.e.} $L \\sim (200-300) {\\rm Mpc}/h$: this would be at odds with the standard scenario of structure formation from Gaussian primordial fluctuations (at more than $10\\sigma$~\\cite{InoueSilk}). We may also mention that according to~\\cite{manyvoids,Szapudi} there are many localized regions in the sky (both underdense and overdense) of about $100 {\\rm Mpc}/h$, which would be responsible for the correlations between the CMB and the Large Scale Structures: as~\\cite{SarkarVoids} has recently stressed the existence of these regions is already at odds with the usual structure formation scenario. Even though we do not address in this paper the issue of the primordial origin of such large objects, we mention some possibilities. For example, one could consider certain models of inflation, such as the ones of the ``extended'' type~\\cite{extended,extendedDMN,extendedBN}, with the possibility of tunneling events and nucleation of bubbles, which would appear as Voids in the sky today. Alternatively, one could also imagine the presence of non-Gaussian features in the primordial fluctuations which would seed a large Void or even non-conventional structure formation histories in the late-time Universe which could enhance the probability of having large Voids.  In the case of an explanation of the Voids via tunneling events, it is interesting to note that if a nucleation process has small probability (per unit volume and time) compared to the Hubble rate (at some time during inflation) then the number of anomalous Voids could well be very small, such as having just one or a few of them in the present observable Universe. In addition, since the LSS is a rather thin shell whose volume is much smaller than the total volume inside it, if there are only few Voids it is more likely that they are located along the line of sight, rather than at the LSS\\footnote{It is possible, also, to imagine that some primordial process could generate a coherent spherical region on the LSS surface, with small density contrast, on top of the primordial Gaussian spectrum, and which would produce a $10^{-5}$ fluctuation. However, we do not consider here such possibilities.}. This is an important point because a Void at the LSS would have a huge impact on the CMB and, as a consequence, its size would be strongly constrained: $L\\lesssim 4 {\\rm Mpc}/h$~\\cite{emptyVoids,extendedDMN}. On the contrary, Voids on the line of sight are not subject to the latter constraint.  Note also that other authors have put forward alternative ideas, such as the idea that the Cold Spot could be due to a topological defect, in particular a cosmic texture~\\cite{texture}. As we are going to discuss, some of our considerations apply in a similar fashion if the origin of the Cold Spot is a texture, rather than a Void.  The paper is organized as follows. In sect.~\\ref{profiles} we calculate the RS-induced temperature profile of the CMB photons passing through a large Void with the physical characteristics of the Cold Spot. The impact on the power spectrum is discussed in sect.~\\ref{powerspectrum}, while sect.~\\ref{bispectrum} deals with the impact on the bispectrum and on the overestimation of $f_{NL}$ that would be done by neglecting the presence of such a Void. Finally, in sect.~\\ref{multivoids} we extend some of the considerations to the case of several objects in the sky. In fact, it is conceivable that if there is one such structure, there may be other ones, maybe of smaller size and less visible -- see {\\it e.g}.~\\cite{manyvoids, Szapudi} for some other candidates in the CMB sky. Our conclusions are drawn in sect.~\\ref{concl}. ",
        "conclusions": "\\label{concl}  Motivated by the so-called Cold Spot in the WMAP data, we have studied in this paper the impact on statistical CMB predictions, in particular the two and three point correlation functions, of the presence of an anomalously large Void along the line of sight. Indeed, the existence of such a Void could be at the origin of the Cold Spot, due to the Rees-Sciama effect (the Lensing effect is analyzed in the companion paper~\\cite{LENS}).  First, we have computed the temperature profile using an LTB solution of the Einstein equations, matched to an FLRW metric. Then, we have computed its impact on statistical predictions for the CMB, {\\it assuming} this structure to be uncorrelated with the Primordial fluctuations. As suggested by~\\cite{texture}, we consider the angular size of such a Void to be about $10^\\circ-18^\\circ$ and its temperature at the centre to be characterized by $\\Delta T=-(190\\pm 80) \\mu K$.  For the  power spectrum the results are as follows. The RS effect is non-negligible: % we predict a bump of $5 \\% -25\\%$ to be added to the Primordial spectrum, localized at multipoles $5 \\leq \\ell \\leq 50$. This should lead to a variation in the $\\chi^2$ for the WMAP fits of order unity.  Then we have studied the impact on the bispectrum coefficients. For the RS effect we have found that the Signal-to-Noise ratio is larger than unity at $\\ell \\gtrsim 40$ for most of the parameter space, and therefore this should already be visible in the available data. Through the bispectrum, we have studied also the impact of such a structure on the determination of the primordial non-gaussianity. The RS bispectrum signal is large but localized at low multipoles ($10 \\leq \\ell \\leq 50 $), so it has a small impact on high resolution experiments, which can go up to very large multipoles: the overestimation of $f_{NL}$ due to the RS effect turns out to be $\\Delta f_{NL}^{(RS)}\\simeq 1$ for WMAP and $\\Delta f_{NL}^{(RS)}\\simeq 0.1$ for Planck. Using the already existing WMAP1-year constraints~\\cite{WMAP1GAUSS} on $f_{NL}$ at low $\\ell$, one can exclude extreme values of the temperature contrast of the Void. For example, values for $\\Delta T/T$ larger than $8\\times 10^{-5}$ for a Cold Spot with diameter of $18^\\circ$ are likely to be excluded, via a full analysis. So, we conclude that the bispectrum is a valuable tool for constraining an anomalously large Void, through the RS effect.  We have also considered the 50 Voids whose detection has been claimed in~\\cite{Szapudi}. These Voids have mean angular diameter of about $10^\\circ$ and average temperature at the centre characterized by $\\Delta T \\approx -11 \\mu K$. The effect on the two-point correlation function is about $0.5\\%$, hence much smaller than the one due to the large Void considered to explain the Cold Spot. The effect on the three-point correlation function has not been studied in detail but it is expected to lead to a Signal-to-Noise ratio smaller than one.  Finally, we have also applied our considerations to the case in which the Cold Spot is explained by a cosmic texture~\\cite{texture} rather than a large Void: the effect on the power spectrum is similar but somewhat smaller; the three-point correlation function leads also to a smaller Signal-to-Noise ratio, which however turns out to be above unity for some part of the parameter space."
    },
    "0808/0808.2435_arXiv.txt": {
        "abstract": "We have carried out a spatial-kinematic study of three proto-planetary nebulae, IRAS 16594$-$4656, Hen 3-401, and Rob 22. High-resolution \\mh images were obtained with NICMOS on the {\\it HST} and high-resolution spectra were obtained with the Phoenix spectrograph on Gemini-South. IRAS 16594$-$4656 shows a ``peanut-shaped'' bipolar structure with \\mh emission from the walls and from two pairs of more distant, point-symmetric faint blobs.  The velocity structure shows the polar axis to be in the plane of the sky, contrary to the impression given by the more complex visual image and the visibility of the central star, with an ellipsoidal velocity structure. Hen 3-401 shows the \\mh emission coming from the walls of the very elongated, open-ended lobes seen in visible light, along with a possible small disk around the star.  The bipolar lobes appear to be tilted 10$-$15$\\arcdeg$ with respect to the plane of the sky and their kinematics display a Hubble-like flow. In Rob 22, the \\mh appears in the form of an ``S''-shape, approximately tracing out the similar pattern seen in the visible. \\mh is especially seen at the ends of the lobes and at two opposite regions close to the unseen central star.  The axis of the lobes is nearly in the plane of the sky.  Expansion ages of the lobes are calculated to be $\\sim$1600 yr (IRAS 16594$-$4656), $\\sim$1100 yr (Hen 3-401), and $\\sim$640 yr (Rob 22), based upon approximate distances. ",
        "introduction": "\\label{intro}  Proto-planetary nebulae (PPNs) are objects in transition between the asymptotic giant branch (AGB) and planetary nebula (PN) stages of stellar evolution.  Studies over the past decade, particularly imaging studies with the {\\it Hubble Space Telescope} ({\\it HST}), have shown that this transition stage is key to understanding the shaping of PNs \\citep{balick02}.  In the interacting stellar winds model \\citep{kwok82}, which was later generalized to include an equatorial density gradient \\citep{balick87}, it was assumed that a fast wind from the central star interacted with the detached, slowly-moving remnant AGB envelope to shape the nebula.  {\\it HST} images of PPNs show that many and perhaps most of them display a basic bipolar structure \\citep{ueta00,su01,sahai07,siod08}. These are shapes that appear to be further developed in the PN stage. In addition, a point symmetry is often seen; \\citet{sahai98a} attribute this to collimated jets, which may be episodic and change their direction, and which they advocate as the main shaping agents of PNs.  Studying \\mh in PPNs is one of the best ways to probe the presence of a fast, perhaps collimated, post-AGB wind and its effect in shaping the nebula. \\mh is the main constituent of the detached circumstellar envelope around a PPN. While a fast wind (V$>$100 km s$^{-1}$) has the energy to dissociate the \\mh on direct impact, it can also produce shocks in the medium that move at a slower speed and collisionally excite the \\mh. Recent surveys have shown the presence of shock-excited \\mh in bipolar PPNs \\citep{garher02, kelly05}. Detailed high-resolution \\mh studies have been published of approximately a half dozen PPNs \\citep{sahai98, kastner01, cox03, davis05, hri06, vds08}.  In this paper we present the results of a detailed \\mh emission study of three bipolar PPNs, IRAS 16594$-$4656 (``Water Lily Nebula''), Hen 3$-$401 (IRAS 10178$-$5958), and Rob 22 (IRAS 10197$-$5750).  All three have relatively hot central stars for PPNs, B and A spectral types, and all possess large infrared excesses due to circumstellar dust \\citep{partha89,garlar99a}. All three are or would be classified as DUPLEX (``DUst Prominent, Longitudinally-EXtended'') nebulae in the classification system of \\citet[see also \\citet{siod08}]{ueta00}.  Rob 22 has a nebula that looks like a pair of butterfly wings.  At high spatial resolution, the lobes show a filamentary structure in visible light.  In addition to bipolar structure, some point symmetry is apparent in the lobes.  A large halo extends out to 25$\\arcsec$ \\citep{sahai99a}. The nebula appears to be viewed essentially edge on, with a dark lane obscuring the star in visible light. In the more detailed morphological classification system of \\citet{sahai07}, this object is classified as Bcw,ml,ps(s),h(e): bipolar nebula with closed lobes, central obscuring waist, minor lobes present, point-symmetric shape, and enlongated halo. The spectral type is A2~Ie, based on light reflected off the lobes of the nebula \\citep{allen78}. The circumstellar envelope appears to have a dual chemistry, with an oxygen-rich region evidenced by OH maser emission \\citep[OH284.18-0.79;][]{allen80} and crystalline silicate emission \\citep{mol02} and a carbon-rich region evidenced by PAH emission \\citep{mol97}.  The nebula of Hen 3-401 possesses long, narrow lobes.  It seems to be viewed at a slightly larger angle with respect to the plane of the sky, with the central star appearing faintly between the lobes \\citep{sahai99b}. The morphological classification of this object is Bow*(0.6),sk: bipolar nebula with open lobes, central obscuring waist with star evident (at 0.6 $\\mu$m), with a skirt-like structure \\citep{sahai07}. The spectral type of the star is Be, and the spectrum shows strong Balmer emission lines and lots of permitted and forbidden low-excitation emission lines \\citep{allen78, garlar99b}. The circumstellar envelope appears to be carbon-rich, since it shows CO emission \\citep{loup90, buj91} and PAH emission \\citep{partha01} but not OH \\citep{silva93} emission or crystalline silicates \\citep{partha01}.  In IRAS 16594$-$4656, the star is clearly seen and relatively bright compared to the nebula, giving the impression that the nebula is at a larger orientation angle. The nebula has the appearance of a basic bipolar structure with a pronounced point symmetry, consisting of three pairs of oppositely-directed, slightly curved, thin lobes.  Also seen is a smaller elliptical structure oriented from northwest to southeast \\citep{hri99}. The morphological classification of this object is Mcw*(0.6),an,ps(m,an),h(a): multipolar nebula with closed lobes, central obscuring waist with star evident (at 0.6 $\\mu$m), ansae present, point symmetry with two or more pairs of diametrically opposed lobes and with ansae, and a halo with centrosymmetric arcs \\citep{sahai07}. Its spectrum shows Balmer emission lines \\citep{garlar99a} and has been classified as B7 based on the H$\\alpha$ line \\citep{vds00}.  \\citet{rey02}, on the basis of high-resolution spectra, fit the spectrum with T$_{\\rm eff}$=14,000 K and log{\\it g}=2.1, slightly hotter than $\\beta$ Ori (B8~I); this indicates a spectral type of B5-7 Ie. He also observed many emission lines.   The nebula appears to be carbon rich, showing PAH features and displaying the 21 $\\mu$m and 30 $\\mu$m features seen in carbon-rich evolved stars \\citep{garlar99a,volk02}. \\citet{vds08} recently published a detailed kinematic and morphologic study of this object using H$_2$ emission, and we will refer to their results in the present study.  We begin with an examination of the high-resolution \\mh images (see \\citet{sahai00} for preliminary images of Rob 22 and Hen 3$-$401 and \\citet{hri04} for preliminary images of IRAS 16594$-$4656), then present new high-resolution, long-slit \\mh spectra, and then follow with a discussion of the kinematics. From these, we are able to determine properties of the collimated winds in these three PPNs and how they have shaped the surrounding nebulae. ",
        "conclusions": "The high-resolution {\\it HST}-NICMOS \\mh images of these three PPNs have allowed us to determine the spatial location of their \\mh emission and the spatially-resolved \\mh spectroscopic observations have allowed us to determine the kinematics of the \\mh emitting regions.  In all three cases, the systemic and expansion velocities are similar to those determined from the molecular-line CO or OH measurements.  We find the following results for these three PPNs.  {\\it IRAS 16594$-$4656}: While the V and H-band images show a complex multi-lobe structure, the \\mh images shows a clear bipolar, peanut shape, although with variations in density (``holes'') and structure \\citep[see][]{vds08}.  The \\mh emission originates along the walls of the lobes (sides and ends), with fainter emission from more distant (ejected?) clumps.  The PV diagram shows the \\mh to arise in an expanding ellipsoidal velocity structure, which is in contrast to the bilobes of the density structure.  The kinematics indicates that the bipolar lobes are nearly in the plane of the sky (i$\\approx$10$\\arcdeg$); this differs from the earlier interpretations of the lobes as being at some intermediate orientation, but it is consistent with recent mid-IR imaging and near-IR polarization studies.  The lobes are estimated to have an age of $\\sim$1600 yr.  {\\it Hen 3-401}: The \\mh emission originates from the sides of lobes, which have open ends.  The lobes are tilted somewhat to the plane of the sky ($\\sim$10$-$15$\\arcdeg$), with the western lobe moving toward us, and they show an increasing velocity with radial distance (Hubble flow). An estimated age for the lobes is $\\sim$1100 yr. The open ends and unhindered outflow may be the consequence of its higher degree of collimation (greater linear momentum).  {\\it Rob 22}: The \\mh emission originate primarily from ends of the S shaped nebula and from regions of the S shape near the obscured central star.  The nebula is nearly in the plane of the sky, consistent with the absence of a visible star due to obscuration by a disk.  An age of $\\sim$630 yr is estimated for the lobes.  \\mh surveys of PPNs have shown that \\mh emission is commonly found in those with a bipolar morphology.  As can be seen from this study, the combination of \\mh high-resolution images and spatially-resolved, high-resolution spectra provides valuable insight into the structure and shaping mechanisms for these bipolar nebulae."
    },
    "0808/0808.2603_arXiv.txt": {
        "abstract": "We analyze archived {\\it Chandra} and {\\it XMM-Newton} \\mbox{X-ray} observations of 536 Sloan Digital Sky Survey (SDSS) Data Release 5 (DR5) quasars (QSOs) at $1.7 \\le z \\le 2.7$ in order to characterize the relative UV and \\mbox{X-ray} spectral properties of QSOs that do not have broad UV absorption lines (BALs).  We constrain the fraction of \\mbox{X-ray} weak, non-BAL QSOs and find that such objects are rare; for example, sources underluminous by a factor of 10 comprise $\\la$2\\% of optically-selected SDSS QSOs.  \\mbox{X-ray} luminosities vary with respect to UV emission by a factor of $\\la$2 over several years for most sources.  UV continuum reddening and the presence of narrow-line absorbing systems are not strongly associated with \\mbox{X-ray} weakness in our sample.  \\mbox{X-ray} brightness is significantly correlated with UV emission line properties, so that relatively \\mbox{X-ray} weak, non-BAL QSOs generally have weaker, blueshifted \\ion{C}{4}~$\\lambda$1549 emission and broader \\ion{C}{3}]~$\\lambda$1909 lines.  The \\ion{C}{4} emission line strength depends on both UV and \\mbox{X-ray} luminosity, suggesting that the physical mechanism driving the global Baldwin effect is also associated with \\mbox{X-ray} emission. ",
        "introduction": "}  It has been known for some time that the UV and \\mbox{X-ray} luminosities of quasars (QSOs) are correlated \\citep[e.g.,][and references therein]{at82}, and recent studies have carefully quantified this relation across $\\approx$5 orders of magnitude in UV luminosity \\citep[e.g.,][]{sbsvv05, ssbaklsv06, jbssscg07}.  Such studies inform ongoing efforts to understand the structure and physics of QSO nuclear regions, providing quantitative constraints on models of physical associations between UV and \\mbox{X-ray} emission.  Because UV photons are generally believed to be radiated from the QSO accretion disk while \\mbox{X-rays} are produced in the disk corona \\cite[e.g.,][and references therein]{rn03}, the \\mbox{UV/X-ray} luminosity relation is an indication of the balance between accretion disks and their coronae.  For example, a large fraction of intrinsically X-ray weak sources would suggest that coronae may frequently be absent or disrupted in QSOs.  In this study, we combine results from recent optical/UV and \\mbox{X-ray} observations of hundreds of QSOs in order to constrain the fraction of sources that are anomalously \\mbox{X-ray} weak, and also to test for additional physical effects that contribute to the scatter observed in the relation between UV and \\mbox{X-ray} luminosities.  Radio-loud QSOs are well known to be relatively \\mbox{X-ray} bright \\citep[e.g.,][]{wtgz87, blvsbbwg00}, while QSOs with broad UV absorption lines (BALs) are \\mbox{X-ray} faint \\citep[e.g.,][]{gsahfbftm95, blw00}.  We wish to understand the emission processes of QSOs without the added complexity of the strong X-ray absorption associated with UV BAL outflows or the enhanced X-ray emission that may be associated with a radio jet.  In order to accomplish this, BAL and radio-loud QSOs must be carefully removed from samples.  This process can be complicated in cases where available spectra do not extend to wavelengths where strong UV BAL absorption may occur.  In this study, we compare the UV and \\mbox{X-ray} properties of optically-selected, radio-quiet Sloan Digital Sky Survey \\citep[SDSS; e.g.,][]{y+00} QSOs that are known not to host BAL outflows along the line of sight.  Careful screening enables us to study the ways in which these ``ordinary'' QSOs deviate from the general relation between UV and \\mbox{X-ray} luminosities.  We quantify the scatter about the best-fit relation between UV and \\mbox{X-ray} luminosities, and in particular constrain the fraction of optically-selected QSOs that are intrinsically weak \\mbox{X-ray} emitters.  Secondly, we search for correlations between UV emission/absorption properties and {\\it relative} \\mbox{X-ray} brightness, where the term ``relative'' indicates the observed \\mbox{X-ray} brightness compared to that expected for an average QSO with the same UV luminosity.  Beyond the UV/X-ray luminosity relation, our findings place additional constraints on physical models that relate QSO UV and \\mbox{X-ray} emission processes.  The significant amount of scatter in the \\mbox{X-ray} brightness of individual sources compared to that of ``average'' QSOs with the same UV luminosity \\citep[e.g.,][]{ssbaklsv06} indicates that some unmodeled physical factors influence the relation between UV and \\mbox{X-ray} luminosities.  Besides the cases of radio-loud and BAL QSOs mentioned above, possible additional causes of scatter include additional \\mbox{X-ray} absorption that is not associated with UV BALs, intrinsically weak (or strong) \\mbox{X-ray} emission, and time variability.  Studies of \\mbox{X-ray} brightness in QSOs use different methods to quantify UV absorption and \\mbox{X-ray} weakness (complicating comparisons between different samples), and intensive UV and \\mbox{X-ray} spectroscopic campaigns have been performed for only a few QSOs.  As a result, we do not have strong constraints on how frequently \\mbox{X-ray} absorption, intrinsically faint \\mbox{X-ray} emission, and time variability cause \\mbox{X-ray} weakness in QSO samples.  In the remainder of this section, we briefly review published cases involving these physical scenarios.  This discussion includes all the non-BAL QSOs listed in \\citet{blw00} and five of the six sources from \\citet{ybsv98} with $\\alpha_{OX} < -2$.\\footnote{The sixth source, QSO~$0316-346$, was identified from optical spectra by \\citet{mppb88}.  UV spectra are required to search for BAL absorption, but no UV spectra of this source are present in the MAST or HEASARC archives.}  We also discuss cases of unusual X-ray weakness in recent studies of individual sources.  \\subsection{X-Ray Absorption\\label{xAbsCauseXWSec}}  Even for QSOs that do not show UV BAL absorption, strong \\mbox{X-ray} absorption may be the simplest explanation for \\mbox{X-ray} weakness.  Mrk~304 (PG~$2214+139$) is an \\mbox{X-ray} weak source that is often classified as a Seyfert~1 galaxy.  It does not show UV BAL absorption, although \\citet{lb02} found some evidence for moderate \\ion{C}{4} and \\ion{N}{5} line absorption.  In the \\mbox{X-rays}, it is strongly absorbed by a multizone, ionized absorber with a total column density $N_H \\sim 10^{23}$~cm$^{-2}$ \\citep{pjgsrs04, bpf04}.  Another \\mbox{X-ray} weak source, PG~$1126-041$, shows evidence for an ionized \\mbox{X-ray} absorber.  However, low-velocity, broad \\ion{C}{4} absorption has been found in some observations, indicating that PG~$1126-041$ should be classified as a (variable) BAL QSO \\citep{wbwyw99}.   \\subsection{Intrinsic X-Ray Weakness\\label{intrinsicCauseXWSec}}  \\mbox{X-ray} weakness cannot always be attributed to absorption.  The nearby ($z = 0.192$), narrow-line type 1 QSO PHL~1811 shows no evidence of BALs in its UV spectrum.  With an upper limit for a neutral \\mbox{X-ray} absorbing column of $N_H < 8.7 \\times 10^{20}$~cm$^{-2}$, the source is not strongly \\mbox{X-ray} absorbed \\citep{lhjgcp07}.  It has been consistently \\mbox{X-ray} weak since it was first observed almost 20 years ago with {\\it ROSAT}, although the \\mbox{X-ray} flux has varied by a factor $\\approx$5.  This variation (small compared to the degree of \\mbox{X-ray} weakness) suggested to \\citet{lhjgcp07} that the \\mbox{X-ray} emission was not scattered into view, as the medium responsible for the scattering would need to be implausibly small.  PHL~1811 therefore appears to be an intrinsically weak \\mbox{X-ray} emitter.  The optical/UV spectrum of PHL~1811 is also unusual \\citep{lhjc07}.  It is very blue and shows no forbidden or semi-forbidden line emission.  The \\ion{C}{4} $\\lambda 1549$ line emission is weak by a factor of $\\approx5$ compared to the composite spectrum of \\citet{fhfcwm91}.  In this study, we will test whether intrinsically \\mbox{X-ray} weak objects like PHL~1811 (which would have been flagged for SDSS spectroscopy based on its optical colors alone) are common, and whether these unusual emission line characteristics are associated with \\mbox{X-ray} weakness.  Another source, PG~$1011-040$, is relatively \\mbox{X-ray} weak by a factor of $\\sim$10 \\citep{blw00, gblemwi01}.  There is some evidence for \\ion{C}{4} absorption in the UV spectrum, but no strong indication of a BAL \\citep{blw00, lb02}.  It appears to be relatively unabsorbed in \\mbox{X-rays}, with an upper limit on the absorbing column density of $N_H \\le 5\\times 10^{21}$~cm$^{-2}$.  In contrast to PHL~1811, multiple observations of PG~$1011-040$ have not found much \\mbox{X-ray} luminosity variation \\citep{gblemwi01}.   \\subsection{UV and X-Ray Variability\\label{varCauseXWSec}}  Source variability may be an important factor contributing to the scatter in the UV/\\mbox{X-ray} luminosity relation.  Additional scatter is introduced by the fact that optical/UV and \\mbox{X-ray} observations are usually taken some time apart, and this time may reach up to years.  The narrow-line Seyfert~1 galaxy Mrk~335 has recently been reported to have decreased in \\mbox{X-ray} brightness by a factor of $\\approx$30 \\citep{gkg07}.  The authors suggested the variation may have been caused by the onset of heavy \\mbox{X-ray} absorption, which may also lead to the appearance of UV BALs.  PG~$0844+349$ has been observed to vary in the \\mbox{X-rays} by up to 60\\% on short (20~ks) time scales, and by a factor of $\\sim$10 on longer (multi-year) time scales \\citep[e.g.,][]{gblemwi01, bgbf03}.  The narrow-line Seyfert 1 galaxy WPVS~007 has dimmed in the \\mbox{X-rays} by a factor of 100 or more, while a \\ion{C}{4} BAL appears to be forming in the UV spectrum of this source \\citep{gslkon07}.   Finally, we note that our understanding of the causes of \\mbox{X-ray} weakness is limited by observing time scales.  Over longer time scales, \\mbox{X-ray} weak QSOs may show different properties (such as long-term emission or absorption variation), and may even cease to be \\mbox{X-ray} weak.  We have also noted in \\S\\ref{introSec} cases where UV BALs may be transient, and recent study has demonstrated that BALs evolve on multi-year (rest-frame) time scales \\citep[e.g.,][]{gbsg08}.  For these reasons, we are only able to place limits on the extent of \\mbox{UV/X-ray} variability (\\S\\ref{uVXVarConstraintSec}).  Future near-simultaneous, multi-wavelength observations of QSOs will improve on these constraints.  Throughout this work we use a cosmology in which $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M = 0.3,$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "}  We have analyzed the relationship between \\mbox{X-ray} luminosity and UV properties of SDSS DR5 QSOs at redshift $1.7 \\le z \\le 2.7$.  At these redshifts, we are able to identify \\ion{C}{4} BAL QSOs in our sample.  This has enabled us to place strong constraints on the fraction of intrinsically \\mbox{X-ray} weak QSOs, as well as to search for additional trends beyond the well-studied relationship between UV and \\mbox{X-ray} luminosities.  In particular, we find that: \\begin{enumerate} \\item{Non-BAL, radio-quiet QSOs which are very (intrinsically) \\mbox{X-ray} weak are rare in optically-selected samples, with $\\la$2\\% of SDSS QSOs having $\\Delta\\alpha_{OX} < -0.4$ and $\\la$1\\% of SDSS/BQS QSOs having $\\Delta\\alpha_{OX} < -0.67$.  Figure~\\ref{binomialProbOfltDAOXFig} provides upper limits on the fraction of sources that may be relatively \\mbox{X-ray} weaker than a given value of $\\Delta\\alpha_{OX}$.} \\item{The rms of the $\\Delta\\alpha_{OX}$ distribution for radio-quiet, non-BAL QSOs is about $0.1$, corresponding to a factor of $\\approx$2 spread in \\mbox{X-rays} relative to that estimated from the UV luminosity.  This places an upper limit on typical \\mbox{X-ray} variability with respect to the UV continuum.} \\item{While the amplitude of and relation between UV and X-ray variation in QSOs is not well-understood, estimates based on recent studies indicate that most, if not all, of the observed spread of $\\Delta\\alpha_{OX}$ can be attributed to variability.} \\item{The distribution of $\\Delta\\alpha_{OX}$ may not be normally distributed (at 97.5\\% confidence) in our sample $B$.  Perhaps physical processes such as \\mbox{X-ray} absorption are significant in some sources, broadening the wings of the distribution.} \\item{We find no strong evidence that reddening or the presence of narrow \\ion{Mg}{2} absorption is related to \\mbox{X-ray} weakness for non-BAL QSOs.} \\item{UV emission line properties such as \\ion{C}{4}~EW, \\ion{C}{3}]~FWHM, and perhaps \\ion{C}{4} wavelength are correlated to relative \\mbox{X-ray} brightness.} \\item{The \\ion{C}{4} emission EW depends on both UV and \\mbox{X-ray} luminosity.  The physics that drives the global Baldwin effect is apparently associated with \\mbox{X-ray} emission as well as UV emission.} \\item{Even after correcting for secondary trends (such as weak \\ion{C}{4} emission, which is associated with relative \\mbox{X-ray} weakness), objects that are as intrinsically \\mbox{X-ray} faint as PHL~1811 are rare.} \\end{enumerate}  We have quantitatively shown that luminous X-ray emission is essentially a universal property of optically-selected QSOs.  Future studies can test whether this result holds true for QSOs selected in other wavebands, such as the radio and infrared.  New optical surveys will extend the luminosity range of our sample, while UV (and infrared) spectroscopy will allow BAL QSOs to be identified (and $L_{2500~\\mathring{A}}$ to be measured) for sources in a wider redshift range.  The high detection rate for relatively short exposures in our sample has also demonstrated that wide, shallow X-ray surveys at high angular resolution are an effective way to study the X-ray properties of bright, optically-selected QSOs."
    },
    "0808/0808.2868_arXiv.txt": {
        "abstract": "We report the measurement of $\\nu$-$e$ elastic scattering  from \\bor\\  solar neutrinos with 3\\,MeV energy threshold by the Borexino detector in Gran Sasso (Italy).  The rate of solar neutrino-induced electron scattering events above this energy in Borexino is $0.217\\pm 0.038 (stat)\\pm 0.008 (syst)$~cpd/100\\,t, which corresponds to $\\Phi^{\\rm ES}_{\\rm ^8B}$ = {2.4 $\\pm$ 0.4$\\pm$ 0.1}$\\times$10$^6$~cm$^{-2}$ s$^{-1}$, in good agreement with measurements from SNO and SuperKamiokaNDE.  Assuming the \\bor\\ neutrino flux predicted by the high metallicity Standard Solar Model, the average \\bor\\ \\nue\\ survival probability above 3 MeV is measured to be 0.29$\\pm$0.10.  The survival probabilities for \\ber\\ and \\bor\\ neutrinos as measured by Borexino differ by  1.9 $\\sigma$.  These results are consistent with the prediction of the MSW-LMA solution of a transition in the solar \\nue\\ survival probability \\Pee\\ between the low energy vacuum-driven and the high-energy matter-enhanced solar neutrino oscillation regimes. ",
        "introduction": "\\label{sec:intro}  Solar \\bor-neutrino  spectroscopy  has been so far performed by the water \\che\\ detectors KamiokaNDE, SuperKamiokaNDE, and SNO~\\cite{Hir89,SKII08,SKI05,SNO07}.  The first two experiments used elastic $\\nu$-$e$ scattering for the detection of neutrinos, whereas SNO also exploited nuclear reaction channels on deuterium with heavy water as  target.  These experiments provided robust spectral measurements with $\\sim$5\\,MeV threshold or higher for scattered electrons; a recent SNO analysis reached a 3.5\\,MeV threshold~\\cite{SNO09}.  We report the first observation of solar \\bor-neutrinos with a liquid scintillator detector, performed by the Borexino experiment~\\cite{BXD08,BX09} via elastic $\\nu$-$e$ scattering.  Borexino is the first experiment to succeed in suppressing all major backgrounds, above the 2.614\\,MeV $\\gamma$ from the decay of \\tal,  to a rate below that of electron scatterings from solar neutrinos.  This allows to reduce  the energy threshold for scattered electrons  by \\bor\\ solar neutrinos to 3\\,MeV, the lowest ever reported for the electron scattering channel. To facilitate a comparison to the results of SuperKamiokaNDE \\cite{SKI05} and SNO  D$_2$O phase \\cite{SNO07}, we also report the measured \\bor\\ neutrino interaction rate with 5\\,MeV threshold.  Since Borexino also detected low energy solar \\ber\\ neutrinos~\\cite{BX07,BX08}, this is the first experiment where both branches of the solar \\pp-cycle have been measured simultaneously  in the same target.  The large mixing angle solution (LMA) of the MSW effect~\\cite{MSW} predicts a transition in the $\\nu_e$ survival probability from the vacuum oscillation regime at low energies to the matter dominated regime at high energies.  Results on solar \\ber\\ and \\bor\\ neutrinos from Borexino, combined with prediction on the absolute neutrino fluxes from the Standard Solar Model~\\cite{BS07,Pen08,Ser09}, confirm that our data are in agreement with the MSW-LMA prediction within 1$\\sigma$. ",
        "conclusions": ""
    },
    "0808/0808.0576_arXiv.txt": {
        "abstract": "To date, fully cosmological hydrodynamic disk simulations to redshift zero have only been undertaken with particle-based codes, such as {\\tt GADGET}, {\\tt Gasoline}, or {\\tt GCD+}.  In light of the (supposed) limitations of traditional implementations of smoothed particle hydrodynamics (SPH), or at the very least, their respective idiosyncrasies, it is important to explore complementary approaches to the SPH paradigm to galaxy formation. We present the first high-resolution cosmological disk simulations to redshift zero using an adaptive mesh refinement (AMR)-based hydrodynamical code, in this case, {\\tt RAMSES}.  We analyse the temporal and spatial evolution of the simulated stellar disks' vertical heating, velocity ellipsoids, stellar populations, vertical and radial abundance gradients (gas and stars), assembly/infall histories, warps/lopsideness, disk edges/truncations (gas and stars), ISM physics implementations, and compare and contrast these properties with our sample of cosmological SPH disks, generated with {\\tt GCD+}.  These preliminary results are the first in our long-term Galactic Archaeology Simulation program. ",
        "introduction": "The ability to form and evolve (correctly!) a disk galaxy with the aid of massively parallel computers and optimised algorithms remains an elusive challenge for astrophysicists.  Ameliorating the non-physical effects associated with overcooling, overmerging, angular momentum loss, and the capture of accurate phenomenological prescriptions for the sub-grid physics governing galaxy evolution (star formation, feedback, etc.) has been achieved through rapid advancements in both hardware and software algorithms, but their complete elimination has yet to be realised.  Fully self-consistent cosmological hydrodynamic simulations of Milky Way-like disk galaxies, with sufficient resolution ($\\simlt$500~pc) to decompose and analyse various galactic sub-components (eg. halo, bulge, and thin + thick disks) have only really appeared over the past $\\sim$5 years (Sommer-Larsen et~al. 2003; Abadi et~al. 2003; Governato et~al. 2004,2007; Robertson et~al. 2004; Bailin et~al. 2005; Okamoto et~al. 2005).  A common thread linking these studies is the use of a particle-based approach to representing and solving the equations of hydrodynamics - usually through the use of a smoothed particle hydrodynamics (SPH) code, such as {\\tt GADGET}, {\\tt Gasoline}, or {\\tt GCD+}. Where there is no disputing the impact that SPH has had on the field, it is important to be aware of both the strengths {\\it and} weaknesses of any specific approach - as O'Shea et~al. (2005) and Agertz et~al. (2007) have shown, both subtle {\\it and} overt differences can be introduced when employing a particle-based, as opposed to a mesh-based (or grid-based) approach (and \\it vice versa\\rm), when simulating galaxy formation and evolution.  To address these concerns, we have initiated a long-term Galactic Archaeology Simulation programme aimed at complementing the aforementioned particle-based studies (including our own) with a comprehensive suite of simulations generated with a grid-based N-body + hydrodynamical code employing adaptive mesh refinement (AMR) -  our software tool of choice has been {\\tt RAMSES} (Teyssier 2002).  \\it These simulations represent (to our knowledge) the first to be generated through to redshift zero, with a grid code, within a fully cosmological and hydrodynamic framework.\\rm\\footnote{The beautiful simulations of Ceverino \\& Klypin (2008), generated with the {\\tt ART} grid code were not (again, to our knowledge) run to redshift zero.}  In this contribution, we provide a brief summary of the methodology adopted, and highlight several {\\it preliminary} results associated with our analyses of the simulations' disk kinematics, chemistry, disk edges / truncations, and assembly / infall histories.  \\vspace{-4mm} ",
        "conclusions": "We have realised (to the best of our knowledge) the first fully self-consistent cosmological hydrodynamic disk simulations to $z$=0 with a mesh code; the resolution attained is 435~pc.  Several preliminary results include: \\begin{itemize} \\item the saturated vertical disk heating seen in semi-cosmological SPH simulations has not yet been clearly replicated in our cosmological simulations; \\item the neutral gas disks show ``edges'' at comparable column densities to those observed; ionised gas disks extend beyond the neutral gas, again in agreement with those observed; \\item the stellar surface brightness profiles show ``breaks'' in the exponential profiles, with associated increases (reddening) in the age (colour) of the stellar populations beyond the break, in agreement with observation; little evidence is seen for an associated break in the stellar surface density profile, also as inferred from observations; stars of the same age beyond and interior to the break do not appear to have the same metallicity, which may prove problematic for radial migration scenarios; \\item gas accretion is not smooth, but does appear to be more-of-less ``inside-out''; \\item the disk-halo ``circulation flux'' is $\\sim$10-50$\\times$ that of the ``infalling flux'' (again, consistent with the broad numbers associated with the Milky Way). \\end{itemize}  Beyond the analysis of the extant simulations, we have a number of planned enhancements, including full chemical evolution / tagging, a ten-fold increase in the number of simulations (to examine scaling relations, environmental dependencies, and assembly history variations), a range of ISM physics implementations (various polytropic equations of state, blast wave parametrisations), quantifying warp and lopsidedness statistics (Mapelli et~al. 2008), 2d IFU Lick index-style maps, dusty radiative transfer, high-velocity clouds, radial gas flows, and detailed SPH vs AMR comparisons with identical initial conditions.  \\vspace{-4mm}"
    },
    "0808/0808.2090_arXiv.txt": {
        "abstract": "We explore large-scale hydrodynamics of H \\Rmnum{2} regions for various self-similar shock flows of a polytropic gas cloud under self-gravity and with quasi-spherical symmetry. We formulate cloud dynamics by invoking specific entropy conservation along streamlines and obtain global self-similar ``champagne flows\" for a conventional polytropic gas with shocks as a subclass. Molecular cloud cores are ionized and heated to high temperatures after the onset of nuclear burning of a central protostar. We model subsequent evolutionary processes in several ways and construct possible self-similar shock flow solutions. We may neglect the mass and gravity of the central protostar. The ionization and heating of the surrounding medium drive outflows in the inner cloud core and a shock travels outwards, leading to the so-called ``champagne phase\" with an expanding outer cloud envelope. Complementarily, we also consider the expansion of a central cavity around the centre. As the inner cloud expands plausibly due to powerful stellar winds, a cavity (i.e., `void' or `bubble') can be created around the centre, and when the cavity becomes sufficiently large, one may neglect the gravity of the central protostar. We thus present self-similar shock solutions for ``champagne flows\" with an expanding central void. We compare our solutions with isothermal solutions and find that the generalization to the polytropic regime brings about significant differences of the gas dynamics, especially for cases of $n<1$, where $n$ is a key scaling index in the self-similar transformation. We also compare our global polytropic self-similar solutions with numerical simulations on the expansion of H \\Rmnum{2} regions. We further explore other possible dynamic evolutions of H \\Rmnum{2} regions after the initiation of nuclear burning of the central protostar, for example asymptotic inflows or contractions far from the cloud centre and the ongoing infall around a central protostar. In particular, it is possible to use the downstream free-fall solution with shocks to describe the dynamic evolution of H \\Rmnum{2} regions shortly after the nascence of the central protostar. We also give an analysis on the invariant form of self-similar polytropic flows by ignoring self-gravity. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.1570_arXiv.txt": {
        "abstract": "Superhorizon perturbations induce large-scale temperature anisotropies in the cosmic microwave background (CMB) via the Grishchuk-Zel'dovich effect.  We analyze the CMB temperature anisotropies generated by a single-mode adiabatic superhorizon perturbation.  We show that an adiabatic superhorizon perturbation in a \\lcdm universe does not generate a CMB temperature dipole, and we derive constraints to the amplitude and wavelength of a superhorizon potential perturbation from measurements of the CMB quadrupole and octupole.  We also consider constraints to a superhorizon fluctuation in the curvaton field, which was recently proposed as a source of the hemispherical power asymmetry in the CMB. ",
        "introduction": "The finite age of the Universe implies the existence of a cosmological particle horizon beyond which we cannot observe.  Inhomogeneities with wavelengths longer than the horizon are not completely invisible, however.  The generation of large-scale temperature fluctuations in the cosmic microwave background (CMB) by superhorizon perturbations is known as the Grishchuk-Zel'dovich effect \\cite{GZ78}.  Through this effect, measurements of the low-multipole moments of the CMB \\cite{COBE, WMAP1} place constraints on the amplitudes and wavelengths of superhorizon perturbations.  A well-known application of the Grishchuk-Zel'dovich effect uses CMB observations to place a lower bound on the size of the nearly homogeneous patch that contains the observable Universe. This bound was first derived for an Einstein-de Sitter universe \\cite{GZ78, Turner91}, and then for an open universe \\cite{KTF94, GLLW95}.  Most recently, an analysis of the WMAP first-year data \\cite{WMAP1} found that our nearly homogeneous patch of the Universe extends to 3900 times the cosmological horizon \\cite{CDF03}.  All of these analyses considered a statistically isotropic distribution of power in superhorizon perturbations and then asked how large the wavelength of order-unity perturbations needed to be in order to be consistent with the observed CMB anisotropies.  In this paper, we analyze the CMB anisotropies induced by a single superhorizon adiabatic perturbation mode rather than an isotropic distribution of superhorizon inhomogeneities.   A single-mode superhorizon perturbation to the gravitational potential would naively be expected to generate a dipolar CMB anisotropy with an amplitude comparable to the perturbation amplitude across the observable Universe.  This is not the case in an Einstein-de Sitter universe, however, because the intrinsic dipole in the CMB produced by the perturbation is exactly cancelled by the Doppler dipole induced by our peculiar motion \\cite{GZ78, Turner91, BL94}.  We show that the same cancellation occurs for an adiabatic superhorizon perturbation in a flat universe with a cosmological constant ($\\Lambda$), cold dark matter (CDM), and radiation.  The strongest constraints to the amplitude and wavelength of a single superhorizon mode therefore arise from measurements of the CMB quadrupole and octupole.  These constraints are less stringent than those derived for modes in a realization of a random-phase random field because it is possible to choose the phase of a single sinusoidal perturbation in such a way that there is no resulting quadrupole anisotropy.  Single-mode superhorizon perturbations have received attention recently \\cite{GHHC05, Gordon07, DDR07, EKC08} because they introduce a special direction in our Universe and could be responsible for observed deviations from statistical isotropy in the CMB \\cite{TOH03, OTZH04, LM05, Eriksen04, HBG04, Eriksen07, GE08}.    In particular, we investigated in a recent paper \\cite{EKC08} how a superhorizon perturbation during slow-roll inflation can generate an anomalous feature of the CMB: the fluctuation amplitude on large scales ($\\ell \\lsim 40$) is 10\\% larger on one side of the sky than on the other side \\cite{Eriksen04, HBG04, Eriksen07}.  We first considered a perturbation to the inflaton field, but we found that the perturbation required to generate the observed power asymmetry induces large-scale anisotropies in the CMB that are too large to be consistent with measurements of the CMB octupole.  We then considered a multi-field model of inflation in which a subdominant field, called the curvaton, is responsible for generating primordial perturbations \\cite{Mollerach90, LM97, LW02, MT01}.   We found that a superhorizon perturbation in the curvaton field can generate the observed power asymmetry without inducing prohibitively large CMB anisotropies.  We will use Ref.~\\cite{EKC08} as an example of how one may apply the CMB constraints to single-mode superhorizon perturbations derived here.  We begin in Section \\ref{sec:basics} by reviewing the Grishchuk-Zel'dovich effect for adiabatic perturbations.  In Section \\ref{sec:potential}, we derive the CMB anisotropy induced by a sinusoidal superhorizon perturbation in the gravitational potential, as would arise from a sinusoidal inflaton fluctuation.   We also show in Section \\ref{sec:potential} that a superhorizon adiabatic perturbation does not generate a large dipolar anisotropy in a \\lcdm universe because the leading-order intrinsic dipole anisotropy is cancelled by the anisotropy induced by the Doppler effect.  A sinusoidal curvaton fluctuation generates a potential perturbation that is not sinusoidal, and we derive the constraints to single-mode perturbations to the curvaton field in Section \\ref{sec:curvaton}.  We summarize our results in Section \\ref{sec:conclusions}.    Finally, an analytic demonstration of the dipole cancellation in a \\lcdm universe is presented in Appendix \\ref{app:cancel}, and the cancellation is shown to occur in flat universes containing a single fluid with an arbitrary constant equation of state in Appendix \\ref{app:cancel2}. ",
        "conclusions": "\\label{sec:conclusions} Superhorizon perturbations generate large-scale anisotropies in the CMB through the Grishchuk-Zel'dovich effect \\cite{GZ78}.   In this paper, we have derived the constraints to single-mode adiabatic superhorizon perturbations that arise from measurements of the CMB quadrupole and octupole.  These constraints differ from those previously derived for an isotropic distribution of superhorizon perturbations \\cite{GZ78, Turner91, KTF94,  GLLW95, CDF03} because the CMB anisotropies generated by a single-mode perturbation depend on the perturbation's phase.  We started by considering a sinusoidal superhorizon gravitational potential perturbation with wavenumber $k\\ll H_0$.  Since the leading-order term in the potential perturbation is proportional to $(\\vec{k}\\cdot\\vec{x})$, it would be expected to generate a dipolar anisotropy of comparable amplitude in the CMB through the Sachs-Wolfe effect.  However, the superhorizon perturbation also gives us a velocity with respect to the CMB, and the resulting Doppler dipole exactly cancels the leading-order intrinsic anisotropy generated by the SW and ISW effects, provided that the perturbation is adiabatic.  This cancellation was known to occur in an Einstein-de Sitter universe \\cite{GZ78, Turner91, BL94}, but we found that it also applies to flat \\lcdm universes with and without radiation, as well as in more exotic flat cosmological models.  Due to this cancellation of the CMB dipole, the leading-order constraints on adiabatic superhorizon fluctuations arise from measurements of the CMB quadrupole and octupole.  If the potential perturbation is sinusoidal, as would be created by a sinusoidal fluctuation in the inflaton, then putting ourselves at the node of the sine wave maximizes the difference in potential across the Universe while also eliminating the induced quadrupole anisotropy in the CMB.  In this case, the CMB octupole provides the strongest constraint on the amplitude of the superhorizon perturbation: $\\Delta \\Psi \\lsim 0.095$, where $\\Delta \\Psi$ is the variation of the potential $\\Psi$ across the surface of last scattering.  A fluctuation in a field that contains only a small fraction of the energy density of the Universe generates a smaller potential perturbation and, consequently, smaller CMB anisotropies.  We consider a multi-field model of inflation in which a subdominant curvaton field generates primordial perturbations \\cite{Mollerach90, LM97, LW02, MT01}.   For a given superhorizon fluctuation in the curvaton field, the measured values of the CMB quadrupole and octupole place upper bounds on the fraction of the total energy density contained in the curvaton field prior to its decay.  Since a sinusoidal perturbation in the curvaton field generates a potential perturbation that is not sinusoidal, there is no value for the phase of the curvaton fluctuation that eliminates the induced CMB quadrupole for any superhorizon curvaton fluctuation.  However, once the amplitude of the curvaton perturbation is specified, it is possible to choose a phase for which the induced CMB quadrupole vanishes.  In this case, measurements of the CMB octupole still place an upper bound on the curvaton energy density, but this bound is significantly weaker than the bound from the CMB quadrupole that applies to curvaton fluctuations with different phases.  Superhorizon perturbations have generated interest recently because they are a simple way to introduce a preferred direction in the Universe and may generate the deviations from statistical isotropy that have been observed in the CMB.  In particular, in Ref. \\cite{EKC08}, we showed that a superhorizon perturbation to an inflationary field can generate the hemispherical power asymmetry found in the WMAP data \\cite{Eriksen04, HBG04, Eriksen07}.   In this paper, we have demonstrated how the CMB constrains such superhorizon perturbations: the octupole constraint on $\\Delta \\Psi$ is sufficient to rule out an inflaton perturbation as the source of the observed power asymmetry, but it is possible to generate the observed power asymmetry with a superhorizon curvaton perturbation.  These constraints may also be applied to other scenarios that invoke superhorizon perturbations.  For instance, order-unity superhorizon fluctuations in the mean value of the curvaton may be a generic feature of the curvaton model \\cite{LM06}."
    },
    "0808/0808.0007_arXiv.txt": {
        "abstract": "Cluster galaxies moving through the intracluster medium (ICM) are expected to lose some of their interstellar medium (ISM) through ram pressure stripping and related ISM-ICM interactions.  Using high-resolution cosmological simulations of a large galaxy cluster including star formation, we show that the ram pressure a galaxy experiences at a fixed distance from the cluster center can vary by well over an order of magnitude  We find that this variation in ram pressure is due in almost equal parts to variation in the ICM density and in the relative velocity between the galaxy and the ICM.  We also find that the ICM and galaxy velocities are weakly correlated for in-falling galaxies. ",
        "introduction": "X-ray observations of clusters have shown that substructure in the intracluster medium (ICM) is common (e.g. Mohr, Mathiesen, \\& Evrard 1999; Schuecker et al 2001).  In a sample of 470 clusters, Schuecker et al. (2001) measure substructure in more than 50{\\%} of their sample.  Detailed examinations of nearby clusters like Perseus and Virgo have discovered substructure and/or asymmetry in both the temperature and density profiles of these clusters (e.g. Bohringer et al. 1994; Shibata et al. 2001; Churazov et al. 2003; Dupke \\& Bregman 2001; Furusho et al 2001). Even Coma, considered a relaxed cluster, has ICM irregularities (White, Briel \\& Henry 1993).  The importance of substructure on cluster mass measurements has been examined (Mohr, Mathiesen \\& Evrard 1999; Bohringer et al 2000), which in turn affects the use of cluster measurements as cosmological constraints (Jeltema et al. 2005; Nagai et al. 2007).  However, the importance of substructure in the ICM is rarely considered when studying ram pressure stripped galaxies.  Common assumptions are that the ICM is static, has a smooth density profile, and is only dense enough very near the center of a cluster to affect galaxies.  Treu et al. (2003), in their evaluation of possible environmental evolutionary mechanisms in Cl 0024 + 16, assume that ram pressure is only effective to 0.6 virial radii.  Solanes et al. (2001) find HI deficiency in galaxies out to two Abell radii, but only discuss the possibility that these galaxies are on highly radial orbits that have already carried them through the cluster center.  Previous simulations studying galaxy evolution in clusters use a static, smooth ICM profile when studying the orbits of galaxies in clusters (e.g. Vollmer 2001; Roediger \\& Br{\\\"u}ggen 2007; J{\\'a}chym et al. 2007).  These authors use different galaxy orbits in order to sample a variety of galaxy velocities at a fixed ICM density.  Although the use of simple assumptions is widespread, there is at least one possible case in which ICM substructure had to be invoked to explain observations of the Virgo galaxy NGC 4522, a galaxy with a truncated gas disk (Kenney et al 2004; Vollmer et al. 2004; Vollmer et al. 2006).  NGC 4522 is located at a projected distance of 1 Mpc from the center of the Virgo cluster, and assuming a static ICM with standard density values, the ram pressure is not strong enough to cause the observed truncation.  Thus, the authors propose that the nearby ICM is either moving relative to the galaxy or overdense.  In a recent paper studying the environmentally-driven evolution of galaxies in clusters using a detailed cosmological simulation (Tonnesen, Bryan \\& van Gorkom 2007), we examined the evolution of cool gas (i.e. ISM) in galaxies within and around the cluster, demonstrating that most gas loss from galaxies was due to ISM-ICM interactions (i.e. ram pressure and related processes), rather than galaxy-galaxy interactions or cluster tidal effects.  We also found that ram pressure stripping occurs out to the virial radius of the cluster (measured using $r_{200}$).  In this paper, we examine this result more closely and show that the ram pressure a galaxy experiences varies substantially, even at fixed distance from the cluster center.   As we will see, this arises both from the density and velocity substructure of the ICM.   First, we briefly introduce our code in \\S \\ref{sec:sim} and explain how we measure ram pressure in our simulation (\\S \\ref {sec:sample} and \\S \\ref{sec:rpsample}).  We then present our results:  a comparison of the standard deviations of ram pressure, ICM density, and velocity difference squared (\\S \\ref{sec:density}), followed by a more detailed look at the velocity of the ICM (\\S \\ref{sec:velocity}). ",
        "conclusions": "\\label{sec:discussion}  In this paper we have presented a detailed examination of the intracluster medium with which a galaxy interacts as it falls into a simulated galaxy cluster.  We find that substructure in the ICM is more important in varying ram pressure than is often assumed and used when modeling ram pressure stripping.  Specifically, we highlight three main points:  \\begin{enumerate}  \\item  In our simulated cluster we measure a range of ram pressure values for any given radius in the cluster.  This ranges from an order of magnitude at 1 Mpc, to two orders of magnitude at the virial radius (1.8 Mpc), to even larger deviations further from the cD.  Therefore, ram pressure can be effective at larger radii wherever there is an overdensity.  \\item  The scatter in ram pressure at different distances from the cD is due equally to the variation in the ICM density and the relative ICM-galaxy velocity ($v_{\\Delta}^2$) within the virial radius.  This is true even when considering only higher ram pressure values.  In fact, the normalized standard deviation in galaxy velocity is smaller than that of $v_{\\Delta}^2$.  It is therefore not only the variety of orbital velocities that causes different values of ram pressure at fixed cluster radius, but also the density and velocity structure of the ICM.  Further from the cD, ICM density variations dominate those of $v_{\\Delta}^2$.  \\item  The ICM velocity is correlated with galaxy velocity, resulting in a smaller $v_{\\Delta}$ than $v_{galaxy}$.  This indicates that the ICM tends to move with in-falling galaxies, which then experience somewhat less ram-pressure than one would expect from a static ICM (although this is less true for high-velocity galaxies that are likely near the cluster center).  \\end{enumerate}  We emphasize that although we determine the ICM properties from a simulation, it is well-known in the X-ray cluster field that ICM substructure in density is common and in good agreement with simulations (Mohr, Mathiesen, \\& Evrard 1999; Jeltema et al. 2005; Nagai et al. 2007).  Because ram pressure stripping is a fast process, even an overdensity with a relatively small extent can strip a galaxy that might otherwise retain its gas, or strip a galaxy more than predicted by its cluster position.  Our results should galvanize the community currently studying galaxy evolution in clusters to look more closely at the intracluster medium."
    },
    "0808/0808.2975_arXiv.txt": {
        "abstract": "The presence of atomic gas mixed with molecular species in a ``molecular'' cloud may significantly affect its chemistry, the excitation of some species, and can serve as probe of the cloud's evolution. Cold neutral atomic hydrogen (HI) in molecular clouds is revealed by its self absorption of background galactic HI 21-cm emission. The properties of this gas can be investigated quantitatively through observation of HI Narrow Self-Absorption (HINSA). In this paper, we present a new technique for  measuring atomic gas physical parameters from HINSA observations that  utilizes molecular tracers to guide the HINSA extraction. This technique offers a significant  improvement in the precision with which HI column densities can be determined over previous methods, and it opens several new avenues of study of relevance to the field of star formation. ",
        "introduction": "Star formation occurs in molecular clouds which are thought to have evolved from diffuse atomic hydrogen (HI) regions to form dense, cold, well--shielded regions composed primarily of molecular Hydrogen (\\H2). Our quantitative understanding of cloud evolution and specifically the conversion process (from HI to \\H2 gas) has been hindered by our inability to  measure confidently the HI abundance in evolving clouds. In this paper we present a new technique for measuring HI column densities in dark clouds offering a significant improvement over previous methods. In the interests of brevity this paper includes the results from only a few clouds on which this technique has been applied in order to demonstrate the technique. The results of a much larger survey of observational data and analysis are to follow in a subsequent publication.     The ability to determine accurately the HI component of molecular clouds could have a variety of benefits. Measurement of  HI/\\H2 ratios in clouds, used in conjunction with astrochemical models, allows us to determine the chemical ages of individual clouds or entire molecular complexes (e.g. Taurus, Perseus, etc.). This will greatly expand our understanding by constraining star formation models thus yielding insights into the collapse process and the interplay of magnetic fields, ambipolar diffusion, turblence, and various potential sources of cloud support. By studying age distributions in large scale regions we can learn about the processes that may trigger the collapse of large complexes. In contrast to previous methods,  our technique allows us to determine the HI/\\H2 ratios for individual velocity components within a cloud thus yielding unique information about cloud kinematics. Further, the technique allows for the absolute measurement of quantities such as HI column density, in contrast to many previous studies which were limited to comparative measurements.     HI is the dominant constituent of the diffuse ISM and its 21cm emission line is prevalent everywhere throughout the sky, especially near the galactic plane. Typical HI emission spectra are composed of numerous superimposed velocity components. The emission linewidths of the overall features are typically on the order of a few 10s of \\kms. They include velocity variations owing to galactic rotation as well as very significant peculiar velocities resulting from localized phenomena. It is rare to find a molecular region for which one can confidently claim that the HI emission observed along the line of sight to the cloud is associated with the cloud, and is not due to background or foreground sources. Owing to such velocity crowding it is difficult or impossible to disentangle HI emission originating from a particular cloud from the background galactic HI emission.    The situation for HI absorption is different. HI within a galactic cloud of any type can absorb the continuum emission from distant (galactic or extragalactic) radio sources. Because the HI optical depth varies inversely with the cloud temperature, the absorption by galactic HI is stronger in cold, galactic HI clouds. The integrated optical depth of all the clouds along the line of sight through the galactic disk can exceed unity as demonstrated in \\cite{Kolpak}, and \\cite{Garwood}.    These cold, interstellar HI clouds may also be identified through their absorption of warm background HI emission originating within the galaxy. Because the emission being absorbed by cold HI clouds in this case is galactic HI emission, the resultant spectral absorption features are refered to as HI self-absorption (HISA). There have been many surveys with HISA detections over the years including \\cite{garzoli}, \\cite{Heiles}, \\cite{Knapp}, \\cite{Heiles2}, \\cite{Wilson}, \\cite{McCutcheon}, \\cite{Myers}, \\cite{Bowers}, \\cite{Batrla}, \\cite{Shuter}, \\cite{vanderWerf}, \\cite{Feldt}, \\cite{Montgomery}, \\cite{Gibson}, and \\cite{Kavars03}, among others. It is important to emphasize that the term HISA refers to an observable spectral absorption feature rather than being a description of a specific physical process.    Molecular spectral emission lines provide an independent view of that subset of interstellar clouds that are cold and composed primarily of molecular species. These ``molecular clouds'' are expected to maintain a residual abundance of atomic hydrogen if for no other reason than the cosmic ray disassociation of \\H2. This applies even when the chemical evolution of the cloud has reached equilibrium \\citep{Solomon1971, HINSA2}. With the residual HI co--existing throughout the cloud with molecular species, the observed spatial and velocity (kinematic) structure of the molecular cloud will be similar whether observed in molecular emission or HI self-absorption lines. Thus, we expect to observe HI self-absorption features along the lines of sight of molecular clouds which share the spatial distribution and kinematics (non--thermal line width) of molecular emission lines. Such localized association between molecular emission and HI self-absorption is observed for many nearby clouds as reported by \\cite{HINSA1} and \\cite{HINSA2}. The specific case in which the HI absorption features observed in the direction of a molecular cloud share the spatial and kinematic structure seen in the molecular lines is called HI Narrow Self-Absorption (HINSA) as defined in \\cite{HINSA1}. The term {\\it Narrow} arises from the typically small nonthermal linewidths (on the order of 0.1 \\kms) of HINSA features, very similar to the similarly small nonthermal linewidths of molecular tracers along the same lines of sight. HINSA can be considered to be a subset of HISA, but it is a subset derived from an understanding of a specific, observable physical phenomenon in molecular clouds.    While both are simply acronyms for spectral absorption features, HISA can be caused by a variety of different conditions and processes, but HI Narrow Self-Absorption (HINSA) is that subset of HISA in which the atomic HI absorption correlates well with molecular emission of certain tracers (most notably $^{13}$CO) in sky position, central velocity, and nonthermal line width. Based on our current understanding of cold molecular clouds, the most satisfactory picture is that HINSA features are a result of HI gas located within these cold, dense, well--shielded regions. Some early examples of HINSA studies that predate the use of this term are those of \\cite{Wilson}, \\cite{vanderWerf}, and \\cite{Jackson}. The technique which we describe in this paper pertains only to the extraction of HI data from HINSA features.     The general picture which emerges, as found in \\cite{HINSA1} and \\cite{HINSA2}, is that the HI gas located within cold, quiescent cores of dark clouds produces HINSA absorption features. The well--defined center velocities and narrow line widths allow us to separate the HI gas associated with individual clouds from the galactic background. However, the complexity of the background emission spectra that are frequently encountered makes extracting accurate data (especially in terms of obtaining the cold HI column density) from the absorption features difficult. Several methods (discussed in \\S \\ref{technique}) have been used previously, but all are recognized to  introduce significant uncertainties in the results. We here present a new technique that aims to improve the situation by using the properties of molecular emission to characterize the region producing the HINSA features, and then employes the HINSA spectral features to derive HI column densities.    In \\S \\ref{real} we present selected data and show the results of applying the new technique to them, and in \\S \\ref{previous} we contrast this with previous methods for analyzing HINSA data. In \\S \\ref{technique} and \\S \\ref{beamsizes} we describe the technique and the combination of molecular data with the HINSA spectra. In \\S \\ref{examples} and \\S \\ref{ambiguity} we verify the validity of our technique using simulated data and also examine its limitations. ",
        "conclusions": "\\label{summary} The significance of accurately measuring the HI content of dark clouds and star forming regions has long been recognized. HINSA features have been shown as a promising method to achieve this goal, but the difficulty in confidently disentangling HINSA from the galactic background emission has limited research efforts in this field. The technique described here builds upon previous work to provide new opportunities. By utilizing a second derivative representation in which HINSA becomes the dominant feature in the spectrum, and using information from associated molecular tracers to constrain our fits, we are able to obtain HI column densities with greater confidence than previously possible. This technique enables us to study individual velocity components within a cloud. Several uncertainties and sources of error still remain, most notably the errors inherent in temperature determination through $^{12}$CO, and the precise relations between the properties of cold HI and molecular tracers such as $^{13}$CO. As shown in \\S \\ref{examples} the purely statistical errors derived from our constructed data were only a few percent in simple cases which represent the majority of observed spectra.  The nonlinearity error may be as large as 50 percent in the extreme cases.     While the scope of this paper is limited to a demonstration of the new technique, the improved confidence in the results have the potential to yield significant scientific advances in the field of molecular cloud and star formation studies which are deferred to a subsequent publication. The two astronomical sources discussed briefly here for demonstration purposes are part of a much larger survey of over 30 dense cloud cores. In conjunction with numerical modeling these data provide us with the chemical ages of the clouds and individual components therein, thus providing a significant constraint on theoretical models of star formation. Large maps have been collected with the Arecibo L-Band Feed Array (ALFA) of the Taurus and Perseus star forming regions. These regions show plentiful HINSA features whose analysis may shed light on the dynamics of such complexes including the processes which may have triggered their formation."
    },
    "0808/0808.0141_arXiv.txt": {
        "abstract": "We present new spectroscopic data for twenty six stars in the recently-discovered Canes Venatici~I (CVnI) dwarf spheroidal galaxy. We use these data to investigate the recent claim of the presence of two dynamically inconsistent stellar populations in this system~\\citep{Ibata2006}. We do not find evidence for kinematically distinct populations in our sample and we are able to obtain a mass estimate for CVnI that is consistent with all available data, including previously published data. We discuss possible differences between our sample and the earlier data set and study the general detectability of sub-populations in small kinematic samples. We conclude that in the absence of supporting observational evidence (for example, metallicity gradients), sub-populations in small kinematic samples (typically fewer than 100 stars) should be treated with extreme caution, as their detection depends on multiple parameters and rarely produces a signal at the 3$\\sigma$ confidence level. It is therefore essential to determine explicitly the statistical significance of any suggested sub-population. ",
        "introduction": "\\label{sec:intro}  It is now widely accepted that the dwarf spheroidal (dSph) satellite galaxies of the Milky Way and Andromeda are the most dark matter dominated stellar systems known in the Universe~\\citep[e.g.][]{Mateo1998}. Over the past two decades, a significant amount of observational work has focussed on quantifying both the amount of dark matter in these systems, and its spatial distribution~\\citep[e.g.][]{Gilmore2007,Walker2007}.  Although recent numerical simulations have shown that many of the dSphs may not be immune to tidal disturbance by the Milky Way~\\citep[e.g.][]{Munoz2008,Lokas2008}, their observed properties still require the presence of massive dark matter haloes which protect them against complete tidal disruption. The dSphs thus provide us with nearby laboratories in which to test dark matter theories.  Given that dSphs occupy the low luminosity end of the galaxy luminosity function, their star formation histories provide useful insights into the star formation process. Analyses of spatial variations in colour-magnitude diagram morphology provided early evidence of population gradients in a number of dSphs~\\citep[e.g.][]{Harbeck2001}. More recently, evidence of metallicity gradients has been found using spectroscopic estimates of [Fe/H]~\\citep[e.g.][]{Tolstoy2004,Koch2006,Battaglia2006}. In at least one case, that of the Sculptor dSph, the metal-rich and metal-poor populations have significantly different spatial distributions and kinematics~\\citep{Tolstoy2004,Battaglia2008}. Although little evidence of similar features has been found in other dSphs~\\citep[e.g.][]{Koch2006, Koch2007a, Koch2007b}, the presence of dynamically distinct stellar populations within dSphs, as well as the complex interplay between the dynamical, spatial and chemical properties of their stars, is of great interest as it has implications for star formation and galaxy evolution.  It is, however, important to note that although the hierarchical build-up of structure in the standard $\\Lambda$-Cold Dark Matter ($\\rm{\\Lambda\\,CDM}$) paradigm implies that satellite galaxies contribute significantly to the stellar haloes of their hosts, detailed abundance studies of stars in the more luminous dSphs have demonstrated that their properties are significantly different from those of the Milky Way halo~\\citep[e.g.][]{Shetrone2001,Helmi2006}. Among the significant differences between the halo and the dSphs, the more important chemical differences are in the alpha-elements~\\citep{Unavane1996,Venn2004}. The observed gradients in the heavy element distributions are reproduced by the models of supernova feedback in dSphs developed by~\\cite{Marcolini2008}. Thus, it appears that the primordial dwarf satellites, which were disrupted to form the Milky Way halo, had stellar populations distinct from those seen in the present-day dSphs~\\citep{Robertson2005, Font2006}.  Given their high estimated mass-to-light ratios, the observed dSphs are usually identified with the large population of sub-haloes which are observed to surround Milky Way-sized haloes in cosmological simulations assuming a standard $\\rm{\\Lambda\\,CDM}$ universe. However, it was noted early on that the number of dSphs around the Milky Way was much lower than the expected number of satellite dark matter haloes~\\citep[e.g.][]{Moore1999}. A number of possible explanations for the apparent lack of Milky Way satellites have been presented in the literature~\\citep[e.g.][]{Stoehr2002, Diemand2005, Moore2006, Strigari2007, Simon2007, Bovill2008}. All these models are based on the reasonable postulate that out of the full population of substructures around the Milky Way, the observed dSphs are merely the particular subset which (for reasons of mass, orbit, formation epoch, re-ionisation, etc.) were able to capture gas, form stars and survive any subsequent tidal interactions with the Milky Way.  In addition, the ratio between the predicted and observed numbers of dwarf galaxies has decreased significantly in the past few years due to the discovery of nine new Milky Way dSph satellites~\\citep{Willman2005,Zucker2006a,Zucker2006b,Belokurov2006,Belokurov2007,Walsh2007} in the data from the Sloan Digital Sky Survey~\\citep[SDSS;][]{York2000}.  Since the SDSS covers only about one fifth of the sky, it is thus likely that the total number of satellites surrounding the Milky Way may be at least a factor of five larger than previously thought, although the extrapolation from the SDSS survey to the whole sky requires careful analysis~\\citep[see e.g.][]{Tollerud2008} . In order to compare the properties of the newly discovered satellites with those of sub-haloes in cosmological simulations, as well as to confirm their nature as true satellite galaxies of the Milky Way, as opposed to star clusters or disrupted remnants, spectroscopic observations of their member stars are essential in order to estimate dynamical masses from the observed stellar kinematics. The extremely low luminosities of these objects ~\\citep[in some cases as low as $10^3$L$_\\odot$: ][]{Martin2008b}, present significant observational challenges as the kinematic data sets are small, making it difficult to obtain statistically significant results.  The Canes Venatici~I (CVnI) dSph is the brightest of the newly discovered population of very faint SDSS dSphs~\\citep{Zucker2006a}. ~\\cite{Ibata2006} presented spectra for a sample of CVnI member stars obtained using the DEIMOS spectrograph mounted on the Keck telescope. They identified two kinematically distinct stellar populations in this data set: an extended metal-poor population with high velocity dispersion and a centrally-concentrated metal-rich population with a dispersion of almost zero. Their analysis of the mass of CVnI suggested that the two populations might not be in equilibrium as the mass profiles obtained based on the individual populations were inconsistent with each other.  However, a subsequent study of CVnI by ~\\cite{Simon2007}, using a larger sample of Keck spectra, did not reproduce this bimodality.  An important outstanding question is whether the ultra-faint dSphs represent the low-luminosity tail of the dSph population, or are instead the brightest members of a population of hitherto unknown faint stellar systems, distinct from both dSphs and star clusters. The presence of multiple, distinct kinematic populations in a low-luminosity dSph would set it apart from the majority of low-luminosity star clusters. In addition, the presence of a spread in the stellar abundances would suggest an association with the brighter dwarf galaxies and would also be interesting in terms of its implications for star formation. It is thus important to determine whether the sub-population identified by ~\\cite{Ibata2006} in CVnI is real. One goal of our study was to shed some light on this issue by using spectra obtained with a different spectrograph to those in the previous two studies of CVnI. In addition, we wanted to investigate the extent to which sub-populations can be reliably detected in the very small kinematic data sets which are observable for the ultra-faint dSphs.  In addition to their potential importance for probing the star formation histories of dSphs, kinematic substructures can be used to test another key feature of the hierarchical structure formation paradigm. The fact that dark matter clustering occurs on all scales means that the dSph satellites of the Milky Way are likely to be in the process of accreting their own population of smaller satellites. Although these substructures may not have been able to form their own stars, they may be able to acquire stars from their host dSph. They would then be detectable as localised populations with mean velocity and/or velocity dispersion distinct from that of the dSph. Populations with these properties have, in fact, been detected in the Ursa Minor and Sextans dSphs~\\citep{Kleyna2003,Walker2006}. Once a dSph halo begins to fall into the Milky Way, it will cease to accrete new satellites as any nearby substructures will rapidly be removed by the tidal field of the Milky Way and the high relative velocities in the Milky Way halo will preclude the capture of new satellites. Due to the short internal dynamical timescales in dSphs (typically a few hundred Myr), any remaining internal substructures will subsequently be destroyed on timescales of at most a few Gyrs if dSph haloes are cusped, although they can survive much longer if their haloes are cored~\\citep{Kleyna2003}. In the standard cusped-halo picture, only those satellites which have been interacting with the Milky Way for less than a few internal dynamical times, either because they are currently passing the Milky Way for the first time as may be the case for the Leo~I dSph~\\citep[]{Mateo2008} or the Magellanic Clouds ~\\citep[]{Kallivayalil2006, Besla2007,Piatek2008} or because their crossing times are larger~\\citep[e.g. the Magellanic Clouds:][]{vanderMarel2002}, would be expected to exhibit localised kinematic substructure. If localised substructures were found to be common in dSphs, this could be difficult to reconcile with a picture in which dSphs occupy cusped haloes. Given that the level of substructure above a given mass fraction is a function of halo mass \\citep{Gao2004}, the expected numbers of sub-haloes per dSph requires further investigation by means of cosmological simulations. However, the importance of comparing the level of substructure in dSphs with the results of numerical simulations adds further motivation to our goal of establishing the level of confidence with which sub-populations can be detected in small data sets.  The outline of the paper is as follows. In Section~\\ref{sec:cvn}, we present a new kinematic data set for stars in CVnI, based on spectra obtained with the Gemini telescope, and calculate a mass estimate for the galaxy from these data. In Section~\\ref{sec:pop}, we look for kinematic sub-populations in our data, and compare our findings with those of ~\\cite{Ibata2006}. Section~\\ref{sec:det} discusses the general detectability of sub-populations in small kinematic data sets. Finally, in Section~\\ref{sec:conc} we draw some general conclusions and suggest possible differences between the two data sets for CVnI that we have compared. ",
        "conclusions": "\\label{sec:conc}  In this paper, we have presented a new data set of velocities and metallicities for the Canes Venatici I (CVnI) dSph based on spectra taken with the GMOS-North spectrograph. A maximum likelihood fit to the velocity distribution yields a mean velocity of $v= 25.8\\pm0.3$\\,km\\,s$^{-1}$ and a dispersion of $\\sigma =7.9^{+1.3}_{-1.1}$\\,km\\,s$^{-1}$. Assuming a constant, isotropic velocity dispersion and a Plummer profile for the mass distribution, we find a mass of $4.4^{+1.6}_{-1.1}\\times10^7 M_\\odot$ in the volume where our tracer stars are located. Although this value is larger than the value $2.7\\pm 0.4\\times 10^7 M_\\odot$ calculated by ~\\cite{Simon2007},this is most likely due to the assumptions made for our models and the distribution of our particular subsets of stars.   One of the original aims of our study was to investigate the claimed multiple stellar populations in CVnI. As we discussed above, the two previous studies by ~\\cite{Ibata2006} and ~\\cite{Simon2007} did not agree on the existence of a cold sub-population in CVnI. The two populations found in the former study were puzzling as they led to two different mass estimates. The authors suggested that this might indicate that the system had recently accreted a younger population and was not yet in equilibrium.  In this paper we looked for evidence of multiple populations in our data under the assumption that each population was Gaussian. Based on this analysis, we concluded that there was no reason to suspect the presence of a second population in our data. We also applied our analysis to the \\cite{Ibata2006} data where we found evidence of a statistically significant sub-population with a dispersion of $\\sigma =0.6$\\,km\\,s$^{-1}$ (compared to $\\sigma =13.6$\\,km\\,s$^{-1}$ for the main population).  Our analysis suggests that there is a qualitative difference between our data and those of \\cite{Ibata2006}. Although further data would be necessary to resolve this issue, we note that the spatial distributions of these two data sets are different, which could potentially account for the differences in the detected populations. However, our central field is centred close to the blue/young star population which \\cite{Martin2008a} find in their photometry from the Large Binocular Telescope, and which they identify with the cold population of \\cite{Ibata2006}. The exact fraction of stars in each population found by \\cite{Martin2008a} is currently unclear, however, and so it is possible that we have not picked up any stars associated with the cold population.  We have also carried out a study of the detectability of sub-populations in small kinematic data sets.  Under the assumption of Gaussian populations, we studied the effects of four parameters. We obtained confidence limits for the detection of sub-populations in samples with different numbers of stars, different population ratios and velocity dispersions. We found that reasonable errors on the observed velocities do not affect the detectability of the sub-populations. For a given sample size, our ability to detect two populations increased as the ratio of their dispersions $\\sigma_{1}/\\sigma_{2}$ increased. However, even for large $\\sigma_{1}/\\sigma_{2}$ and equal population size, a sample of 30 stars yielded a $3\\sigma$ detection in only $\\sim35$ per cent of cases. As expected, for larger sample sizes, this detection rate was significantly higher. We also showed that a cold population needs to constitute a larger fraction of the total sample than is required to detect a hot sub-population. This suggests that the robust detection of the sub-populations associated with any surviving sub-haloes within a dSph would require samples of many hundreds of velocities. In this case, localised substructures could be detected by windowing the data, provided that a window whose spatial size coincided with plausible sub-halo scales would contain a sample of at least 100 stars. As such data sets are now becoming available for many of the larger dSphs, this test may soon be feasible. We note that the claim of multiple global populations in Sculptor~\\citep[]{Tolstoy2004} was based on a large data set and is therefore still robust.  Finally, we note that all our significance tests were based on the assumption of Gaussian populations, which was the case for all our Monte Carlo samples. However, for real data, the true distributions will not be known, and are not necessarily well-approximated by Gaussians. It is therefore difficult in a real case to assign a robust statistical significance to a particular detection of a sub-population.  As we have shown, for small data sets, many Monte Carlo realisations do not yield significant detections of the sub-populations. In the absence of a robust estimate of the confidence level of a particular detection, or additional, independent evidence of the presence of multiple populations, we conclude that one should exercise great caution in decomposing data sets of fewer than $100$ stars into multiple populations."
    },
    "0808/0808.0188_arXiv.txt": {
        "abstract": "{In recent years, increasing evidence for chemical complexity and multiple stellar populations in massive globular clusters (GCs) has emerged, including extreme horizontal branches (EHBs) and UV excess.} % {Our goal is to improve our understanding of UV excess in compact stellar systems, covering the regime of both ultra-compact dwarf galaxies (UCDs) and massive GCs.} {We use deep archival GALEX data of the central Fornax cluster to measure NUV and FUV magnitudes of UCDs and massive GCs. } {We obtain NUV photometry for a sample of 35 compact objects that cover a range $-13.5<M_V<-10$ mag. Of those, 21 objects also have FUV photometry.  Roughly half of the sources fall into the UCD luminosity regime ($M_V\\lesssim-$11 mag). We find that seven out of 17 massive Fornax GCs exhibit a NUV excess with respect to expectations from stellar population models, both for models with canonical and with enhanced Helium abundance. This suggests that not only He-enrichment has contributed to forming the EHB population of these GCs.  The GCs extend to stronger UV excess than GCs in M31 and massive GCs in M87, at the 97\\% confidence level. Most of the UCDs with FUV photometry also show evidence for UV excess, but their UV colours can be matched by isochrones with enhanced Helium abundances and old ages 12-14 Gyrs. We find that Fornax compact objects with X-ray emission detected from Chandra images are almost disjunct in colour from compact objects with GALEX UV detection, with only one X-ray source among the 35 compact objects.  However, since this source is one of the three most UV bright GCs, we cannot exclude that the physical processes causing X-ray emission also contribute to some of the observed UV excess.} {} ",
        "introduction": "\\label{introduction} In recent years, increasing evidence for chemical complexity and multiple stellar populations in several massive Galactic globular clusters (GCs) has emerged (e.g. Gratton et al. 2004, Bedin et al. 2004, Piotto et al. 2007, Milone et al. 2008). For some globular clusters (e.g. $\\omega$\\,Centauri and NGC 2808) there is evidence that these multiple populations may be due to an enhancement in Helium content (up to $Y\\simeq 0.4$), as deduced from the presence of extremely blue horizontal branches (EHBs) and multiple main sequences (D'Antona et al. 2005, Lee et al. 2005).   The link between occurence of EHBs and He-enhancement is, however, under debate. EHB stars may also originate directly from fast rotating stars with an enhanced mass loss during the red giant branch (RGB) evolution. Those stars leave the RGB before the He flash towards the (He-core) white dwarf cooling curve where they ignite Helium and end up on the EHB (Castellani \\& Castellani 1993, D'Cruz et al. 1996, Brown et al. 2001). This is sometimes called the 'hot-flasher' scenario. A significant fraction of the EHB stars in $\\omega$\\,Centauri seem to have followed this evolutionary path (Moehler et al. 2007).   It has been shown that EHB stars contribute most of the light in the UV bands (e.g.  NGC 2808: Dieball et al. 2005). The presence of an EHB in extragalactic - unresolved - GCs reveals itself in the integrated light by a UV-excess compared to GCs with a `normal' horizontal branch (HB). For example, Rey et al. (2007) find three metal-rich ([Fe/H]$>-$1) GC candidates in M31 with significant FUV flux which are thought to be analogs of two peculiar Galactic GCs, NGC 6388 and NGC 6441 (Yoon et al. 2008).  Sohn et al. (2006) and Kaviraj et al. (2007) analysed the UV properties of massive globular clusters associated with M87 in the Virgo cluster, and found that many of them show a UV-excess with respect to canonical stellar population models. These findings support the idea that EHBs may be a common feature to the most massive compact stellar systems.  In this Research Note we focus on the UV properties of compact stellar systems in the Fornax cluster. In contrast to the studies of Sohn et al. and Kaviraj et al. on Virgo GCs, we extend our analysis to compact stellar systems beyond the mass range of GCs ($M\\lesssim 3 \\times 10^6$ M$_{\\sun}$), including the so-called ultra-compact dwarf galaxies (UCDs, Drinkwater et al. 2003), which cover the mass range up to $\\sim10^8$M$_{\\sun}$, having $M_V \\lesssim -11$ mag. We analyse how the UV properties of UCDs compare to those of both massive and normal GCs, in order to improve our knowledge of EHB occurence in compact stellar systems.  Throughout this study we adopt (m-M)=31.4 mag (Freedman et al. 2001) as distance modulus to the Fornax cluster. \\begin{figure*}[] \\begin{center} \\epsfig{figure=f1a.eps,width=8.6cm} \\epsfig{figure=f1b.eps,width=8.6cm}  \\caption{{\\bf Left:} Map of the central Fornax cluster. The large circles indicate the FoV of the two archival GALEX pointings used for this study. The large dotted circle corresponds to the Deep Imaging Survey (DIS), the large dashed circle corresponds to the Near Galaxies Survey (NGS). The small dots indicate our sample of spectroscopically confirmed compact objects down to $V\\simeq 22$ mag ($M_V=-9.4$ mag). The filled circles indicate the sources with visually verified GALEX matches in the NUV, the asterisks are those sources with matches in the FUV (matching radius 3$''$). {\\bf Right:} V,I colour-magnitude diagram of the same sources as in the left panel. The optical data is from Mieske et al.  (2004 \\& 2006, 2007), and Dirsch et al. (2003). Magenta dots indicate GCs in M31 with UV colours available (Fig.~\\ref{cmds}; Rey et al. 2007), red circles indicate GCs in M87 with FUV colours available (Fig.~\\ref{cmds}; Sohn et al. 2006). Green squares indicate sources with X-ray counterparts from the Chandra study of Scharf et al. (2005). Note that only one of the GALEX UV detections has a detected X-ray counterpart.}  \\label{map} \\end{center} \\end{figure*} ",
        "conclusions": "A UV excess in an old stellar population is likely due to EHB stars. As pointed out in Sect.~\\ref{introduction}, an EHB may be linked to helium-enriched stars (e.g. Ventura et al. 2001, D'Antona et al. 2002). The strong UV excess of the seven massive Fornax GCs beyond the He-enhanced isochrones, especially the NUV excess of the three most extreme GCs with (NUV-V)$<$2.4 mag (see Fig.~\\ref{colcol}), suggests that in these objects, EHB formation is also driven by other processes. In this context, a plausible explanation may be enhanced mass loss of evolved stars, triggered by high stellar densities (Decressin et al.  2007; Huang \\& Gies 2006) and/or large binary fractions.  Excess radiation at short wavelengths can in principle also arise from accretion onto a black hole (King et al. 1993), which can be traced by low-mass X-ray binaries (Jord\\'{a}n et al. 2004).  We have cross-checked the positions of all GALEX UV detections with X-ray source detections in the Chandra Fornax Survey data (Scharf et al. 2005 and private communication), the deepest available wide-field X-ray survey of Fornax (50ks integration with ACIS). The sensitivity of these images is a few $ 10^{38}$ erg/sec, allowing to detect the most luminous LMXBs (Jord\\'{a}n et al.  2004).  In Fig.~\\ref{map} { (right panel)} we indicate the (V-I) optical colours of those compact objects with X-ray matches. At a given magnitude, the X-ray matches { in GCs} are biased towards red optical colours (see also Jord\\'{a}n et al.  2004), while GALEX UV detections are biased towards blue optical colours.  This suggests that generally, the UV- and X-ray-emission of the compact stellar systems are not caused by the same physical processes.  However, there is one GALEX UV detection with an X-ray counterpart (Fig.~\\ref{map} and~\\ref{colcol}), which happens to be one of the three GCs with largest UV excess. We can therefore not exclude that the UV excess in some of the GCs is linked to accretion processes.  We finally note that comparing the probability of UV excess between UCDs and GCs allows to test whether EHBs are more likely associated with present-day deep potential wells (i.e. UCDs) or high stellar densities (i.e. GCs; Dabringhausen et al. 2008, Mieske et al. 2008). One would expect deep potential wells to favour self-enrichment (e.g. Ventura et al. 2001, D'Antona et al. 2002), and high stellar densities to favour mass-loss scenarios (Decressin et al.  2007; Huang \\& Gies 2006). Such a comparison may therefore help to constrain the efficiency of EHB formation channels, provided that the present-day density and mass of the systems investigated have not experienced significant changes during the past, which could have been the case due to core collapse (Noyola \\& Gebhardt 2006, de Marchi et al. 2007) or tidal stripping (e.g. Lee et al.  2007). To properly perform this comparison, deeper UV imaging data will be required that allow detection of UV intermediate-bright to faint GCs down to $M_V\\simeq-$10 mag ($V\\simeq 21.5$ mag at the Fornax distance).  In this respect, the outcomes of the HST observations in Cycle 15 (GO10901, PI O'Connell) of GCs belonging to NGC 1399 are highly anticipated."
    },
    "0808/0808.3374_arXiv.txt": {
        "abstract": "At a distance of about 130~pc, the Corona Australis molecular cloud complex is one of the nearest regions with ongoing and/or recent star formation. It is a region with highly variable extinction of up to $A_{V} \\sim 45$ mag, containing, at its core, the \\textsl{Coronet} protostar cluster. There are now 55 known optically detected members, starting at late B spectral types. At the opposite end of the mass spectrum, there are two confirmed brown dwarf members and seven more candidate brown dwarfs. The CrA region has been most widely surveyed at infrared wavelengths, in X-rays, and in the millimeter continuum, while follow-up observations from centimeter radio to X-rays have focused on the \\textsl{Coronet} cluster. ",
        "introduction": "{\\em Corona Australis (southern crown)}, abbreviated {\\em CrA}, is one of the 48 ancient constellations listed by Ptolemy, who called it {\\em Stephanos Notios} (in Greek). The latin translation was {\\em corona notia, austrina, australis, meridiana,} or {\\em meridionalis}, from which \\citet{all36} picked {\\em corona australis}; see also \\citet{bot80} for a discussion of the name. For an overview, see Fig.~\\ref{lokekuntan}.  \\subsection{The Discovery} \\label{sec_disc}  Given the convention for naming variable stars, R CrA was the first variable star noticed in the CrA constellation. Marth from Malta and Schmidt from Athens (quoted in \\citealp{rey16}) were the first to notice the variability of the nebula NGC 6729 in CrA; Schmidt also discovered the variable stars R, S, and T CrA with the nebula attached to R CrA. The variability of stars and nebula were confirmed by Innes at Cape Observatory (quoted in \\citealp{rey16}) and \\citet{kno16}; the latter gave $\\Delta m$ = 1.4 mag for the variability of R CrA in images taken from 21 Aug 1911 to 23 May 1915 at Helwan Observatory (Egypt). \\citet{rey16} also noticed that the behaviour of R CrA (variability and environment) is similar to that of the star T Tau. \\citet{luy27} found the three bright variables UZ CrA, VW CrA, and VX CrA. Then, \\citet{vge32,vge33} found more than 100 variable stars in or near CrA, based on plates taken at the Union Observatory in Johannesburg, mostly by himself and Hertzsprung, a few of them young stars.  \\citet{her60} estimated the age of the two early-type emission-line stars R CrA (Ae) and T CrA (F0e) as their main sequence contraction time ($10^{5}$ to $10^{7}$ years) and concluded that they are young, and that therefore also the associated nebula and the other stars should be young. \\citet{kna73} found infrared (IR) excess in some of the variable stars in CrA, indicative of circumstellar material, and derived $\\sim 1$ Myr as the age of the group.  The Corona Australis molecular cloud complex is today known as one of the nearest regions with ongoing and/or recent intermediate- and low-mass star formation. The dark cloud near R CrA is the densest cloud core in the region with extinction of up to $A_{V} \\sim 45$ mag \\citep{dau87,wil92}. This cloud is also called {\\em condensation A} \\citep{ros78}, {\\em FS~445 -- FS~447} \\citep{fes84}, DoH {\\em 2213} \\citep[][see Fig.~\\ref{dobashifig}]{dob05}, and sometimes also just {\\em CrA} or {\\em R CrA cloud}; we use {\\em CrA} for {\\em Corona Australis} as the name for the star-forming region.  \\begin{figure*} \\includegraphics[width=\\linewidth]{cra_fig1.eps} \\caption{Stars and dust in CrA. Credit: Loke Kun Tan (StarryScapes).} \\label{lokekuntan} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics*[width=0.95\\linewidth,bb=20 54 576 546]{cra_fig2.eps} \\caption{The CrA cloud in an extinction map from \\citet{dob05}. Note the location relative to the Galactic plane at $b\\approx-18^\\circ$.} \\label{dobashifig} \\end{figure*}  \\begin{figure*} \\includegraphics*[width=\\linewidth,bb=19 23 575 266]{cra_fig3.eps} \\caption{Extinction map of CrA from $B$-band star counts (from \\citealp{cam99}).} \\label{cambresy_fig6} \\end{figure*}   \\subsection{The Galactic Environment}  The CrA dark cloud is located $\\sim 18^{\\circ}$ below the Galactic plane (see Fig.~\\ref{dobashifig}). According to \\citet[see also Fig. 1.10 in \\citealp{poe97}]{ola82}, the CrA cloud is not part of the Gould Belt nor of the Lindblad ring \\citep{poe97}, but is located within the massive HI shell called {\\em Loop I} \\citep{har93}. The 3D space motion of the CrA T~Tauri stars would be consistent with the dark cloud being formed originally by a high-velocity cloud impact onto the Galactic plane, which triggered the star formation in CrA \\citep{neu00}. When a high-velocity cloud impacts onto the gas and dust in the Galactic plane, this high-velocity cloud may get distroyed, but then forms a new (lower velocity) cloud. Alternatively, the star formation in CrA may have been caused by the expansion of the UpperCenLupus (UCL) superbubble: \\citet{mam01} found that the CrA complex is radially moving away from UCL with 7~km~s$^{-1}$  so that it was located near its center $\\sim 14$~Myrs ago.  \\citet{har93} resolved cloud {\\em A} into five condensations, the star R~CrA being located in {\\em A2}. \\citet{anv96} and \\citet{cam99} mapped the cloud using optical star counts (see Fig.~\\ref{cambresy_fig6}). \\citet{juv08} study the relationship between scattered near-infrared light and extinction in a quiescent filament of the CrA cloud and compare the result to extinction measurements. Several IR surveys revealed a large population of embedded IR sources \\citep{tas84,wil86,wil89,wil92,wil97} some of which are IR Class~I objects. Based on their $^{13}$CO map, \\citet{tac02} estimated the star formation efficiency in CrA to be $\\sim 4~\\% $.  \\subsection{The Distance}  The distance towards the CrA star forming region was estimated by \\citet{gap36} to be $150 \\pm 50$~pc and later by \\citet{mar81} to be $\\sim 129$~pc (assuming a ratio of total to selective absorption of $R=4.5$, see \\citealp{vrb81}). The \\textsl{Hipparcos} satellite attempted to measure the parallax of the star R CrA and found $122 \\pm 68$ mas, i.e. no reliable solution. \\citet{knh98} gave $\\sim 170$ pc as distance. Three likely CrA members have reliable \\textsl{Hipparcos} parallax measurements and thus distances; these are HR~7170 ($d$=$85\\pm26$~pc), SAO~210888 ($d$=$193\\pm39$~pc), and HD~176386 ($d$=$139\\pm22$~pc). R~CrA and HR~7169 do not have reliable parallax measurements. The weighted mean of these three distances is $128\\pm26$~pc. \\citet{cas98} determined the distance towards the eclipsing double-lined spectroscopic binary TY CrA to be $129 \\pm 11$~pc from their orbit solution. The orbit solution of TY~CrA provides the currently best estimate for the distance to the star-forming region, and we therefore suggest to use a distance of 130~pc for CrA, see also \\citet{dez99}. ",
        "conclusions": ""
    },
    "0808/0808.3835_arXiv.txt": {
        "abstract": "% NGC\\,2264 is a young Galactic cluster and the dominant component of the Mon OB1 association lying approximately 760 pc distant within the local spiral arm. The cluster is hierarchically structured, with subclusters of suspected members spread across several parsecs. Associated with the cluster is an extensive molecular cloud complex spanning more than two degrees on the sky. The combined mass of the remaining molecular cloud cores upon which the cluster is superposed is estimated to be at least $\\sim$3.7$\\times$10$^{4}$ M$_{\\odot}$. Star formation is ongoing within the region as evidenced by the presence of numerous embedded clusters of protostars, molecular outflows, and Herbig-Haro objects. The stellar population of NGC\\,2264 is dominated by the O7 V multiple star, S~Mon, and several dozen B-type zero-age main sequence stars. X-ray imaging surveys, H$\\alpha$ emission surveys, and photometric variability studies have identified more than 600 intermediate and low-mass members distributed throughout the molecular cloud complex, but concentrated within two densely populated areas between S~Mon and the Cone Nebula. Estimates for the total stellar population of the cluster range as high as 1000 members and limited deep photometric surveys have identified $\\sim$230 substellar mass candidates. The median age of NGC\\,2264 is estimated to be $\\sim$3 Myr by fitting various pre-main sequence isochrones to the low-mass stellar population, but an apparent age dispersion of at least $\\sim$5 Myr can be inferred from the broadened sequence of suspected members. Infrared and millimeter observations of the cluster have identified two prominent sites of star formation activity centered near NGC\\,2264 IRS1, a deeply embedded early-type (B2--B5) star, and IRS2, a star forming core and associated protostellar cluster. NGC\\,2264 and its associated molecular clouds have been extensively examined at all wavelengths, from the centimeter regime to X-rays. Given its relative proximity, well-defined stellar population, and low foreground extinction, the cluster will remain a prime candidate for star formation studies throughout the foreseeable future. ",
        "introduction": "Few star forming regions surpass the resplendent beauty of NGC\\,2264, the richly populated Galactic cluster in the Mon OB1 association lying approximately 760 pc distant in the local spiral arm. Other than the Orion Nebula Cluster, no other star forming region within one kpc possesses such a broad mass spectrum and well-defined pre-main sequence population within a relatively confined region on the sky. Estimates for the total stellar population of the cluster range up to $\\sim$1000 members, with most low-mass, pre-main sequence stars having been identified from H$\\alpha$ emission surveys, X-ray observations by {\\it ROSAT}, {\\it Chandra}, and {\\it XMM-Newton}, or by photometric variability programs that have found several hundred periodic and irregular variables. The cluster of stars is seen in projection against an extensive molecular cloud complex spanning more than two degrees north and west of the cluster center. The faint glow of Balmer line emission induced by the ionizing flux of the cluster OB stellar population contrasts starkly with the background dark molecular cloud from which the cluster has emerged. The dominant stellar member of NGC\\,2264 is the O7 V star, S~Monocerotis (S~Mon), a massive multiple star lying in the northern half of the cluster. Approximately 40\\arcmin\\ south of S~Mon is the prominent Cone Nebula, a triangular projection of molecular gas illuminated by S~Mon and the early B-type cluster members. NGC\\,2264 is exceptionally well-studied at all wavelengths: in the millimeter by Crutcher et al. (1978), Margulis \\& Lada (1986), Oliver et al. (1996), and Peretto et al. (2006); in the near infrared (NIR) by Allen (1972), Pich\\'{e} (1992, 1993), Lada et al. (1993), Rebull et al. (2002), and Young et al. (2006); in the optical by Walker (1956), Rydgren (1977), Mendoza \\& G\\'omez (1980), Adams et al. (1983), Sagar \\& Joshi (1983), Sung et al. (1997), Flaccomio et al. (1999), Rebull et al. (2002), Sung et al. (2004), Lamm et al. (2004), and Dahm \\& Simon (2005); and in X-rays by Flaccomio et al. (2000), Ramirez et al. (2004), Rebull et al. (2006), Flaccomio et al. (2006), and Dahm et al. (2007).  NGC\\,2264 was discovered by Friedrich Wilhelm Herschel in 1784 and listed as H VIII.5 in his catalog of nebulae and stellar clusters. The nebulosity associated with NGC\\,2264 was also observed by Herschel nearly two years later and assigned the designation: H V.27. The Roman numerals in Herschel's catalog are object identifiers, with `V' referring to very large nebulae and `VIII' to coarsely scattered clusters of stars. One of the first appearances of the cluster in professional astronomical journals is Wolf's (1924) reproduction of a photographic plate of the cluster and a list of 20 suspected variables. Modern investigations of the cluster begin with Herbig (1954) who used the slitless grating spectrograph on the Crossley reflector at Lick Observatory to identify 84 H$\\alpha$ emission stars, predominantly T Tauri stars (TTS), in the cluster region. Herbig (1954) postulated that these stars represented a young stellar population emerging from the dark nebula. Walker's (1956) seminal photometric and spectroscopic study of NGC\\,2264 discovered that a normal main sequence exists from approximately O7 to A0, but that lower mass stars consistently fall above the main sequence. This observation was in agreement with predictions of early models of gravitational collapse by Salpeter and by Henyey et al. (1955). Walker (1954, 1956) proposed that these stars represent an extremely young population of cluster members, still undergoing gravitational contraction. Walker (1956) further noted that the TTSs within the cluster fall above the main sequence and that they too may be undergoing gravitational collapse. Walker (1956) concluded that the study of TTSs would be ``of great importance for our understanding of these early stages of stellar evolution.''  \\begin{figure}[!tbh] \\begin{center} \\includegraphics[angle=90,width=\\linewidth, draft=False]{fg1.eps} \\caption[fg1.eps]{A 12.5$^{\\circ}$ square false-color IRAS image (100, 60 \\& 25~$\\mu$m) of the Mon OB1 and Mon R1 associations. NGC\\,2264 lies at the center of the image with several nearby IRAS sources identified, including the reflection nebulae NGC\\,2245 and NGC\\,2247, and NGC\\,2261 (Hubble's Variable Nebula). South of NGC\\,2264 is the Rosette Nebula and its embedded cluster NGC\\,2244, lying 1.7 kpc distant in the Perseus arm. \\label{f1}} \\end{center} \\end{figure}  The molecular cloud complex associated with NGC\\,2264 was found by Crutcher et al. (1978) to consist of several cloud cores, the most massive of which lies roughly between S~Mon and the Cone Nebula. Throughout the entire cluster region, Oliver et al. (1996) identified 20 molecular clouds ranging in mass from $\\sim10^{2}$ to $10^{4}$ M$_{\\odot}$. With NGC\\,2264 these molecular clouds comprise what is generally regarded as the Mon OB 1 association. Active star formation is ongoing within NGC\\,2264 as evidenced by the presence of numerous embedded protostars and clusters of stars, as well as molecular outflows and Herbig-Haro objects (Adams et al. 1979; Fukui 1989; Hodapp 1994; Walsh et al. 1992; Reipurth et al. 2004a; Young et al. 2006). Two prominent sites of star formation activity within the cluster are IRS1 (also known as Allen's source), located several arcminutes north of the tip of the Cone Nebula, and IRS2, which lies approximately one-third of the distance from the Cone Nebula to S~Mon. New star formation activity is also suspected within the northern extension of the molecular cloud based upon the presence of several embedded IRAS sources and giant Herbig-Haro flows (Reipurth et al. 2004a,c). From 60 and 100~$\\mu$m IRAS images of NGC\\,2264, Schwartz (1987) found that the cluster lies on the eastern edge of a ring-like dust structure, 3$^{\\circ}$ in diameter. Shown in Figure 1 is a 12\\fdg5$\\times$12\\fdg5 false-color IRAS image (100, 60, and 25~$\\mu$m) centered near NGC\\,2264. The reflection nebulae NGC\\,2245 and NGC\\,2247, members of the Mon R1 association, are on the western boundary of this ring (see the chapter by Carpenter \\& Hodapp). Other components of the Mon R1 association include the reflection nebulae IC\\,446 and IC\\,2169, LkH$\\alpha$215, as well as several early type (B3--B7) stars. It is generally believed that the Mon R1 and Mon OB 1 associations are at similar distances and are likely related. The Rosette Nebula, NGC\\,2237-9, and its embedded young cluster NGC\\,2244 lie 5$^{\\circ}$ southwest of NGC\\,2264, 1.7 kpc distant in the outer Perseus arm (see the chapter by Rom\\'an-Z\\'u\\~niga \\& Lada). Several arcs of dust and CO emission have been identified in the region, which are believed to be supernovae remnants or windblown shells. Many of these features are apparent in Figure 2, a wide-field H$\\alpha$ image of NGC\\,2264, NGC\\,2244, and the intervening region obtained by T. Hallas and reproduced here with his permission. It is possible that star formation in the Mon OB1 and R1 associations was triggered by nearby energetic events, but it is difficult to assess the radial distance of the ringlike structures evident in Figure 2, which may lie within the Perseus arm or the interarm region. Shown in Figure 3 is a narrow-band composite image of NGC\\,2264 obtained by T.A. Rector and B.A. Wolpa using the 0.9 meter telescope at Kitt Peak. S~Mon dominates the northern half of the cluster, which lies embedded within the extensive molecular cloud complex.   \\begin{figure} \\centering \\includegraphics[angle=0,width=\\linewidth, draft=False]{fg2.eps} \\caption[fg2.eps]{An extraordinary widefield, narrow-band H$\\alpha$ image of NGC\\,2264 (upper center), the Rosette Nebula and NGC\\,2244 (lower right), and the numerous windblown shells and supernova remnants possibly associated with the Mon OB1 or Mon OB2 associations. The Cone Nebula is readily visible just above and left of image center as is S~Mon. Also apparent in the image is the dark molecular cloud complex lying to the west of NGC\\,2264. This image is a composite of 16 20-minute integrations obtained by T. Hallas using a 165 mm lens and an Astrodon H$\\alpha$ filter. \\label{f2}} \\end{figure}  To summarize all work completed over the last half-century in NGC\\,2264  would be an overwhelming task and require significantly more pages than alloted for this review chapter. The literature database for NGC\\,2264 and its members has now grown to over 400 refereed journal articles, conference proceedings, or abstracts. Here we attempt to highlight large surveys of the cluster at all wavelengths as well as bring attention to more focused studies of the cluster that have broadly impacted our understanding of star formation. The chapter begins with a review of basic cluster properties including distance, reddening, age, and inferred age dispersion. It then examines the OB stellar population of the cluster, the intermediate and low-mass stars, and finally the substellar mass regime. Different wavelength regions are examined from the centimeter, millimeter, and submillimeter to the far-, mid-, and near infrared, the optical, and the X-ray regimes. We then review many photometric variability studies of the cluster that have identified several hundred candidate members. Finally, we consider future observations of the cluster and what additional science remains to be reaped from NGC\\,2264. The cluster has remained in the spotlight of star formation studies for more than 50 years, beginning with the H$\\alpha$ survey of Herbig (1954). Its relative proximity, low foreground extinction, large main sequence and pre-main sequence populations, the lack of intense nebular emission, and the tremendous available archive of observations of the cluster at all wavelengths guarantee its place with the Orion Nebula Cluster and the Taurus-Auriga molecular clouds as the most accessible and observed Galactic star forming region.  \\begin{figure}[!tbh] \\vspace{8mm} \\begin{center} \\includegraphics[angle=0,width=\\linewidth, draft=False]{fg3.eps} \\caption[fg3.eps]{Narrow-band, three-color image of NGC\\,2264 obtained by T.A. Rector and B.A. Wolpa (NOAO/AURA/NSF) using the 0.9 meter telescope at Kitt Peak National Observatory. The filters used for this composite image are: O III (light blue), H$\\alpha$ (red-orange), and [S II] (blue-violet). The field of view is approximately 0.75$^{\\circ}$$\\times$1$^{\\circ}$. S~Mon lies just above the image center and is believed to be the ionizing source of the bright rimmed Cone Nebula. \\label{f2}} \\end{center} \\vspace{-8mm} \\end{figure} ",
        "conclusions": ""
    },
    "0808/0808.3718_arXiv.txt": {
        "abstract": "We present a system of X-ray photometry for the \\chandra satellite. X-ray photometry can be a powerful tool to obtain flux estimates, hardness ratios, and colors unbiased by assumptions about spectral shape and independent of temporal and spatial changes in instrument characteristics. The system we have developed relies on our knowledge of effective area and the energy-to-channel conversion to construct filters similar to photometric filters in the optical bandpass. We show that the filters are well behaved functions of energy and that this X-ray photometric system is able to reconstruct fluxes to within about 20\\%, without color corrections, for non-pathological spectra. Even in the worst cases it is better than 50\\%. Our method also treats errors in a consistent manner, both statistical as well as systematic. ",
        "introduction": "\\label{sec:intro} Photometry is one of the most widely used, relatively simple, tools used in describing and categorizing astronomical objects. Standardization by \\citet{johnson:53} and subsequent additions in the optical and infrared (see e.g. \\citet{bessel:05}) have allowed comparisons between measurements by different telescopes and instruments without bias. Important applications of optical photometry include stellar classifications (see e.g. \\citet{johnson:53}), galaxy redshifts \\citep{puschell:81}, and the discovery of the most distant quasars in the SDSS \\citep{fan:99}. In particular photometry is important for sources too faint to extract detailed spectra, i.e. most sources, given the almost universal increase of source numbers to faint fluxes.  Although optical astronomy is a much older discipline than X-ray astronomy, optical photometry was established only about 20 years before the beginning of X-ray astronomy in the 1960s. Given the enormous success of optical photometry it seems obvious to use it as a precedent and try to duplicate its success in other wavebands beyond UV and infrared. The X-ray band can be defined as reaching from about 0.1 keV to a few hundred keV, spanning almost 4 decades of frequency -- although most work concentrates on the 0.2--20 keV band -- compared with 2 octaves in the optical. Moreover, an important difference between optical and X-ray is the typically low number of photons in X-ray astronomy. The X-ray range is photon starved such that sources with a few hundred counts are considered bright in X-rays. This limitation increases the importance of broad-band photometry in the X-ray band.  While X-ray astronomy has used relative, mission-specific photometry for most of its existence, there is, as yet, no standard X-ray photometric system. The usefulness of an X-ray photometric system is evident already from the use of these somewhat idiosyncratic energy bands. The bands used in the past have been chosen for specific purposes, e.g. to use color--color diagrams to diagnose X-ray binary spectral states \\citep{white:84,hasinger:89} where, in the latter, the energy bands are different for each source. Even so, the resulting color--color diagrams have immensely increased our knowledge of X-ray binary spectral/accretion states \\citep[e.g.][]{prestwich:03,gierlinski:06}. Thus a standard X-ray photometric system is highly desirable in X-ray astronomy in order to cross-compare observations of the hundreds of thousands of sources being cataloged by XMM \\citep{watson:07}, \\chandra \\citep{fabbiano:07}, and other missions. Even within a given mission different types of CCDs (XMM) or changes in operating temperature, gain or contamination (\\chandra) mean that simple count rates cannot be used.  However, there have been complicating factors in establishing photometric energy bands beyond individual observations. One cause of this lack of standardization is that the energy ranges covered by different X-ray satellites and instruments differ widely. For example, RXTE/PCA, GINGA and EXOSAT bands have practically no overlap with ROSAT bands; and ASCA, XMM and \\chandra bands are somewhere in the middle. Fig. \\ref{fig:ebands} shows a selection of energy bands used by different authors for different X-ray satellites\\footnote{For more information on X-ray satellites, see the HEASARC web page http://heasarc.gsfc.nasa.gov/docs/observatories.html. A list of energy bands and references is given in the on-line version of this paper.}. Most X-ray missions with focusing optics cover the energy range from $\\sim$0.1 keV to 10 keV.  Another reason for the lack of a standard photometric system is that in X-rays there are no bright constant point sources in the way that stars can be used for calibration like in the optical.  But the most fundamental cause for the lack of an X-ray photometric system has been the limited spectral resolving power (R$=$E/$\\Delta E\\sim$1) of proportional counters, which were used in X-ray astronomy from the earliest days through to ROSAT and RXTE. A resolution of R$\\sim$1 allows no clean separation of energy bands, and different spectra with similar flux will give widely different flux estimates in any chosen band. In optical terms, the ``color correction'' is very large. However, with the introduction of X-ray CCDs in ASCA \\citep{burke:93} this limitation has largely gone away. X-ray CCDs have R$>$10, so comparable to the R$\\sim$6 of broad band optical photometry. It seems thus timely to consider the introduction of an X-ray photometric system. Therefore we have investigated how good a photometric system can be created for \\chandra ACIS observations and, by extension, for all other X-ray CCDs. We report the encouraging results in this paper.  \\bfig[h] \\resizebox{\\hsize}{!}{\\includegraphics{f1.eps}} \\caption{Energy bands used in various publications for X-ray satellites. The different energy bands for different satellites are due to energy coverage of the instruments.  The colors represent the width of the soft/medium/hard bands used in the corresponding papers. A table of energy bands and references are given in the on-line version of the paper.\\label{fig:ebands}} \\efig ",
        "conclusions": "We have presented a system of X-ray photometry for the \\chandra satellite. The system we have developed relies on a knowledge of effective area and the energy-to-channel conversion to construct X-ray filters, but is unbiased by assumptions about the spectral shape of a source. We have shown that the filters are comparable to filters in the optical and infrared, and that our photometric system in X-rays is able to estimate fluxes to within about 20\\%. Even in the worst cases it is better than 50\\%. We have incorporated methods to estimate systematic errors and consistently propagate statistical as well as systematic errors.  Due to the construction method employed our filter system is very flexible and can be adapted readily to other CCD X-ray detectors, in particular to XMM-Newton EPIC. The code to compute fluxes is available at a web page \\footnote{http://hea-www.cfa.harvard.edu/~jcm/xray/index.html}. A table with a selection of correction factors for ACIS-S3 is given in Appendix A in the electronic edition and available at the same URL. The table contains only every fourth tile because of the generally slowly varying correction factor with chip location. In the future we will explore potential improvements to X-ray photometry by means of:  \\begin{itemize} \\item Making color corrections using band ratios. A preliminary investigation suggests that extreme spectra (e.g. highly absorbed or high photon indexes) will gain significantly in the accuracy of flux estimates and even normal spectra will have a reduced error range.  \\item Optimizing the choice of bands. The properties of our current method show that there is a correlation between the accuracy of flux estimates and the filter shape. The more ``boxy'' a filter is the better the flux estimate will be. This suggests a limit for the width of a filter at which point deviations from boxiness result in an accuracy of the flux estimate below a certain value.  \\item Comparing results for \\chandra ACIS and XMM-Newton EPIC data. \\chandra and XMM have similar instrumental setups (ARF and RMF) and overlapping science capabilities. And given the large numbers of sources observed by these two X-ray missions it is very important to be able to compare results for the two. \\end{itemize}"
    },
    "0808/0808.2414_arXiv.txt": {
        "abstract": "Broad absorption line quasars (commonly termed BALQSOs) contain the most dramatic examples of AGN-driven winds. The high absorbing columns in these winds, $\\sim$~$10^{24}$~cm$^{-2}$, ensure that BALQSOs are generally X-ray faint. This high X-ray absorption means that almost all BALQSOs have been discovered through optical surveys, and so what little we know about their X-ray properties is derived from very bright optically-selected sources. A small number of X-ray selected BALQSOs (XBALQSOs) have, however, recently been found in deep X-ray survey fields. In this paper we investigate the X-ray and rest-frame UV properties of five XBALQSOs for which we have obtained XMM-Newton EPIC X-ray spectra and deep optical imaging and spectroscopy. We find that, although the XBALQSOs have an $\\alpha_{ox}$ steeper by $\\sim$~0.5 than normal QSOs, their median $\\alpha_{ox}$ is nevertheless flatter by 0.30 than that of a comparable sample of optically selected BALQSOs (OBALQSOs). We rule out the possibility that the higher X-ray to optical flux ratio is due to intrinsic optical extinction. We find that the amount of X-ray and UV absorption due to the wind in XBALQSOs is similar, or perhaps greater than, the corresponding wind absorption in OBALQSOs, so the flatter $\\alpha_{ox}$ cannot be a result of weaker wind absorption. We conclude that these XBALQSOs have intrinsically higher X-ray to optical flux ratios than the OBALQSO sample with which we compare them. ",
        "introduction": "In seeking to understand any role that AGN winds might play in the coevolution of supermassive black holes and their host galaxies, we need to study the physical properties of those winds at z~$\\sim$~2 where star formation and black hole growth were at their height. It is probable that the majority of mass and energy carried in AGN winds is transported by the X-ray absorbing part of the outflow \\citep[e.g.][]{blustin2007}, as the X-ray absorber has the highest column density, at least in nearby AGN, but it is usually difficult to obtain satisfactory X-ray spectra from high-redshift AGN.  The most well-known population of distant AGN showing evidence of ionised winds is the Broad Absorption Line quasars (BALQSOs; z~$\\sim$~0.1$-$6), whose winds have outflow speeds of up to 60000 km~s$^{-1}$, with line-of-sight absorbing columns of $10^{23}-10^{24}$ cm$^{-2}$ \\citep[e.g.][]{krolik1999} causing the soft X-ray band to be highly absorbed \\citep[see~e.g.][~and~references~therein]{gallagher2001}. The vast majority of X-ray observations of BALQSOs have been follow$-$ups after their discovery in optical surveys, and most of the \\emph{XMM-Newton} and \\emph{Chandra} observations of BALQSOs, with the exception of LBQS~2212-1759 \\citep{clavel2006} and APM~08279+5255 \\citep{chartas2002}, have been short snapshots. Our knowledge of the X-ray properties of BALQSOs is therefore based upon a small number of the very brightest optically-selected sources; the largest relevant samples to date have been published by \\citet{green2001,punsly2006,gallagher2006}.  Recently a number of X-ray selected BALQSOs (XBALQSOs) have also come to light \\citep[Page~et~al.,~in~preparation;][]{barcons2003}. We use the abbreviation XBALQSOs to refer to serendipitously discovered X-ray sources whose BALQSO nature has been revealed by follow$-$up optical spectroscopic observations. It is unknown whether these represent a distinct section of the BALQSO population. Since the X-ray faintness of BALQSOs is probably due to absorption by the wind, the relative X-ray brightness of XBALQSOs could be due to their winds causing less absorption, through either having lower absorbing columns or higher ionisation parameters; XBALQSOs could also conceivably be exceptionally luminous BALQSOs with high dust extinction in the optical. There is also the possibility that their intrinsic continua simply have a higher ratio of X-ray to optical flux.  In this paper, we examine the properties of five XBALQSOs discovered in two deep X-ray survey fields: four objects in the 1$^{H}$ survey field and one in the Chandra Deep Field South (CDFS). We begin by comparing the X-ray and optical fluxes, and X-ray to optical spectral indices ($\\alpha_{ox}$) of the sources to those of the largest available uniformly-analysed sample of OBALQSOs observed to date in the X-ray \\citep{gallagher2006}. We then investigate whether the intrinsic extinction in the XBALQSOs differs from that in the \\citet{gallagher2006} sample and from QSOs from the SDSS \\citep{vandenberk2001}. In section~\\ref{dynamics}, we compare the extent of intrinsic UV absorption in the XBALQSOs with that in the \\citet{trump2006} sample of OBALQSOs from the SDSS. Next, in section~\\ref{xray_wind_properties} we obtain estimates of the absorbing columns and ionisation parameters of the X-ray absorbing winds of the XBALQSOs, comparing the results with those of similar analyses of OBALQSOs in the literature. We then provide estimates of the mass-energy budgets of the XBALQSOs and their X-ray absorbing winds. Finally we discuss the nature of these XBALQSOs, and the implications for the population of faint BALQSOs discoverable in deep X-ray surveys.  Names, positions and basic parameters of these sources are given in Table~\\ref{source_properties}. Throughout this paper, the objects are referred to by the abbreviations XBQ1-5 (X-ray selected broad absorption line quasars 1 to 5), as identified in Table~\\ref{source_properties}. We use a cosmology with H$_{0}$=70~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{M}$=0.3, and $\\Omega_{\\Lambda}$=0.7, and all uncertainties are 1$\\sigma$ unless otherwise stated.  \\begin{table*} \\centering \\begin{minipage}{120mm} \\caption{Basic properties of the five X-ray selected BALQSOs: IAU name and abbreviated name for use in this paper; RA; Dec; r, offset between the optical and X-ray positions in arcseconds; redshift (see section~\\ref{optical_data}); Galactic extinction using reference pixel in \\citet{schlegel1998}; de-reddened AB $B$ and $K_s$ magnitudes, where the figures for XBQ5 are from the MUSYC survey (\\citealt{gawiser2006}; Taylor et al., in preparation; see section~\\ref{optical_data}); C$_{x}$, number of background-subtracted X-ray counts in the 0.2$-$12~keV range.} \\begin{tabular}{@{}llllll@{}} \\hline Name  & & & RA & Dec & r  \\\\ \\hline CXOU J014546.9-043031    & & XBQ1 & 01:45:46.88 & -04:30:30.50 & 0.50  \\\\ CXOU J014517.1-042503    & & XBQ2 & 01:45:17.05 & -04:25:03.20 & 0.91  \\\\ CXOU J014553.9-043726    & & XBQ3 & 01:45:53.93 & -04:37:26.20 & 0.74   \\\\ CXOU J014546.5-042450    & & XBQ4 & 01:45:46.54 & -04:24:49.50 & 0.78   \\\\ CXOCDFS J033209.5-274807 & & XBQ5 & 03:32:09.46 & -27:48:06.80 & 0.28\\footnote{Table 2, \\citet{giacconi2002}} \\\\ \\hline \\hline Name & z & E(B-V) & $B$ & $K_s$ & C$_{x}$\\\\ \\hline XBQ1 & 2.63  & 0.023 & 21.8 $\\pm$ 0.1 & 21.1 $\\pm$ 0.2 & 208  \\\\ XBQ2 & 1.793 & 0.023 & 21.4 $\\pm$ 0.1 & 20.7 $\\pm$ 0.2 & 154  \\\\ XBQ3 & 1.40  & 0.022 & 21.5 $\\pm$ 0.1 & 19.6 $\\pm$ 0.2 & 74  \\\\ XBQ4 & 2.64  & 0.024 & 20.1 $\\pm$ 0.1 & 19.5 $\\pm$ 0.2 & 137  \\\\ XBQ5 & 2.82  & 0.009 & 21.33 $\\pm$ 0.05 & 20.00 $\\pm$ 0.06 & 1115  \\\\ \\hline \\end{tabular} \\end{minipage} \\label{source_properties} \\end{table*} ",
        "conclusions": "\\label{discussion}  Our goal in this paper was to investigate the X-ray and optical properties of a group of X-ray selected BALQSOs discovered in deep X-ray survey fields, and in particular to see whether they had been discoverable in such surveys due to having unusual spectral properties.  As we showed in section~\\ref{fluxes_alpha_ox}, the XBALQSOs have, on average, slightly higher X-ray to optical flux ratios than the largest available comparable sample of OBALQSOs \\citep{gallagher2006}. This is probably not a result of intrinsic source variability. It is equally possible that all five sources would happen to be especially X-ray faint as X-ray bright when we observed them, and if variability was a major factor, we would expect to see a mixture of high and low $\\alpha_{ox}$ values among the XBALQSOs.  We can also immediately discard the possibility that the flatter $\\alpha_{ox}$ is due to the XBALQSOs being intrinsically very luminous BALQSOs which happen to have high rest-frame optical extinction; as we showed in Section~\\ref{extinction_properties}, their $B-Ks$ colours as a function of redshift are entirely consistent with the OBALQSO population. Given the very low space density of luminous QSOs on the sky, it is also highly unlikely that we should have found four such objects serendipitously in a single 30\\arcmin\\ diameter field.  If XBALQSOs are brighter than expected in the X-rays due to having less intrinsic X-ray absorption than optically selected sources, the ionised X-ray absorbing winds should have lower N$_{H}$ and/or higher $\\xi$ than those in comparable OBALQSOs. In Fig.~\\ref{xray_properties_avg} we compare total absorbing column and (weighted average, where necessary) ionisation parameter with these quantities in OBALQSOs from the literature, as well as from a sample of lower redshift Seyfert galaxies and QSOs with ionised soft X-ray absorbers which we have studied previously \\citep{blustin2005}. The OBALQSOs were PG1115+080 \\citep{chartas2007}, Q1246-057, SBS 1542+541 \\citep{grupe2003}, and PG2112+059 \\citep{gallagher2004}. We find that, on the contrary, these XBALQSOs have similar or even perhaps greater intrinsic absorption than the OBALQSOs. Testing whether that is a general feature of XBALQSOs would require a larger sample of sources.  \\begin{figure} \\includegraphics[width=55mm,angle=-90]{xray_properties_avg_v6.ps} \\caption{A comparison of total Log $N_{H}$ (the sum of $N_{H}$ for the individual phases) and column-weighted average log $\\xi$ for the X-ray selected BALQSOs in our sample (red stars) with four optically-selected BALQSOs from the literature (black circles; see section~\\ref{discussion}) and a sample of low$-$redshift `warm absorber' Seyferts and QSOs (blue squares; \\citealt{blustin2005}).} \\label{xray_properties_avg} \\end{figure}  The amount of absorption from the UV-absorbing part of the wind, as measured by the BI, is consistent with the distribution of BI for OBALQSOs in the SDSS, although there is again some indication that they are towards the more highly absorbed end of the distribution.  The mass-energy transport of the winds, which we estimated in Section~\\ref{mass_energy_budget}, ranges quite widely. For XBQ5, which has the best-constrained values, mass outflow rate via the X-ray absorbing part of the wind is $\\sim$~0.04$-$10 times the mass accretion rate onto its black hole, and the kinetic luminosity of the wind is $\\sim$~0.002$-$0.3\\% of the bolometric luminosity. A complete picture of the mass-energy budget of the wind would require consideration of the mass-energy output via the UV-absorbing part of the wind, which is beyond the scope of this paper, and taking account of the full ionisation range of the wind, which would require far better data.  Regarding the nature of XBALQSOs, we are left with the conclusion that, although they are fundamentally similar to the wider population of OBALQSOs, they may belong to a part of the BALQSO population with an intrinsically higher X-ray to optical flux ratio. This is actually what we would expect for AGN as UV-faint as our XBALQSOs, since, according to the \\citet{vignali2003} relation, for example, AGN with lower restframe 2500~\\AA\\ luminosities should have flatter $\\alpha_{ox}$ indices. The implication of this is that, as hard X-ray surveys get deeper, there should be a lot of faint XBALQSO-type AGN waiting to be discovered. It is interesting to ask whether the XBALQSO $\\alpha_{ox}$ values do in fact exactly follow the \\citet{vignali2003} relation; our XBALQSOs do not span sufficient luminosity space for us to attempt an answer to that question. Investigating this point will require many more sources over a range of luminosities, with sufficient X-ray counts for a reliable estimation of the unabsorbed X-ray flux to be made."
    },
    "0808/0808.1598_arXiv.txt": {
        "abstract": "We describe a data reduction pipeline for VLBI astrometric observations of pulsars, implemented using the ParselTongue AIPS interface.  The pipeline performs calibration (including ionosphere modeling), phase referencing with proper accounting of reference source structure, amplitude corrections for pulsar scintillation, and position fitting to yield the position, proper motion and parallax.  The optimal data weighting scheme to minimize the total error budget of a parallax fit, and how this scheme varies with pulsar parameters such as flux density, is also investigated. The robustness of the techniques employed are demonstrated with the presentation of the first results from a two year astrometry program using the Australian Long Baseline Array (LBA). The parallax of PSR J1559--4438 is determined to be $\\pi = 0.384 \\pm 0.081$\\ mas $(1\\sigma)$, resulting in a distance estimate of 2600 pc which is consistent with earlier DM and HI absorption estimates. ",
        "introduction": "Accurate, model independent distances to pulsars are, like many astronomical distance measurements, highly prized but difficult to obtain.  The dispersion measure (DM) of a pulsar indicates the integrated column density of free electrons between the observer and the pulsar, and can be used in conjunction with Galactic electron distribution models to estimate pulsar distances, but this approach suffers from considerable uncertainty when the electron distribution model is poorly understood, such as at high Galactic latitudes. Timing residuals for millisecond pulsars (MSPs) can be used to determine parallax and proper motion \\citep[e.g.][]{hot06}, but only for the very limited subset of pulsars with extremely accurate timing solutions.  VLBI astrometry offers a means to directly measure the parallaxes and proper motions of nearby pulsars.  By obtaining an independent pulsar distance estimate, a pulsar DM can be used to determine the average electron density along the line of sight; an ensemble of such measurements can be used to improve electron distribution models. Whilst several recent large pulsar parallax programs have significantly increased the number of known VLBI parallaxes \\citep[e.g.][]{cha04}, to date very few VLBI parallaxes have been obtained for southern pulsars, due to the bias of VLBI facilities towards the Northern Hemisphere.  As such, models of Galactic electron distributions are more uncertain at far southern declinations.  Furthermore, binary pulsars offer the opportunity to make exceedingly precise tests of gravitational theories, since relativistic effects contribute to the observed rate of change of binary period (\\pbdot), which can be measured very precisely in systems with a stable pulsar.  However, kinematic effects \\citep{shk70} also contribute to \\pbdot, and can only be accurately subtracted to obtain \\pbdot\\ due to General Relativity (GR) if the pulsar distance and transverse velocity are accurately known.  The effect scales with the square of proper motion, and hence is typically largest for nearby pulsars, where VLBI measurements of parallax are most feasible. Since the uncertainty of a DM--based distance estimate is large for any individual pulsar, a reliable estimate of the error on derived GR quantities requires a direct distance determination. Similarly, luminosities for individual pulsars based on DM--inferred distances are questionable, so deriving accurate an accurate luminosity for any individual pulsar generally requires an independent confirmation of distance.  In this paper, we describe our target selection policy in \\S\\ref{sec:select}, and the observations undertaken in \\S\\ref{sec:obs}.  The data reduction pipeline is described in \\S\\ref{sec:datared}, and we present the results of PSR J1559--4438 in \\S\\ref{sec:results}.  We analyse the optimum visibility weighting scheme to use for VLBI astrometry in \\S\\ref{sec:optweight}, and the magnitudes of different sources of systematic errors are investigated in \\S\\ref{sec:errorbudget}. Our conclusions are presented in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} A pipeline for the reduction of LBA astrometric data has been developed in ParselTongue and verified by calculating the parallax ($\\pi = 0.384 \\pm 0.081$\\ mas) and proper motion ($\\mu_{\\alpha} = 1.52 \\pm 0.14$\\ mas yr$^{-1}$, $\\mu_{\\delta} = 13.15 \\pm 0.05$\\ mas yr$^{-1}$) of PSR J1559--4438.  The calculated values are consistent with the DM distance estimate and earlier HI absorption and proper motion studies.  Full account has been made of the impact of residual systematic errors on the quality of the astrometric fit.  The optimal weighting scheme in the presence of systematic errors and varying thermal errors has been investigated, resulting in the guideline that superior astrometric quality can be obtained for the LBA for typical astrometric observations by using equally weighted, as opposed to sensitivity weighted, visibilities if the target can be detected with S/N $>10$.  The completion of this parallax program will result in the publication of 5 more Southern Hemisphere pulsar parallaxes, which will quadruple the number of Southern Hemisphere pulsars with parallaxes determined directly from VLBI astrometry."
    },
    "0808/0808.1284_arXiv.txt": {
        "abstract": "When double neutron star or neutron star-black hole binaries merge, the final remnant may comprise a central solar-mass black hole surrounded by a $\\sim 0.01-0.1~ M_{\\sun}$ torus. The subsequent evolution of this disc may be responsible for short $\\gamma$-ray bursts (SGRBs).  A comparable amount of mass is ejected into eccentric orbits and will eventually fall back to the merger site after $\\sim 0.01$ seconds.  In this {\\it Letter}, we investigate analytically the fate of the fallback matter, which may provide a luminous signal long after the disc is exhausted.  We find that matter in the eccentric tail returns at a super-Eddington rate and is eventually ($\\gsim 0.1$ sec) unable to cool via neutrino emission and accrete all the way to the black hole.  Therefore, contrary to previous claims, our analysis suggests that fallback matter is {\\it not} an efficient source of late time accretion power and is unlikely to cause the late flaring activity observed in SGRB afterglows.  The fallback matter rather forms a radiation-driven wind or a bound atmosphere. In both cases, the emitting plasma is very opaque and photons are released with a degraded energy in the X-ray band. We therefore suggest that compact binary mergers could be followed by an ``X-ray renaissance\", as late as several days to weeks after the merger. This might be observed by the next generation of X-ray detectors. ",
        "introduction": "Close binaries of compact solar-mass objects are expected to form via the evolution of massive star binaries or by dynamical interaction in dense star clusters.  Neutron star (NS--NS) binaries have been detected as radio pulsars \\citep[e.g.][]{faulkner05}, and while black hole--neutron star (BH-NS) or double black hole (BH-BH) binaries have not been observed directly, they are predicted by population synthesis models.  The compact objects are expected to merge due to gravitational wave emission, with evolutionary scenarios estimating a local rate of NS--NS mergers $10-100$ times higher than for BH-NS and BH-BH systems \\citep[e.g.][]{bel07}. The final remnant for NS--NS and NS--BH coalescence is generally thought to be a BH of a few solar masses surrounded by a $0.01-0.1~ M_{\\sun}$ accreting disc \\citep[e.g.][]{ruffert97,shibata03,ross04,faber06}.  The accretion power immediately following the merger is perhaps the ultimate cause of SGRBs \\citep{blinnikov84,eichler89,paczynski91}.  At early times ($\\lsim 0.1 -1$ sec), the accreting disc is geometrically thin, effectively cooled by neutrino emission \\citep{popham99}. When the accretion rate drops below $\\sim 0.1 ~\\msec$ --- the exact value depending on the accretion parameter $\\alpha$ and BH spin (Chen \\& Beloborodov 2007; Metzger, Piro \\&  Quataert 2008) --- the disc becomes radiatively inefficient and super-Eddington accretion drives a substantial outflow \\citep{metzger08}.   During the dynamical phase of the merger, in which the lighter companion is tidally disrupted, a fraction ($\\sim 10^{-2} M_{\\sun}$) of the debris receives enough energy to be ejected from the system while a comparable amount remains bound in eccentric orbits \\citep[e.g.,][]{ross07,faber06} and will eventually return to the disc site: {\\em fallback} matter.  This weakly bound matter may give rise to interesting phenomena observable on timescales longer than any viscous timescale of the disc.  For example, it has been suggested \\citep{lee07,ross07,metzger08} that it can be responsible for the X-ray flaring, observed in SGRB afterglows on timescales of minutes to hours \\citep[e.g.][]{campana06}.  Unfortunately, numerical investigations have not yet been able to follow the long-term ($>$ minutes) evolution of this eccentric tail, because of time-step limitations \\citep{ross07}.   In this {\\em Letter}, we investigate analytically the fate of matter falling back onto a recent merger.  We argue that energy released during fallback is {\\it not} a promising source of the X-ray flares. The energy liberated during fallback will either lead to a powerful, radiation-driven wind or a more gradually expanding ``breeze'' that could ultimately form a bound cloud around the merged object.  In either case, the expanding gas is so opaque that the radiation is trapped in the expanding flow and degraded to low energies before being released in the X-ray band.  We therefore suggest that compact binary mergers might be accompanied by delayed X-ray emission.  We assess the detectability of this emission when the merger is localized by either a short $\\gamma$-ray burst or a gravitational wave signal. A direct observation of the accretion activity would give us valuable information on how compact-object binaries merge.  This {\\it Letter} is organized as follows. We discuss the behavior of the fallback matter in \\S~\\ref{sec:fb}. Then, we consider two possible scenarios for this material as it rebounds: we model a wind in \\S~\\ref{sec:wind} and a bound atmosphere in \\S~\\ref{sec:atmo}. Prospects for detecting the X-ray emission are discussed in \\S~\\ref{sec:detection} and conclusions are drawn in \\S~\\ref{sec:conclu}. ",
        "conclusions": "\\label{sec:conclu}  In this {\\it Letter}, we investigate the possible fate of fallback matter associated with mergers of compact objects, where a disc is formed by disruption of a NS.  Matter flung to highly eccentric orbits will eventually come back to the disc at a super-Eddington rate, converting its kinetic energy into heat via shocks, and will be unable to cool by neutrino emission. Contrary to previous claims, we think that this implies that fallback matter {\\it cannot} accrete all the way to the central object and be responsible for the late energy injections observed in GRB afterglows. Rather, the fallback matter is likely to be blown off the disc plane, leading to the formation of a radiation-driven wind or a bound atmosphere. For the wind case, we have analytically calculated the time evolution of the temperature and luminosity at the trapping radius: while the luminosity decreases (eq.~\\ref{eq:lxwind}), the wind photosphere becomes hotter (eq.~\\ref{eq:txwind}).  At first, the emission is in the EUV band and absorption will likely prevent us from observing it. After one or two weeks, the emission finally peaks in the soft X-ray band and the wind activity can be observed. The bound cloud is radiation pressure-dominated and emits at the Eddington limit in soft X-rays, if the atmosphere does not extend much further than $10^{8}$ cm. We note that our estimates of luminosities are conservative: factors such as a smaller electron fraction in the ejected plasma and moderate geometrical beaming can substantially increase the expected luminosity.  We also discuss detection prospects for this delayed X-ray activity. Our inspection indicates that only in fortuitous circumstances could the X-ray emission be detected with current instruments, while the planned missions (such as Con-X and XEUS) have a better chance \\S~5.1. Then, the main limiting factors will not be the X-ray detector capability, but rather the tool for localizing the merger \\S~5.2. On the one hand, short $\\gamma$-ray bursts can be easily detected and localized in the whole volume where instruments like {\\it Con-X} and {\\it XEUS} can observe the X-ray emission; however, they are estimated to occur at a rate that is $\\sim 100$ times smaller than the rate at which compact binaries merge. On the other hand, the planned advanced gravitational wave interferometers should be able to detect a signal from any such a merger but within cosmic distances smaller than the maximum distance that {\\it Con-X} and {\\it XEUS} can reach. Moreover, X-ray follow up would require better localization precision than currently estimated.  The net result is that between a few to a few tens of detections per year are expected by {\\it XEUS} with a follow-up of a short GRB. Assuming sufficiently good localization, re-pointing after a gravitational signal detection can result in $\\sim 4-54$ wind detection per year from NS-NS mergers, for both {\\it Con-X} and {\\it XEUS}. Furthermore for {\\it XEUS}, there is the exciting possibility to observe X-ray emission from BH-NS mergers: $\\sim 1-53$ event per year. The X-ray emission from these sources should also be brighter than from a NS-NS mergers, since the mass of the central BH could be much larger. The above rates, however, should be taken as indicative of upper limits. We have not taken into account selection effects such as background/foreground sources and the fact that not all BH-NS and NS-NS mergers seem to lead to an accreting system \\citep[e.g.,][]{ross05,bel08}. Moreover, in some cases, the X-ray afterglow from the burst could outshine the wind emission. Nevertheless, the possibility to get information on mergers of compact objects from electromagnetic signals remains, and it could bring important understanding of the physics of these systems.  Finally, our findings have implications for interpreting late time activity observed in GRB afterglows. We consider unlikely that fallback matter can be held responsible, since most of the mass is blown away. Even if $\\sim 10 \\%$ of this matter can accrete all the way to the hole, it is very unlikely that it could produce the observed flares, which have an energy ($\\sim 10^{49}-10^{46}$ ergs) comparable to that of the prompt emission  \\citep[e.g.][]{campana06}. This would require that  the eccentric tail is far more massive than the main disc (contrary to what is observed in simulations) or that the efficiency in converting accreted mass to energy is somehow strongly enhanced in the late fallback accretion. These arguments also apply  to the late accretion from the main disc, which is highly super-Eddington \\citep{metzger08}. We thus conclude that, in general, standard late time accretion is unlikely to account for the phenomena, like flares and plateaux, observed in GRB afterglows."
    },
    "0808/0808.3138_arXiv.txt": {
        "abstract": "We present a coherent theoretical framework for computing gravitational lensing effects and redshift-space distortions in an inhomogeneous universe and investigate their impacts on galaxy two-point statistics. Adopting the linearized Friedmann-Lema\\^\\i tre-Robertson-Walker metric, we derive the gravitational lensing and the generalized Sachs-Wolfe effects that include the weak lensing distortion, magnification, and time delay effects, and the redshift-space distortion, Sachs-Wolfe, and integrated Sachs-Wolfe effects, respectively. Based on this framework, we first compute their effects on observed source fluctuations, separating them as two physically distinct origins: the volume effect that involves the change of volume and is always present in galaxy two-point statistics, and the source effect that depends on the intrinsic properties of source populations. Then we identify several terms that are ignored in the standard method, and we compute the observed galaxy two-point statistics, an ensemble average of all the combinations of the intrinsic source fluctuations and the additional contributions from the gravitational lensing and the generalized Sachs-Wolfe effects. This unified treatment of galaxy two-point statistics clarifies the relation of the gravitational lensing and the generalized Sachs-Wolfe effects to the metric perturbations and the underlying matter fluctuations. For near future dark energy surveys, we compute additional contributions to the observed galaxy two-point statistics and analyze their impact on the anisotropic structure. Thorough theoretical modeling of galaxy two-point statistics would be not only necessary to analyze precision measurements from upcoming dark energy surveys, but also provide further discriminatory power in understanding the underlying physical mechanisms. ",
        "introduction": "\\label{sec:intro} The standard inflationary models with a single inflaton potential predict a nearly perfect Gaussian spectrum of primordial fluctuations \\citep{BASTTU83,STARO82,HAWKI82,GUPI82,MUFEBR92}. Two-point statistics, correlation function in real space and power spectrum in Fourier space, constitutes a complete description of Gaussian random fields, and it has been widely used to understand the physics of the early universe from measurements of the cosmic microwave background and large-scale structure. The recent discovery \\citep{RIFIET98,PEADET99} of the late time acceleration of the universe has spurred extensive investigations of a mysterious energy component with negative pressure, dubbed  dark energy. Observationally, upcoming dark energy surveys will measure galaxy two-point statistics with unprecedented precision from millions of galaxies, constraining the expansion history and the spatial curvature of the universe. Consequently, accurate theoretical modeling of galaxy two-point statistics would be crucial to take full advantage of the promise that these future surveys will deliver.  In achieving this goal, complications arise notably from the nonlinear evolution of matter and scale-dependence of galaxy bias. In this paper we limit ourselves to the linear bias model \\citep{KAISE84} and study the linear theory predictions and its corrections, considering that recent attention has been paid to measuring galaxy two-point statistics in the linear regime (e.g., \\citep{EIBLET05,TEEIST06,PENIET07,PASCET07}). However, measurement precision is often highest on nonlinear scales, and proper modeling of galaxy bias on nonlinear scales can substantially increase the leverage to constrain the underlying physics (see, e.g., \\citep{JIMOBO98,SELJA00,MAFR00,PESM00,SCSHET01,BEWE02,COSH02}).  Further complication arises from the distortion of redshift-space structure by peculiar velocities, which results in anisotropy from otherwise isotropic two-point statistics \\citep{KAISE87,HAMIL92}. The standard practice is to analyze the angle-averaged correlation function or power spectrum, or to construct a linear combination of their multipole components, suppressing the angular dependence of two-point statistics. However, analyzing the full anisotropic structure, though observationally challenging, can utilize additional information that is lost to some degree in the standard practice \\citep{MATSU04,SEEI07,OKMAET08,GACAHU08}.  Gravitational lensing, often assumed to be negligible in galaxy two-point statistics, deflects the propagation of light rays, displacing the position of observed galaxies, and it alters the unit area on the sky and magnifies the observed flux, changing the observed number density of galaxies. The former effect on two-point statistics is to convolve it with the power spectrum of the lensing potential, smoothing out the features in galaxy two-point statistics \\citep{SELJA96}. The latter effect, known as the magnification bias \\citep{NARAY89}, is often used to measure the galaxy-matter cross-correlation function from two source populations separated by large line-of-sight distance \\citep{SCMEET05,BLPOET06}. Recent work \\citep{MATSU00,VADOET07,HUGALO07} showed that these effects on galaxy two-point statistics are non-negligible at the level of accuracy adequate for upcoming dark energy surveys.  However, it is unclear whether this list of additional contributions on galaxy two-point statistics is exhaustive, and what are the contribution terms that are ignored in the standard method but need to be considered if higher accuracy is dictated by observations. Here we present a coherent theoretical framework for computing gravitational lensing effects and redshift-space distortions, and investigate their impacts on galaxy two-point statistics in an inhomogeneous universe. Our treatment generalizes the early work \\citep{MATSU00} and complements the recent work \\citep{VADOET07,HUGALO07}, providing a unified description of galaxy two-point statistics. However, we emphasize that these effects naturally arise from metric perturbations in our approach, comprising a complete and exhaustive set of additional (linear order) contributions to galaxy two-point statistic.  The rest of this paper is organized as follows. In Sec.~\\ref{sec:formalism}, we describe our notation for the Friedmann-Lema{\\^\\i}tre-Robertson-Walker (FLRW) metric and derive the gravitational lensing and the generalized Sachs-Wolfe effects. In Sec.~\\ref{sec:density}, we study their impacts on source galaxy fluctuations and discuss their correspondence to the standard redshift-space distortion and gravitational lensing effect. In Sec.~\\ref{ssec:two}, we derive the observed galaxy two-point statistics in real space and in Fourier space, and we compare the effects of each contribution term on the observed galaxy two-point statistics in Sec.~\\ref{ssec:com}. We conclude in Sec.~\\ref{sec:discuss} with a discussion of the further improvement of our approach. ",
        "conclusions": "\\eneq From Eq.~(\\ref{eq:mgb}), the magnification bias is defined as \\beeq \\dMB(\\Vang,z,f)=5\\left[p(f)-0.4\\right]\\kappa(\\Vang,z), \\eneq with redshift $z$ being the source redshift of the convergence $\\kappa(\\Vang)$ in Eq.~(\\ref{eq:conv}). Considering $\\varepsilon(z)\\simeq V(z)-V(0)$, we call the volume effect in Eq.~(\\ref{eq:dgSW}) as the redshift-space distortion bias, \\bear \\label{eq:zfive} \\zdist(\\Vang,z)&=&-2~{1+z\\over H\\chi}~\\varepsilon -(1+z)H{d\\over dz} \\left({\\varepsilon\\over H}\\right)-\\varepsilon+\\delta V \\nonumber \\\\ &=&-2~{1+z\\over H\\chi}~\\varepsilon +{1+z\\over H}~\\varepsilon~{dH\\over dz} \\nonumber \\\\ &&-{1+z\\over H}~{\\partial\\varepsilon\\over\\partial\\chi} -\\varepsilon+\\delta V. \\enar Note that the generalized Sachs-Wolfe effect $\\varepsilon(z)$ implicitly depends on the direction $\\Vang$ via the line-of-sight velocity $V(z)=v_\\alpha e^\\alpha=\\Vang\\cdot\\bdv{v}(\\Vang,z)$, but it is independent of the limiting flux $f$, provided that galaxies have no velocity bias (i.e., galaxies and matter follow the same velocity field). Finally, the evolution bias is defined from Eq.~(\\ref{eq:zevo}) as \\beeq \\devo(\\Vang,z,f)=-{1+z\\over z}\\left[\\alpha-\\beta \\left({z\\over z_0}\\right)^\\beta\\right]\\varepsilon, \\label{eq:bevo} \\eneq where the directional dependence comes from $\\varepsilon$ and the evolution coefficients $(\\alpha,\\beta,z_0)$ depend on the galaxy sample selected with the limiting flux $f$. While the evolution bias arising from the difference between $\\bar n(z)$ and $\\bar n(\\zobs)$ was recognized \\citep{KAISE87,HAMIL98,MATSU04}, it has been ignored in the literature. However, we show in Sec.~\\ref{sec:two} that the evolution bias can be significantly enhanced. Last, we want to emphasize that equation~(\\ref{eq:summary}) is gauge-invariant as is written in the conformal Newtonian gauge.   \\label{sec:discuss} Galaxy two-point statistics, correlation function in real space and power spectrum in Fourier space, have been extensively used in cosmology to characterize the underlying matter fluctuations. We have presented a coherent theoretical framework based on the linearized Friedmann-Lema{\\^\\i}tre-Robertson-Walker (FLRW) metric for computing the gravitational lensing and the generalized Sachs-Wolfe effects. Within this framework, the metric perturbations are sourced by the underlying matter fluctuations, and they naturally give rise to perturbations in the observable redshift of source galaxies and their angular position on the sky. The time component of the photon geodesic equations can be used to show the former, the generalized Sachs-Wolfe effect \\citep{SAWO67} that generalizes the standard redshift-space distortion by peculiar velocities in a cosmological context, including the Sachs-Wolfe and the integrated Sachs-Wolfe effects. The spatial components of the photon geodesic equations can be used to derive the latter, the gravitational lensing effect that includes the weak lensing distortion, magnification, and time delay effects. This unified treatment provides a complete description of the relation between these seemingly different effects and the underlying matter fluctuations.  Furthermore, it becomes transparent in this treatment how the gravitational lensing and the generalized Sachs-Wolfe effects affect the observed fluctuation field of source galaxies. To the linear order in perturbations, we have computed all the additional contributions to the intrinsic source fluctuation, arising from the gravitational lensing and the generalized Sachs-Wolfe effects. We can gain more insight on the impact of these effects by separating them as two physically distinct origins: the volume and the source effects. The former effect that involves the change of volume is independent of source galaxy populations and hence regardless thereof the volume effect is always present in galaxy two-point statistics. By contraries, the latter effect depends on the intrinsic properties of source galaxy populations and may vanish for a certain population. All of these contributions to the intrinsic source fluctuations result in numerous additional auto and cross terms in the observed galaxy two-point statistics, and therefore proper account should be taken into these additional terms in interpreting measurements of galaxy two-point statistics from upcoming dark energy surveys.  With the complete list of the contributions of the gravitational lensing and the generalized Sachs-Wolfe effects, separated as two physically distinct origins, we have identified several contributions in the volume effect and one contribution in the source effect, which are ignored in the standard treatment: the evolution bias in the source effect arises from the generalized Sachs-Wolfe effect, when the mean number density of sources changes rapidly in redshift, and its impact on the observed galaxy two-point statistics can be substantially larger than that of the gravitational lensing magnification bias. The ignored contributions in the volume effect are typically of order peculiar velocities and hence they are subdominant, compared to the standard redshift-space distortion effect. However, their impact is comparable to the magnification bias at low redshift. While the cross term of the magnification bias and the intrinsic source fluctuation is more important at low redshift than the contribution of the magnification bias itself in the gravitational lensing effect, further calculations of the additional contributions associated with the volume effect may be needed, if higher accuracy of theoretical modeling is required from observation.  We have investigated the impact of the additional contributions to the anisotropic structure of the observed galaxy two-point statistics, after simplifying some of the contributions to the intrinsic source fluctuations. The redshift-space distortion affects the observed galaxy two-point statistics most, imprinting its well-known feature in the anisotropic structure \\citep{KAISE87,STWI95,HAMIL98}. The gravitational lensing effect is small but non-negligible at a percent level, particularly along the line-of-sight separation and at high redshift, since their contribution increases with longer line-of-sight distance to the source galaxies and the clustering amplitude of the intrinsic source fluctuations decreases in redshift. The evolution bias has an angular pattern similar to the redshift-space distortion, but its impact becomes appreciable, only at fairly large transverse separation. While it is challenging to analyze the observed anisotropic structure of galaxy two-point statistics, its full analysis from upcoming dark energy surveys can provide a great opportunity to separately identify each contribution from the gravitational lensing and the generalized Sachs-Wolfe effects, increasing the leverage to understand the underlying physical mechanism.  However, we note that constraining the underlying cosmological model will require not only accurate theoretical predictions, but also model fitting to measurements, which results in further distortion in galaxy two-point statistics, known as Alcock-Paczy\\'nski effect \\citep{ALPA79}. Furthermore, our current investigation has focused on the linear theory predictions and its additional contributions: nonlinearity and scale-dependent galaxy bias can affect our results, though its impact is expected to be less than at the percent level around the acoustic scale (see, e.g., \\citep{EIWH04,SEEI05,EISEWH07}). However, additional leverage can be gained by modeling scale-dependent galaxy bias on nonlinear scales \\citep{YOWEET08}."
    },
    "0808/0808.4004_arXiv.txt": {
        "abstract": "The near future of astrophysics involves many large solid-angle, multi-epoch, multi-band imaging surveys.  These surveys will, at their faint limits, have data on large numbers of sources that are too faint to be detected at any individual epoch.  Here we show that it is possible to measure in multi-epoch data not only the fluxes and positions, but also the parallaxes and proper motions of sources that are too faint to be detected at any individual epoch.  The method involves fitting a model of a moving point source simultaneously to all imaging, taking account of the noise and point-spread function in each image.  By this method it is possible to measure the proper motion of a point source with an uncertainty close to the minimum possible uncertainty given the information in the data, which is limited by the point-spread function, the distribution of observation times (epochs), and the total signal-to-noise in the combined data.  We demonstrate our technique on multi-epoch Sloan Digital Sky Survey (SDSS) imaging of the SDSS Southern Stripe.  We show that with our new technique we can use proper motions to distinguish very red brown dwarfs from very high-redshift quasars in these SDSS data, for objects that are inaccessible to traditional techniques, and with better fidelity than by multi-band imaging alone.  We re-discover all 10 known brown dwarfs in our sample and present 9 new candidate brown dwarfs, identified on the basis of significant proper motion. ",
        "introduction": "There are many multi-epoch imaging surveys in progress or coming up, which will, among other things, deepen our image of the sky and provide information on source variability and proper motions.  These surveys include the SDSS Southern Stripe \\citep{sdssdr7}, the Dark Energy Survey, PanSTARRS, LSST, and SNAP.  These surveys promise proper-motion precisions for well-detected sources on the order of $\\masperyr$ over large parts of the sky.  For context, a typical halo star at a distance of $10~\\kpc$ moving at a transverse heliocentric speed of $100~\\kmpers$ has a proper motion of $2~\\masperyr$, and a typical disk star at $100~\\pc$ and $10~\\kmpers$ has a proper motion of $20~\\masperyr$.  These surveys therefore have the capability of revolutionizing our view of the Galaxy and of the Solar neighborhood.  In most conceptions of a proper-motion measurement, one imagines measuring the position of a source in each of several images, taken at different times.  A linear trajectory is fitted to the positions, relative to some reference frame or set of fixed sources or sources with well measured proper motions.  In its most straightforward form, this method only works for sources bright enough to be detected independently at every epoch---or at least most epochs.  In a multi-epoch survey like the SDSS Southern Stripe, which has $\\sim 70$ epochs \\citep{sdssdr7}, this limits the sources with measured proper motions to a small subset of all sources detectable in the combined data, since the combined data reach $\\sim 2.3~\\mag$ fainter than any individual epoch; for typical source populations this represents increases in population size by factors of $5$ to $25$ at any given signal-to-noise threshold.  In this paper we present a methodology for measuring in multi-epoch imaging the proper motions of sources too faint to detect at any individual epoch.  There are several different technical regimes for these faint-source proper-motion measurements.  In the ``easy'' regime, the sources of interest move a distance smaller than or comparable to the point-spread function width over the duration of the multi-epoch survey.  In this regime, the sources are easy to detect in the co-added image, even without taking account of their proper motions; proper motions can be determined from processing the individual epoch images after detection in the co-added image.  There is a ``difficult'' regime in which the sources of interest move substantially more than the width of the point-spread function over the duration of the survey.  In this regime, the source will not appear at high significance in the co-added image if it does not appear at high significance at any epoch, because its different appearances in the different individual-epoch images do not overlap. In principle, the difficult regime can be addressed by brute force with large computing resources.  In the context of outer Solar-System bodies, brute-force search in the narrow range of expected motions is feasible (for example, \\citealt{bernstein04a, fuentes2008}).  In this paper, we consider only the easy regime.  \\paragraph{Modeling the data:} The traditional method for measuring a stellar proper motion with a set of images taken at different times is as follows: Detect the star at each observed epoch; measure its centroid (by, for example, finding the peak or first moment of the flux) at each observed epoch; and fit a linear motion to the measured positions and times.  This procedure obtains a proper motion, but it puts an unnecessary requirement on the data: that the star be detectable at every epoch.  It also puts an unnecessary burden on the data analyst: it requires decision making about detection and centroiding of the stars at each epoch, decisions that matter at low signal-to-noise, or when faced with data issues such as bad pixels or strong variations in noise from pixel to pixel.  Our new approach is to \\emph{model} all individual-epoch images simultaneously with a single point source that is permitted to have a non-zero parallax and proper motion.  This approach combines the individual-image positional measurement and the determination of the parallax and proper motion, and determines all of these simultaneously by making a statistically ``good'' model of the union of all the data.  In any well-understood imaging survey, each image will have a per-pixel noise model, photometric calibration parameters, and a model of the point-spread function.  In any sufficiently small patch of the sky, if the foreground-subtracted intensity in that patch is dominated by a small number of point sources, it is possible to make an accurate model of all of the pixels in the data set that contribute signal to that small patch.  In this model of the patch, the fluxes, angular positions, parallaxes, and proper motions of the stars in the patch are simply parameter values in the well-fitted models.  In other words, we are assuming that it is possible to model the set of pixels (from all of the images) that contribute to the patch with a $6N$-dimensional model that consists of a set of $N$ moving point sources.  The proper motions determined by image modeling have several advantages over those determined by the traditional method: They require fewer decisions about measurement techniques (although they do require a good model of the data, including point-spread function); they use all of the information in all of the pixels, not just those pixels involved in traditional centroiding; they gracefully handle missing data due to bad pixels or cosmic rays (assuming the bad pixels have been flagged); they require the investigator to make explicit the assumptions about the physical properties of the image and the noise; they can be made to properly propagate pixel-value uncertainties into parameter uncertainties (in this case, proper motion uncertainties); they are the result of optimization of a well-justified scalar objective function (in this case the likelihood). Most importantly for what follows, they can be determined in data sets in which the stars are not well detected at any individual epoch, but only appear in the \\emph{combination} of the images.  In a data set with $\\sim 70$ similar epochs (such as the SDSS Southern Stripe), this corresponds to an increase in the number of available targets by factors of $5$ to $25$ (assuming source populations double to quadruple with each magnitude of depth).  Here we propose, build, test, and use an image-modeling system for the determination of stellar proper motions.  We show that it can work down to low signal-to-noise ratios and that it makes measurements in real data that fully exploit the information available.  We also use it to discover interesting new astrophysical sources.  An approximation to the technique used here has been used previously in the Solar System literature \\citep{bernstein04a}.  \\paragraph{Proper-motion and parallax uncertainties:} Consider a well-sampled image $i$ with a point-spread function of full width at half maximum $\\fwhm_i$.  The signal-to-noise at which the flux of a point source can be measured, $\\sn_i$, is the sum in quadrature of the signal-to-noise contributions from pixels within the point-spread function.  A point source measured with signal-to-noise $\\sn_i$ in a single image can be centroided with (RMS) uncertainty $\\sigma_{\\theta,i}$ of \\begin{equation} \\sigma_{\\theta,i} \\approx \\frac{\\fwhm_i}{\\sn_i} \\quad ; \\end{equation} details such as the shape of the point-spread function introduce factors of order unity \\citep{king1983}.  If we have $N$ such images spanning some time interval, we might hope to obtain a proper motion estimate with uncertainty $\\sigma_{\\mu}$ limited by the point-spread function, the time interval, and the total signal-to-noise \\begin{equation} \\sntotal^2=\\sum_{\\mathrm{images}\\ i}\\sn_i^2 \\label{eq:sntotal} \\end{equation} in the combination of all the images (we have assumed here that the images $i$ are all independent).  The relevant time ``interval'' is not the total time spanned by the data but rather $\\delta_t\\equiv\\sqrt{\\var{t}}$, the standard deviation (root variance) of the times; the best possible proper-motion estimates will have uncertainties \\begin{equation} \\sigma_{\\mu}\\approx\\frac{\\fwhm}{\\delta_t\\,\\sntotal} \\quad , \\label{eq:muerror} \\end{equation} where properly $\\fwhm$ is the square-signal-to-noise weighted mean point-spread function full width at half maximum, and $\\delta_t$ is the square root of the square-signal-to-noise-weighted variance of the times at which the individual epoch images were taken.  By a similar argument, we hypothesize that the best possible parallax estimates will have uncertainties \\begin{equation} \\sigma_{\\pi}\\approx\\frac{\\fwhm}{\\delta_\\lambda\\,\\sntotal} \\quad , \\end{equation} where $\\delta_\\lambda$ is the square root of the square-signal-to-noise-weighted variance of the trigonometric functions of the ecliptic longitude $\\lambda$ of the Sun (time of year in angle units): \\begin{equation} \\delta_\\lambda^2\\equiv\\sigma_{\\cos\\lambda}^2+\\sigma_{\\sin\\lambda}^2 \\quad . \\end{equation} Essentially, $\\delta_\\lambda$ describes how well the parallactic ellipse is sampled; an ideal survey for parallax measurements will have $\\delta_\\lambda\\approx 1$.  Disk stars move with respect to one another at velocities of $\\sim 30\\,\\kmpers$ \\citep{dehnen98a, hogg05a}, that is, on the same order as the velocity of the Earth around the Sun.  In a multi-epoch survey spanning a small number of years (such as the SDSS Southern Stripe), $\\delta_t$ is of order unity, so for disk stars the parallax and proper motion signal-to-noise ratios ought to be comparable in magnitude.  However, most surveys sample ecliptic longitude $\\lambda$ poorly, because of season and scheduling constraints; therefore $\\delta_\\lambda$ is usually substantially less than unity, so the signal-to-noise of parallax is smaller than that of proper motion. ",
        "conclusions": "We have shown that straightforward image modeling permits the measurement of apparent motions, especially the proper motion and parallax of a source in multi-epoch data, even when the source is too faint to be reliably detected or centroided at any individual epoch. The results of this project are not surprising; indeed what is surprising is how rarely the measurements of stellar motions are made by comprehensive data modeling.  We demonstrated the technique on real and artificial data.  In the process of performing these tests we showed that spectrosopically confirmed quasars and brown dwarfs can be perfectly distinguished with proper motions measured by this technique.  Working without proper motions, but with Co-add Catalog sources and a significant amount of near-infrared imaging follow-up, a group has followed up the $z$-only sources most likely to be high-redshift quasars \\citep{chiu08a, jiang08a}.  This project, even after infrared imaging, found---after expensive spectroscopic follow-up---that some of the high-redshift quasar candidates selected on the basis of visible and near-infrared imaging are in fact nearby brown dwarfs.  We have shown that all of these spectroscopically confirmed brown dwarfs have significantly measured ($>5$~sigma) non-zero proper motions by the technique shown here (and are reported in \\tablename~\\ref{tab:movingsources}).  None of the spectroscopically confirmed high-redshift quasars do.  Use of this technique could have been used to substantially increase the efficiency of either quasar or brown-dwarf searches in this data set.  In performing this demonstration, we have independently identified all 10 known brown dwarfs \\citep{fan00a, geballe02a, hawley02a, berriman03a, knapp04a, chiu08a, metchev08a} in our parent sample, and we have discovered 9 \\emph{new} candidate brown dwarfs, presented in \\tablename~\\ref{tab:movingsources}.  Based on our analysis, these objects have a high probability of being brown dwarfs.  It would be desirable to separate disk dwarfs from halo dwarfs---the fastest angular movers tend to be halo members (for example, \\citealt{lepine03a})---but the time cadence of the SDSSSS data is such that parallaxes are not measured well.  Two of the dwarfs we rediscover---2MASS J010752.42+004156.3 and 2MASS J020742.84+000056.4---have previously measured parallaxes \\citep{vrba04a}; the measurements are consistent with our upper limits.  Our tests show that the uncertainty in the proper-motion measurement made by image modeling is consistent with the best possible uncertainties given the angular resolution and photometric sensitivity of the combination of all images in the multi-epoch data set.  These tests effectively show that such measurements can be made for objects that are fainter than those available to traditional methods that require source detection at every epoch.  In imaging with $N$ equally sensitive epochs, we are able to measure objects that are fainter by $\\Delta m$ magnitudes: \\begin{eqnarray} \\Delta m &=& -\\log_{2.5}(\\sn_i) + \\log_{2.5}(\\sntotal) \\\\ &=& \\log_{2.5}(\\sqrt{N}) \\\\ &\\sim& 0.55\\,\\log N~\\mag \\quad . \\end{eqnarray} This advantage amounts to $1~\\mag$ for surveys with $6$ similar epochs, and $1.6$ to $2.0~\\mag$ in data with $20$ to $40$ epochs (such as the data used here).  In the $\\sim 70$ epochs available in SDSS DR7, it reaches $2.3~\\mag$.  Several of the high-redshift quasars and brown dwarfs analyzed in this study were only detectable in the combination of all of the multi-epoch images.  The depth advantage of image modeling is most dramatic in surveys with very large numbers of epochs, as is expected for LSST.  In general the number of interesting sources is a strong function of depth (factors of $2$ to $4$ per magnitude), so the ``reach'' of the image-modeling technique is a strong function of the number of epochs.  One limitation of the work presented here is that we used the zero-proper-motion image ``stack'' for source detection and therefore will only have in the candidate list objects with small proper motions.  Faint stars and dwarfs with proper motions large enough that they move the width of the PSF between epochs, or some significant fraction of that, are harder to find, because they don't appear in the stack at much higher signal-to-noise than they appear in any individual-epoch image.  In future work we hope to address the detection and measurement of these fast-moving but very faint sources. Approximations have been executed in the search for Solar System bodies (for example, \\citealt{bernstein04a}).  Certainly a reliable system for discovery in this regime would have a big impact on future surveys like PanSTARRS and LSST."
    },
    "0808/0808.3630_arXiv.txt": {
        "abstract": "We present results of a periodicity search of 20 intra-day variable optical light curves of the blazar S5 0716$+$714, selected from a database of 102 light curves spanning over three years. We use a wavelet analysis technique along with a randomization test and find strong candidates for nearly periodic variations in  eight light curves, with probabilities ranging from 95\\% to $>$99\\%. This is the first good evidence for periodic, or more-precisely, quasi-periodic, components in the optical intra-day variable light curves of any blazar. Such periodic flux changes support the idea that some active galactic nuclei variability, even in blazars, is based on accretion disk fluctuations or oscillations. These intra-day variability time scales are used to estimate that the central black hole of the blazar S5 0716$+$714 has a mass $> 2.5 \\times 10^6$ M$_{\\odot}$. As we did not find any correlations between the flux levels and intra-day variability time scales, it appears that more than one emission mechanism is at work in this blazar. ",
        "introduction": "Blazars are the subclass of radio-loud active galactic nuclei (AGNs) consisting of BL Lacertae objects (BL Lacs) and flat spectrum radio quasars (FSRQs). BL Lacs show largely featureless optical continua.  All blazars exhibit strong flux variability on diverse time scales varying from a few minutes to many years at all accessible wavelengths of the electromagnetic spectrum. For convenience, blazar variability can be broadly divided into three classes, viz., intra-day or intra-night variability, short-term variability and long-term variability. Variations in flux of up to a few tenths of a magnitude over the course of a night or less is variously known as intra-day variability (IDV) (Wagner \\& Witzel 1995), which is the term we adopt here, microvariability (e.g., Miller, Carini \\& Goodrich 1989), or intra-night optical variability (e.g., Gopal-Krishna, Sagar \\& Wiita 1995) . Short- and long-term variabilities can amount to four or five magnitudes and are usually defined to have time scales from weeks to several months and several months to years, respectively.  It is widely accepted that the central engines of AGNs fundamentally are comprised of super massive black holes (SMBHs) along with the radiating matter accreting onto them.  In the case of radio-loud AGN such as blazars, jets emerge from these central engines. IDV is most likely produced in close proximity to the SMBH (e.g., Miller et al.\\ 1989), although an origin further away within a turbulent jet (e.g., Marscher, Gear \\& Travis 1992) is also possible. Minimum IDV time scales of blazars can be used to place constraints on sizes of the emitting regions, and, if the radiation is indeed emitted close to the center, on the masses of the SMBHs. Detection of periodic or quasi-periodic oscillations in optical IDV light curves of blazars would be strong evidence for the presence of a single dominant orbiting hot-spot on accretion disk, or accretion disk pulsation (e.g., Chakrabarti \\& Wiita 1993; Mangalam \\& Wiita 1993). Hence the search for periodic or quasi-periodic oscillations in the IDV light curves of blazars is of great interest; however, such variations have not yet been well established. The motivation of the present research program is to see if such periodic IDV patterns do indeed exist in the optical light curves of blazars.  For the present work, we extracted data from the published literature for optical IDV of the BL Lac object S5 0716$+$714. It is one of the brightest and therefore most well studied BL Lacs with respect to optical IDV. The optical continuum of this blazar is so featureless that many attempts made to determine its redshift have failed. The non-detection of its host galaxy allowed a lower limit to its redshift, $z > 0.3$ (Wagner et al.\\ 1996), to be set and a later non-detection led to a claim that $z > 0.52$ (Sbarufatti, Treves \\& Falomo 2005).  However, there has been a very recent claim of a host galaxy detection which produces a ``standard candle'' value of $z = 0.31 \\pm 0.08$ (Nilsson et al.\\ 2008). The duty cycle of the source is essentially unity, which implies that the source is always in at least a moderately active state (Wagner \\& Witzel 1995). In the search for optical variability on diverse time scales in this source, it has been observed intensely over at least the last 15 years (e.g.\\ Wagner \\& Witzel 1995; Wagner et al.\\ 1996; Heidt \\& Wagner 1996; Ghisellini et al.\\ 1997; Sagar et al.\\ 1999; Villata et al.\\ 2000; Qian, Tau \\& Fan 2002; Raiteri et al.\\ 2003; Wu et al.\\ 2005, 2007; Nesci et al.\\ 2005; Stalin et al.\\ 2006; Gu et al.\\ 2006; Montagni et al.\\ 2006; Ostorero et al.\\ 2006; Foschini et al.\\ 2006; Villata et al.\\ 2008; Gupta et al.\\ 2008 and references therein). Five major optical outbursts were reported for this source; they occurred at the beginning of 1995, in late 1997, in the fall of 2001, in March 2004 and in the beginning of 2007 (Raiteri et al.\\ 2003; Foschini et al.\\ 2006; Gupta et al.\\ 2008). These five outbursts indicate a timescale of long-term variability of $\\sim 3.0 \\pm 0.3$ years (e.g., Raiteri et al.\\ 2003). Correlated radio/optical short-term variability measured over a month seemed to reveal some quasi-periodicities for this blazar (Quirrenbach et al.\\ 1991; Wagner et al.\\ 1996).  This paper is structured as follows. In \\S 2 we discuss the source of the data we use and we describe the criteria employed to select the  light curves for this study. In \\S 3 we describe the wavelet analysis and randomization technique we employ in our analysis. Our results are presented in \\S 4  and we discuss them and draw conclusions in \\S 5. ",
        "conclusions": "There are two major classes of models for AGN variability, those involving shocks-in-jets and those invoking instabilities on or above accretion disks. The former is expected to dominate in blazars and the latter is most important when jets are absent or weak (e.g., Wagner \\& Witzel 1995; Urry \\& Padovani 1995; Mangalam \\& Wiita 1993). In another jet-based variant it is argued that for low-luminosity AGNs, the accretion disk is radiatively inefficient and any small amplitude variation, even in the low-state of the source, will be due to a weak jet (e.g., Chiaberge et al.\\ 1999, 2006; Capetti et al.\\ 2007, and references therein). This variant may be able  to  explain optical IDV in blazars. The detection of quasi-periodicity or periodicity on IDV time scales can be most easily explained by the presence of a single dominating hot-spot on the accretion disk (e.g., Mangalam \\& Wiita 1993; Chakrabarti \\& Wiita 1993) or perhaps by pulsational modes in the disk (e.g., Espaillat et al.\\ 2008).  Together, the difficulties in obtaining lengthy, high quality, and essentially evenly spaced data for ground based observations of AGN, and the weaknesses of the standard analysis tools (periodograms and Fourier transforms) under these circumstances, has meant that performing good searches for periods or quasi-periods in blazars has been difficult until recently.  So it is not surprising that there have been few strong claims of quasi-periodic variations in blazar LCs. Carini et al.\\ (1992) noted a weak peak in the power spectrum of optical IDV observations of OJ 287 at $\\approx$ 32 minutes, but it  was not statistically significant. In the same blazar, Carrasco et al.\\ (1985) earlier claimed periodic variations on timescales of tens of minutes but they were not convincing. Earlier excellent optical data of the blazar 0716$+$714 has shown possible quasi-periodicity  on the timescales of $\\sim$ 1 day, 4 days and 7 days, and some of these were in apparent synchrony with radio variations, but the time series were too short to provide conclusive evidence (Wagner 1992; Heidt \\& Wagner 1996). Very recently, using X-ray data for the blazar 3C 273, Espaillat et al.\\ (2008) have reported quasi-periodicity on a timescale of 3.3 ks.  If these fluctuations represent orbital periods they would imply central BH masses for 3C 273 some 10 to 100 times smaller than independent determinations have found, so Espaillat et al.\\ (2008) favor the hypothesis that the variations they have seen arise from higher-order modes within the disk. To our knowledge, no other members of the blazar class have shown significant harmonic components in the IDV or short-time scale LCs.  Here we have used a wavelet plus randomization technique (Linnell Nemec \\& Nemec 1985; O'Shea et al.\\ 2001) to search for optical IDV time scales of the blazar S5 0716$+$714. We selected 20 nights of IDV data that were characterized by long strings of approximately uniformly sampled data with high IDV signals and low noise from an initial database of 102 nights of observations by Montagni et al.\\ (2006).  We found high probabilities ($ \\ge 95$\\%) of (at least quasi-) periodic oscillations in eight of these IDV LCs and very high probabilities ($ \\ge 99$\\%) for five of them. This is the first evidence of detection of periodic components in blazar optical IDV LCs. We found quasi-periodic IDV time scales between $\\sim$ 25 and $\\sim$ 73 minutes.  These variability timescales lead to nominal BH  masses ranging from 2.47 -- 7.35 $\\times$ 10$^{6}$ M$_{\\odot}$ assuming the period arises from a Schwarzschild BH, while for a rapidly rotating Kerr BH these mass estimates rise to 1.57 -- 4.67 $\\times$ 10$^{7}$ M$_{\\odot}$. However, if these variations arise from internal disk modes then the corresponding BH mass can be substantially larger, by factors of $\\sim$10 to over 100.   And if they do emerge from an inner portion of the  jet where they are induced by disk fluctuations, the masses would be increased by a factor of $\\delta \\sim 20$; however, if they emerge from jets at large distances from the SMBH, then clearly no information on the BH mass can be extracted from these timescales.  We did not find any correlation between average flux and IDV time scales, which implies that optical emission from the blazar is probably not governed by a single mechanism. Our results appear to be most consistent with models in which accretion disk variability plays a role in the optical emission of blazars. Nonetheless, it is possible that emission from turbulet jet, which is also precessing or swinging (e.g., Camenzind \\& Krockenberger 1992; Gopal-Krishna \\& Wiita 1992) could also explain these observations."
    },
    "0808/0808.1918_arXiv.txt": {
        "abstract": "We have modelled the observed color-magnitude diagram (CMD) at one location in M33's outskirts under the framework of a simple chemical evolution scenario which adopts instantaneous and delayed recycling for the nucleosynthetic products of Type II and Ia supernovae. In this scenario, interstellar gas forms stars at a rate modulated by the Kennicutt-Schmidt relation and gas outflow occurs at a rate proportional to the star formation rate (SFR). With this approach, we put broad constraints on the role of gas flows during this region's evolution and compare its [$\\alpha$/Fe] vs.\\ [Fe/H] relation with that of other Local Group systems. We find that models with gas inflow are significantly better than the closed box model at reproducing the observed distribution of stars in the CMD. The best models have a majority of gas inflow taking place in the last 7 Gyr, and relatively little in the last 3 Gyr. These models predict most stars in this region to have [$\\alpha$/Fe] ratios lower than the bulk of the Milky Way's halo. The predictions for the present-day SFR, gas mass, and oxygen abundance compare favorably to independent empirical estimates. Our results paint a picture in which M33's outer disc formed from the protracted inflow of gas over several Gyr with at least half of the total inflow occurring since $z \\sim 1$. ",
        "introduction": "\\label{sec:intro}  The distribution of stars in a color-magnitude diagram (CMD) is a sensitive function of the past and present star formation rate and chemical composition. This fact allows the extraction of the star formation history (SFH) and chemical enrichment history (CEH) of a galaxy by fitting its observed CMD with a model CMD whose SFH and CEH are known. The results of this synthetic CMD fitting method can ultimately yield insights into the processes shaping galaxy evolution, such as gas flows, energetic feedback, satellite accretion, tidal interactions, and dynamical mixing. Nevertheless, this method is only one of many employed in the domains of near-field cosmology and galactic paleontology \\citep{Freeman02}.  Another, equally important method, is chemical modelling, whose primary goal is to reproduce the elemental abundance distribution in a galaxy's stars. This method involves tracking the effects of gas flows, star formation, and stellar nucleosynthesis on the abundances and abundance ratios of certain elements. For example, the abundance of the $\\alpha$-elements (O, Ne, Mg, Si, S, Ca, Ti) relative to iron, commonly plotted as [$\\alpha$/Fe] vs.\\ [Fe/H], is an extremely useful tool to discriminate between different models for a galaxy's evolution because the $\\alpha$-elements and iron have different production sites and time-scales.  The $\\alpha$-elements are produced mainly in the hydrostatic burning phases and explosive deaths (supernovae Type II, or SNe II) of massive stars ($M \\ga 8\\ M_{\\sun}$), which have short lifetimes ($\\la 10^7$ yr). The majority of iron, on the other hand, is created in supernovae Type Ia (SNe Ia), which are the thermonuclear explosions of carbon-oxygen white dwarfs in binary systems and, therefore, typically occur on a much longer time-scale of $\\sim 1$ Gyr \\citep{Greggio05}.  The methods of CMD fitting and chemical modelling are complementary, but they have evolved mostly independently. Some examples in the recent literature which combine the two methods generally use the results of CMD fitting as external inputs to chemical evolution models. For example, \\citet{Aparicio97} derived the first SFH of the Local Group (LG) dwarf galaxy, LGS 3, from CMD fitting and used the chemical evolution equations to see how much gas inflow and outflow was required to match their results. \\citet{Carigi02} computed the chemical evolution of four dwarf spheroidal (dSph) satellites of the Milky Way (MW).   As external constraints, these authors used the SFHs of these galaxies derived by \\citet{Hernandez00} from a CMD fitting method. \\citet{Dolphin02} calculated the SFHs of 6 dSph MW satellites using synthetic CMD fitting, and then \\citet{Lanfranchi03,Lanfranchi04} adopted his SFHs as inputs to their chemical evolution models for the same galaxies. Similarly, \\citet{Fenner06} modelled the chemical evolution of the Sculptor dSph using the SFH derived by \\citet{Dolphin05} from a CMD analysis. Lastly, after adopting the SFH derived by \\citet{Wyder01,Wyder03}, \\citet{Carigi06} modelled the evolution of NGC 6822 in a cosmological framework.  Conversely, the results of chemical modelling can be used as external inputs to CMD fitting.  For example, \\citet[][hereafter PT98]{Pagel98} modelled the chemical evolution of the Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC), adopting gas inflow and non-selective galactic winds (i.e., the wind efficiency was the same for all chemical elements). These authors tuned the model parameters to match the observed elemental abundances of clusters and supergiant stars in these systems. Some subsequent studies adopted the LMC's age-metallicity relation (AMR) derived by \\citetalias{Pagel98} to extract its SFH from the CMD and luminosity function \\citep[e.g.,][]{Holtzman99,SmeckerHane02}.  All the aforementioned studies combined CMD fitting and chemical modelling in a two-step process, in which the first step was done independently of the second. However, the two steps are inextricably linked because an increase in the star formation rate (SFR) speeds up the chemical enrichment. Gas flows into or out of the system can change this coupling making them important ingredients to include in any model. Therefore, one drawback of using the CMD-derived SFHs as independent constraints on the chemical evolution models is that the SFHs are not necessarily physically self-consistent under the action of processes like nucleosynthesis, galactic winds, and gas accretion, processes which are fundamental to the chemical models themselves.  Similarly, using the results of chemical evolution models as independent constraints on CMD fitting is not necessarily self-consistent because the models have a particular form for the SFH, which is the very thing the CMD fitting is supposed to derive.  The recent works of \\citet{Ikuta02} and \\citet{Yuk07} represent a significant improvement over the studies mentioned above because they incorporate CMD fitting and chemical modelling simultaneously. \\citet{Ikuta02} computed a few closed box evolution models for several dSph satellites of the MW. They performed a qualitative comparison of their model CMD and [Mg/Fe] vs.\\ [Fe/H] relation with what was observed and found a reasonable agreement, but they had to invoke gas stripping via ram pressure or tidal shocks to reconcile the present day gas fraction of their closed box models ($\\sim 97\\%$) with the observed values ($\\sim 0\\%$). \\citet{Yuk07} improve upon the work of \\citet{Ikuta02} by quantitatively fitting a closed box chemical model to the CMD of IC 1613, a relatively isolated LG dwarf irregular galaxy. Their model SFH and AMR are in good agreement with previous independent determinations based on the canonical CMD fitting method, lending support to both the old and new methods. Moreover, their predicted present-day oxygen abundance is consistent with the observed value.  There are many lines of evidence that suggest not all galaxies evolve as closed boxes. As originally hypothesized by \\citet{Larson74} and exemplified by the \\citet{Ikuta02} results, gas outflows could explain the lack of gas in dSphs despite their low metallicities (see also \\citet{Lanfranchi04}). Other evidence for gas outflows includes the abundances of metals in the IGM \\citep{Edge91,Mushotzky97}, the mass-metallicity relation of galaxies at low and high redshift \\citep[e.g.,][]{Garnett02,Tremonti04,Pilyugin04,Erb06}, extended extra-planar HI gas in spirals \\citep{Fraternali04,Oosterloo07}, high velocity clouds (HVCs) around the MW with near solar metallicity \\citep{VanWoerden04,Wakker01,Richter01}, extra-planar optical and X-ray emission around starburst galaxies \\citep{Heckman90,Lehnert96,Martin02,Strickland04}, and velocity shifts of high-ion absorption lines in damped Lyman-$\\alpha$ systems \\citep{Fox07} and Lyman break galaxies at $z \\sim 3 - 4$ \\citep[e.g.,][]{Pettini01,Adelberger03}.  For a more detailed review, we refer the interested reader to \\citet{Veilleux05}.  Evidence for gas inflows includes the so-called G-dwarf problem, which is the fact that the metallicity distribution function of low mass, long lived stars observed in the Solar vicinity (SV) and in many other galaxies is too narrow and contains too few metal poor stars compared to the closed box model \\citep[e.g.,][]{Tinsley75,RochaPinto96,Seth05,Jorgensen00,Wyse95, Harris02,Koch06,Sarajedini05,Mouhcine05,Worthey05}. Additionally, gas inflow over several Gyr is required to explain many characteristics of the SV, including stellar chemical compositions and the present-day gas mass fraction and SFR \\citep[e.g.,][]{Chiappini01,Portinari98}.  Other evidence for inflows includes the kinematics of extra-planar HI gas in some spirals \\citep{Fraternali08}, the low metallicities and large distances of some MW HVCs \\citep{Wakker99,Richter01}, high velocity OVI absorption along various sight-lines through the MW's halo \\citep{Sembach03}, and the prevalence of warps in HI discs \\citep{Bosma91}. We refer the reader to \\citet{Sancisi08} for a review of observational evidence for gas inflow.  There are theoretical expectations for gas inflows, as well. Cosmological simulations of galaxy formation in the $\\Lambda$CDM framework predict disc galaxies to form from the accretion of gas after the last major merger \\citep{Abadi03,SommerLarsen03,Governato04,Governato07}. For galaxies with virialized dark halo masses $\\la 10^{12} M_{\\sun}$, the accreted gas is expected to be mostly cold ($T \\sim 10^{4-5} K$). On the other hand, upon entering the haloes of more massive galaxies, most of the accreted gas experiences shock-heating to the virial temperature ($T \\ga 10^6 K$), creating surrounding reservoirs of hot gas \\citep{Keres05,Dekel06}. Thermal instabilities lead to the condensation of cooler clouds, which then rain down on the discs \\citep{Maller04,SommerLarsen06,Kaufmann06}. In some of the most recent N-body simulations, these clouds have properties similar to the HVCs mentioned above \\citep[][]{Peek08}. Furthermore, this idea is supported by the recent discovery of a hot ($T \\sim 10^6 K$) gaseous halo around the quiescent massive spiral, NGC 5746 \\citep{Pedersen06}.  This gas is too hot to be heated by SNe in the disc and the disc SNe rate is too low to have created the reservoir through the outflow of disc gas. Finally, theoretical simulations suggest the need for extended accretion of dilute gas to keep discs from being destroyed after a succession of minor mergers with mass ratios of 4:1 and even up to 10:1 \\citep{Bournaud07}.  Despite all this evidence, the precise nature and importance of gas flows in the evolution of galaxies, particularly spirals, is still uncertain. To help improve the situation, in the present paper, we develop a chemophotometric CMD fitting method as an extension to the canonical method used in many other studies, but with the goal of examining the role of gas flows in M33's evolution.  Following \\citet{Ikuta02}, we solve the chemical evolution equations to obtain a self-consistent SFH and AMR. We improve upon their work by allowing for gas inflow and outflow and by efficiently searching the full volume of parameter space to make a detailed and quantitative fit to the observed CMD.  The photometric data we use in this paper and the reduction procedure were presented in \\citet[][hereafter \\citetalias{Barker07a}]{Barker07a}. In summary, three co-linear fields located in projection $\\sim 20 - 30\\arcmin$ southeast of M33's nucleus were observed with the Advanced Camera for Surveys on board the {\\it Hubble Space Telescope}. In \\citet[][hereafter \\citetalias{Barker07b}]{Barker07b}, we computed the SFHs for these fields using the canonical synthetic CMD fitting method with age and metallicity as free parameters. Because the outer two fields may have a non-negligible contribution from M33's halo or thick-disc (see \\citetalias{Barker07b} for a discussion), we restrict ourselves to analyzing the innermost field in the current study. This field has a projected galactic area of $\\sim 0.7\\ \\rm kpc^2$ and it lies at a deprojected radius of $R_{dp} \\sim 6$ disk scale lengths \\citep{Ferguson07} or $\\sim 9$ kpc, assuming a distance of 867 kpc \\citep{Tiede04}, inclination of $56\\degr$, and position angle of $23\\degr$ \\citep{Corbelli97}. At this radius in M33, the azimuthally averaged HI and stellar surface densities are, respectively, $\\sim 3\\ \\rm M_{\\sun}\\ pc^{-2}$ and $\\sim 0.3\\ \\rm M_{\\sun}\\ pc^{-2}$ \\citep{Corbelli97,Corbelli03}.  In the next section, we outline the chemical evolution equations, which form the backbone of our models. We describe in \\S \\ref{sec:method} how we link these equations with the synthetic CMD fitting method to build a self-consistent model CMD. In \\S \\ref{sec:chemresults}, we present the results of fitting closed box and inflow/outflow models to the data. We explore the effects of varying the model parameters in \\S \\ref{sec:var}. In \\S \\ref{sec:obs}, we test the model predictions with independent observations and, in \\S \\ref{sec:chemdisc}, we compare the [$\\alpha$/Fe] vs.\\ [Fe/H] relation in M33 with other LG systems. Finally, we summarize our results in \\S \\ref{sec:chemconc}. ",
        "conclusions": "\\label{sec:chemconc}  In the framework of a simple chemical evolution scenario which adopts instantaneous and delayed recycling for the nucleosynthetic products of Type II and Ia SNe, we have modelled the observed CMD in one location in M33's outskirts. In this scenario, interstellar gas forms stars at a rate modulated by the KS relation and gas outflow occurs at a rate proportional to the SFR. Compared to the common CMD fitting method of allowing age and metallicity to be free parameters, this scenario yields a more physically self-consistent SFH with fewer free parameters and makes more predictions that can be tested with independent observations. Moreover, this method puts broad constraints on the role of gas flows and extends the method of chemical fingerprinting to stellar systems, like M33, that are beyond the reach of current high resolution spectrographs.  The precise details of our results depend on which stellar evolutionary tracks are used, Padova or Teramo. Nevertheless, the broad trends appear to be robust. We found that, when the star formation efficiency, $\\epsilon$, is constant, the canonical closed box model fails to reproduce the observed distribution of stars in the CMD. Models with an exponentially declining IFR from 14.1 Gyr ago exhibit similar discrepancies, but to a lesser magnitude.  Instead, the inflow models which best reproduce the observed CMD have a significant fraction ($\\ga 50\\%$) of gas inflow taking place in the last 7 Gyr and a smaller fraction ($< 10\\%$) taking place within the last 3 Gyr. This leads to a SFH in overall agreement with what we found in \\citetalias{Barker07b}, and suggests the traditional method of synthetic CMD fitting can be physically self-consistent.  Allowing $\\epsilon$ to vary with time, as might be expected if the star formation time-scale depends on local ISM properties, significantly improves the closed box model if the initial metallicity is [Fe/H] $\\sim -1.3$, but the resulting present-day gas mass surface density is too high compared to the observed value. A model with a varying $\\epsilon$ and constant, non-zero IFR reproduces the CMD and gas mass about as well as the best constant $\\epsilon$ models. In this case, more inflow can take place in the last 3 Gyr, but there is still a significant fraction ($\\sim 50\\%$) of inflow occurring over the last 7 Gyr. The amount of variation in $\\epsilon$ required by these models, however, could be larger than the intrinsic scatter in the KS relation of other galaxies.  We also examined the [$\\alpha$/Fe] vs.\\ [Fe/H] relation, a key diagnostic of the evolutionary history of stellar systems. Like the MW's dwarf satellites, the bulk of stars at this location in M33's outskirts have [$\\alpha$/Fe] ratios lower than the MW's halo field. Stars formed over 8 Gyr ago have $\\langle$[$\\alpha$/Fe]$\\rangle$ $\\sim 0.2 \\pm 0.1$ and stars forming today have $\\langle$[$\\alpha$/Fe]$\\rangle$ $\\sim -0.1 \\pm 0.1$. The mean [$\\alpha$/Fe] ratio of all stars ever formed at this location in M33 was found to be $\\sim 0.1 \\pm 0.1$. From the tests we have conducted, we estimate that the systematic errors of these values are $\\sim \\pm 0.2$ dex, comparable to that of some high resolution spectroscopic measurements of other systems \\citep[e.g.][]{Monaco05,Geisler07}.    Our results paint a picture in which M33's outer disc formed from the protracted inflow of gas over several Gyr with at least half of the total inflow occurring since $z \\sim 1$, but relatively little since $z \\sim 0.25$. All the acceptable inflow models have a similar IFR $\\sim 2 \\times 10^{-4}\\ \\rm M_{\\sun}\\ yr^{-1}\\ kpc^{-2}$ averaged over the lifetime of the Universe. This estimate is uncertain by at least a factor of 2 and the IFR could have intrinsic variations of a factor of $\\sim 3$ when averaged over 2 $-$ 3 Gyr time-scales. Nevertheless, it still gives a rough indication of the mean gas IFR that could occur in the outskirts of low mass spiral galaxies.  A useful baseline for comparison is the lifetime average IFR in the SV, predicted by \\citet{Colavitti08} to be $\\sim 4 \\times 10^{-3}\\ \\rm\\ M_{\\sun}\\ yr^{-1}\\ kpc^{-2}$. This estimate comes from several chemical evolution models which adopt different shapes for the IFH and which reproduce many local observables, like the abundances of various elements, MDF of long-lived stars, and present-day gas fraction, total mass density, SFR, IFR, and SN rate. That the mean IFR in the SV has been higher than in M33's outer disc is not surprising since the MW is more massive than M33 and the SV is at a much smaller radius (in terms of disc scale lengths) than the M33 field we have studied.  \\citet[][]{SommerLarsen03} ran an N-body cosmological simulation and selected 12 dark matter haloes for more detailed simulations which included baryons. Two Milky Way-type spirals, S1 and S2, that resulted from the simulations experienced gas accretion at a mean rate of $\\sim 10^{-3}\\ M_{\\sun} \\rm\\ yr^{-1}\\ kpc^{-2}$ at radii of $6 - 7$ disc scale lengths. This rate is about 5 times larger than the mean IFR we have found at a similar location in M33's outer disc, but S1 and S2 were about 5 times more massive than M33. The total disc averaged accretion rates of the least massive simulated galaxies, with rotation velocities comparable to M33's, were $\\sim 15 - 35\\%$ smaller than those of S1 and S2 at $z = 0$. Therefore, if the {\\it outer} disc accretion rate scales with a galaxy's rotation velocity in the same way as the {\\it total} disc averaged rate, and if this scaling holds for all redshifts, then our results are consistent with the simulations of \\citet[][]{SommerLarsen03}.  What is the source and nature of the gas inflow? If it occurs in the disc plane, it could be caused by spiral density waves, viscosity in the disc gas, or gas with lower angular momentum falling onto the disc at larger radii and flowing toward M33's nucleus \\citep{Lacey85b,Lin87,Bertin96b,Portinari00,Roskar08}. Gas flowing from above the disc plane could come from the condensation of a hot halo corona as described in \\S \\ref{sec:intro}. However, such a process is expected to be more efficient in a massive spiral like the MW than in a late-type spiral like M33 \\citep{Dekel06}.  The condensed clouds predicted by numerical simulations to fall onto a disc galaxy have properties similar to the MW's HVC population \\citep{Peek08}. Recent surveys of HI emission around M31 and M33 have revealed an analogous HVC population around these galaxies \\citep{Thilker04,Braun04,Grossi08}. \\citet{Thilker04} detected 20 clouds within 50 kpc of M31's disc and the maps presented by \\citet{Grossi08} show many within $\\sim 20$ kpc of M33 with an inferred total mass $\\geq 5 \\times 10^7\\ M_{\\sun}$. The M31/33 clouds appear to have a mean mass $\\sim 10^5\\ M_{\\sun}$ and size $\\sim 1$ kpc \\citep{Westmeier05,Grossi08}. The lifetime integrated inflow mass of $\\sim 10^6\\ M_{\\sun}$ in the M33 region we have studied requires the accretion of $\\sim 10$ such clouds. Extrapolation to M33's entire disc is risky given how little we know about its evolution and the population of clouds around M33. In addition to condensed halo gas, the clouds could also be low-mass dark matter subhaloes that never formed stars or the tidal debris from recent mergers and interactions \\citep{Thilker04,Grossi08}. In the future, we plan to apply the techniques presented in this paper to other regions of M33 and other galaxies. This approach can provide insights on the origin and amount of accreted gas and coupling (or lack thereof) between the baryonic and dark matter accretion histories of these systems."
    },
    "0808/0808.3295_arXiv.txt": {
        "abstract": "Five planets are known to orbit the star 55 Cancri.  The recently-discovered planet {\\it f} at 0.78 AU (Fischer \\etal 2008) is located at the inner edge of a previously-identified stable zone that separates the three close-in planets from planet {\\it d} at 5.9 AU.  Here we map the stability of the orbital space between planets {\\it f} and {\\it d} using a suite of n-body integrations that include an additional, yet-to-be-discovered planet $g$ with a radial velocity amplitude of 5 $m \\, s^{-1}$ (planet mass = 0.5-1.2 Saturn masses).  We find a large stable zone extending from 0.9 to 3.8 AU at eccentricities below 0.4. For each system we quantify the probability of detecting planets $b-f$ on their current orbits given perturbations from hypothetical planet $g$, in order to further constrain the mass and orbit of an additional planet.  We find that large perturbations are associated with specific mean motion resonances (MMRs) with planets $f$ and $d$.  We show that two MMRs, 3f:1g (the 1:3 MMR between planets $g$ and $f$) and 4g:1d cannot contain a planet $g$. The 2f:1g MMR is unlikely to contain a planet more massive than $\\sim 20 \\mearth$.  The 3g:1d and 5g:2d MMRs could contain a resonant planet but the resonant location is strongly confined.  The 3f:2g, 2g:1d and 3g:2d MMRs exert a stabilizing influence and could contain a resonant planet.  Furthermore, we show that the stable zone may in fact contain 2-3 additional planets, if they are $\\sim 50 \\mearth$ each. Finally, we show that any planets exterior to planet {\\it d} must reside beyond 10 AU. ",
        "introduction": "In a remarkable study, Fischer \\etal (2008) have measured the orbits of five planets orbiting the the star 55 Cancri, the most planets of any exoplanet system to date.  The system contains two strongly-interacting, near-resonant giant planets at 0.115 and 0.24 AU (Butler \\etal 1997; Marcy \\etal 2002), a 'hot Neptune' at 0.038 AU (McArthur \\etal 2004), a Jupiter analog at 5.9 AU (Marcy \\etal 2002) and a newly-discovered sub-Saturn-mass planet at 0.78 AU (Fischer \\etal 2008).  Table 1 lists Fischer \\etal's self-consistent dynamical fit of the orbits of the five known planets in 55 Cancri.  The fast-paced nature of exoplanet discoveries can lead to interesting interactions between theory and observation.  Prior to the discovery of planet 55 Cancri {\\it f}, several groups had mapped out the region between planets $c$ and $d$ to determine the most likely location of additional planets.  Most studies used massless test particles to probe the stable zone (Barnes \\& Raymond 2004 -- hereafter BR04; Jones, Underwood \\& Sleep 2005; Rivera \\& Haghighipour 2007).  Test particles are good proxies for small, Earth-sized planets because they simply react to the ambient gravitational field. However, they are not good substitutes for fully-interacting, real planets. Thus, Raymond \\& Barnes (2005; hereafter RB05) mapped out this zone using Saturn-mass test planets.  The stable zone from BR04 and RB05 extended from 0.7 AU to 3.2-3.4 AU, a region that includes the star's habitable zone (Raymond, Barnes \\& Kaib 2006).  The planet 55 Cnc $f$ was discovered by Fischer \\etal at the inner edge of that stable zone.   The ``Packed Planetary Systems'' (PPS) hypothesis asserts that if a zone exists in which massive planets are dynamically stable, then that zone is likely to contain a massive planet (BR04, RB05, Raymond \\etal 2006; Barnes, Godziewski \\& Raymond 2008).  Although the idea behind the PPS hypothesis is not new (see, for instance, Laskar 1996), the large number of planetary systems being discovered around other stars allows PPS to be tested directly. Indeed, the $\\sim$ 1.4 Saturn mass planet HD 74156 $d$ recently discovered by Bean \\etal (2008) was located in the stable zone mapped out in BR04 and RB05, and with the approximate mass predicted by RB05 (Barnes \\etal 2008).  In addition, most of the first-discovered planetary systems are now known to be packed (Barnes \\etal 2008), as well as $\\sim$85\\% of the known two-planet systems (Barnes \\& Greenberg 2007).  The fact that 55 Cancri $f$ lies within the stable zone identified in previous work (BR04; RB05) also supports PPS, especially since planets $e$ through $c$ are packed, i.e. no additional planets could exist between them.  Several other planet predictions have been made and remain to be confirmed or refuted (see Barnes \\etal 2008) -- the most concrete outstanding prediction is for the system HD 38529 (see RB05).  Mean motion resonances (MMRs) are of great interest because they constrain theories of planet formation.  Models of convergent migration in gaseous protoplanetary disks predict that planets should almost always be found in low-order MMRs and with low-amplitude resonant libration (Snellgrove \\etal 2001; Lee \\& Peale 2002).  This may even have been the case for the giant planets in our Solar System (Morbidelli \\etal 2007).  On the other hand, planet-planet scattering can produce pairs of resonant planets in $\\sim 5\\%$ of unstable systems, but with large-amplitude libration and often in higher-order MMRs (Raymond \\etal 2008).  Thus, understanding the frequency and character of MMRs in planetary systems is central to planet formation theory.  In the context of PPS, 55 Cancri is an important system as it contains many planets, but still appears to have a gap large enough to support more planets. Therefore, PPS makes a clear prediction that another planet must exist between known planets $f$ and $d$.  In this paper we add massive hypothetical planets to the system identified by Fischer \\etal (2008) to determine which physical and orbital properties could still permit a stable planetary system.  We focus our search on the ``new'' stable zone between planets $f$ and $d$.  We also show that certain dynamically stable configurations are unlikely to contain a planet because the large eccentricity oscillations induced in the known planets significantly reduce the probability of Fischer \\etal having detected the known planets on their identified orbits, to within the observational errors.  The orbital regions that perturb the known planets most strongly correlate with specific dynamical resonances, such that we can put meaningful constraints on the masses of planets in those resonances.  Finally, we also use test particle simulations to map out the region of stability for additional planets beyond planet $d$, in the distant reaches of the planetary system. ",
        "conclusions": "We have mapped out the region in 55 Cancri where an additional planet $g$ might exist.  There is a broad region of stability between known planets $f$ and $d$ that could contain a $\\sim$Saturn-mass planet (Fig.~\\ref{fig:ae}). Since observations rule out a very massive planet, our simulations suggest that the region could easily support two or possibly even three additional planets.  In addition, one or more outer planets could be present in the system beyond about 10 AU.  However, such distant planets would not be detectable for many years.  We examined eight mean motion resonances in detail (see Table 2).  For two of these, 3f:1g (i.e., the 1:3 MMR between planet $f$ and hypothetical planet $g$) and 4g:1d, there was no stable region that exhibited regular libration of resonant arguments.  Therefore, these resonances can not contain planets in the mass range that we explored.  Given the very low FTD values, the 2f:1g MMR is unlikely to contain a resonant planet more massive than $\\sim 20 \\mearth$. Two other MMRs, 3g:1d and 5g:2d, may contain a stable, high-FTD resonant planet but the location of the MMRs is constrained to a very small region of ($a_g,e_g$) space which is surrounded by a chaotic region.  Finally, three MMRs, 3f:2g, 2g:1d, and 3g:2d, have a stabilizing influence and may contain planets near or even across the collision line with planet $f$ or $d$.  Each of these MMRs contains broad regions of stable libration of resonant angles, although the locations of low-FTD libration can vary with $M_g$.  We can therefore only weakly constrain the presence of an additional planet in one of these resonances.  The region between planets $f$ and $d$ contains many MMRs which display a wide range of behavior. In addition to stable and unstable resonances, the behavior of resonant arguments is also diverse. In some regions we would expect all resonant angles to librate regularly, but in others only some librate.  In two instances, planet $g$ could be in the apsidal corotation resonance (Michtchenko \\& Beauge 2003; Ferraz-Mello \\etal 2003): for large $M_g$ in the 2g:1d MMR at the $g-d$ collision line (see Fig.~\\ref{fig:libd21}), or in 5g:2d MMR (Fig.~\\ref{fig:resd52}).  Moreover, we also see cases of ``asymmetric'' libration in which the equilibrium angle is neither 0$^\\circ$ or 180$^\\circ$ (see Fig.~\\ref{fig:evold31}). Even if there are no additional planets in the $f-d$ gap, there could be an asteroid belt in which this diverse and exotic dynamical behavior is on display.  55 Cancri is a critical test of the ``Packed Planetary Systems'' (PPS) hypothesis, which asserts that any large contiguous stable region should contain a planet (BR04; RB05; Raymond \\etal 2006; Barnes \\etal 2008).  To date, two planets have been discovered in the three stable zones mapped out by BR04 and RB05 (in HD 74156 and 55 Cnc).  Given the width of the stable zone between planets $f$ and $d$, PPS indicates that at least one, and possibly two or three, more planet(s) should exist in 55 Cancri.  We look forward to further observations of the system that may find such planets, or perhaps show evidence of their absence.  Our results may be used to guide observers searching for planet $g$ and beyond."
    },
    "0808/0808.0543_arXiv.txt": {
        "abstract": "The temperatures of electrons and ions in the post-shock accretion region of a magnetic cataclysmic variable (mCV) will be equal at sufficiently high mass flow rates or for sufficiently weak magnetic fields. At lower mass flow rates or in stronger magnetic fields, efficient cyclotron cooling will cool the electrons faster than the electrons can cool the ions and a two-temperature flow will result. Here we investigate the differences in polarized radiation expected from mCV post-shock accretion columns modeled with one- and two-temperature hydrodynamics. In an mCV model with one accretion region, a magnetic field \\raisebox{-0.2em}{$\\stackrel{>}{\\sim}$} 30 MG and a specific mass flow rate of $\\sim$0.5 g cm$^{-2}$ s$^{-1}$, along with a relatively generic geometric orientation of the system, we find that in the ultraviolet either a single linear polarization pulse per binary orbit or two pulses per binary orbit can be expected, depending on the accretion column hydrodynamic structure (one- or two-temperature) modeled. Under conditions where the physical flow is two-temperature, one pulse per orbit is predicted from a single accretion region where a one-temperature model predicts two pulses. The intensity light curves show similar pulse behavior but there is very little difference between the circular polarization predictions of one- and two-temperature models. Such discrepancies indicate that it is important to model some aspect of two-temperature flow in indirect imaging procedures, like Stokes imaging, especially at the edges of extended accretion regions, were the specific mass flow is low, and especially for ultraviolet data. ",
        "introduction": "Magnetic cataclysmic variables (mCVs), consisting of the classes known as polars and intermediate polars, are composed of a Roche-lobe-filling M type main sequence star in orbit about a magnetic white dwarf (see \\citet{warner95} for a review). Mass is lost through the inner Lagrangian point, $L_{1}$, and flows toward the magnetosphere of the white dwarf either predominately in a stream (polars) or after forming a truncated accretion disk circulating around the white dwarf (intermediate polars). In either case, the ionized gas follows magnetic field lines to the surface of the white dwarf after the gas reaches the magnetosphere where the magnetic pressure exceeds the gas ram pressure. Upon reaching the white dwarf surface the gas will be essentially at ``free fall'', with highly supersonic velocities. The abrupt stop of the radial inflow near the surface of the white dwarf leads to the formation of a shock, which heats the inflowing material  \\citep{fabian76,king79,lamb79,wu00}. The hot subsonic post-shock flow settles gradually onto the white dwarf, and cools via emitting bremsstrahlung X-rays and optical/infra-red cyclotron radiation.  The hydrodynamic structure of the post-shock settling flow is determined by radiative and particle energy processes, which are essentially characterized by the bremsstrahlung cooling time $t_{\\mbox{\\small br}}$, the cyclotron cooling time $t_{\\mbox{\\small cy}}$, the electron-ion energy-exchange time  $t_{\\mbox{\\small ei}}$, the electron-electron collisional time $t_{\\mbox{\\small ee}}$, and the ion-ion collisional time $t_{\\mbox{\\small ii}}$ \\citep{king79,lamb79,imamura96,saxton05}. For weakly magnetic systems (with $B \\sim 10^6~$G or weaker) with accretion luminosities $10^{31}-10^{33}$ erg cm$^{-2}$ s$^{-1}$, typical of mCVs, $t_{\\mbox{\\small cy}} > t_{\\mbox{\\small br}} > t_{\\mbox{\\small ei}} > t_{\\mbox{\\small ee}}$. As the strength of the magnetic field increases to $B \\: \\raisebox{-0.2em}{$\\stackrel{>}{\\sim}$} \\: 10^7$~G, cyclotron cooling may dominate bremsstrahlung cooling, $t_{\\mbox{\\small cy}} < t_{\\mbox{\\small br}}$. For sufficiently strong magnetic fields and low specific accretion rates, $t_{\\mbox{\\small cy}}$ is so short ($t_{\\mbox{\\small cy}} < t_{\\mbox{\\small ei}}$) that collisions between electrons and ions cannot maintain an equal temperature between the two types of particles. The accretion flow is therefore in a two-fluid regime which requires a two-temperature (2T) hydrodynamic description. A strong magnetic field can also result in a situation where cyclotron radiative loss is so rapid ($t_{\\mbox{\\small cy}} < t_{\\mbox{\\small ee}}$) that electron-electron collisions are not efficient enough to maintain a Maxwellian distribution. In the extreme situation where $t_{\\mbox{\\small cy}} < t_{\\mbox{\\small ii}}$, the accretion flow is no longer hydrodynamic.  Previous calculations of cyclotron radiation from the post-shock settling flow in mCVs have either assumed a uniform density and temperature (\\citet{chanmugam81}; \\citet{meggitt82};\\\\ \\citet{barrett84};\\\\ \\citet{wickramasinghe85};\\\\ \\citet{wu88,wu89}) or a one-temperature (1T) structure (\\citet{wu90,wu92}; \\citet{potter02}). However, detailed 1D calculations of the hydrodynamic structure of a post-shock accretion column that self-consistently include cyclotron and bremsstrahlung cooling clearly show that a 2T structure is to be expected in many physical situations relevant to mCVs (\\citet{imamura96}; \\citet{woelk96}; \\citet{saxton01}; \\citet{wu03}; \\citet{saxton05}).  Here we have computed and compared the cyclotron radiation from a cylindrical post-shock accretion column, with a uniform cylindrical radial structure, assuming both a 1T hydrodynamic structure and a 2T hydrodynamic structure. The resulting cyclotron spectra for a grid of three white dwarf masses (0.5, 0.7 and 1.0 $M_{\\odot}$), three magnetic field strengths (10, 30 and 50 MG) and two mass flow rates ($10^{16}$ and $10^{14}$ g s$^{-1}$) were computed for various viewing inclination angles. For each case, using the computed viewing-angle dependent cyclotron spectra, Johnson bandpass \\citep{johnson65,bessell90} filtered light curves over an orbital period were computed for a mCV with an orbital inclination of 45$^{\\circ}$ and with the given accretion column at a co-latitude of 30$^{\\circ}$.  This work is organized as follows. In \\S \\ref{methods} we outline the hydrodynamic formulation used to determine the density and temperature structure of the post-shock flows and the radiative transfer through the ionized accreting gas. In \\S \\ref{results} we present the results of the polarized radiative transfer calculations, and in \\S \\ref{discussion} we examine the differences between the spectro-polarization properties of the emission from 1T and 2T accretion flows and discuss their implications. Concluding remarks are made in \\S \\ref{conclusion}. ",
        "conclusions": "Assuming a one-temperature hydrodynamic post-shock accretion column as the source for polarized radiation in models of magnetic cataclysmic variables can lead to erroneous predictions for the radiation when the cyclotron cooling efficiency is greater than the electron-ion energy exchange efficiency. This effect shows up at the lower mass flow rate modeled here ($\\dot{m} = 0.5$ g cm$^{-2}$s$^{-1}$) at higher cyclotron harmonics for the fairly generic white dwarf masses of 0.7 and 1.0 $M_{\\odot}$ and magnetic fields of 30 MG or greater. So the interpretation of light curve data obtained at higher frequencies, in the ultraviolet, needs to take into account the effect of two-temperature flow on the production of polarized radiation. In particular, the number of linear polarization pulses observed in the ultraviolet can be misinterpreted if a one-temperature accretion column flow is assumed."
    },
    "0808/0808.3226_arXiv.txt": {
        "abstract": "We present a spectral line survey of the C-rich envelope \\object{CIT\\,6} in the $\\lambda$ 2\\,mm and 1.3\\,mm bands carried out with the Arizona Radio Observatory (ARO) 12\\,m telescope and the Heinrich Hertz Submillimeter Telescope (SMT).  The observations cover the frequency ranges of 131--160\\,GHz, 219--244\\,GHz, and 252--268\\,GHz with typical sensitivity limit of $T_R<10$\\,mK.  A total of 74 individual emission features are detected, of which 69 are identified to arise from 21 molecular species and isotopologues, with 5 faint lines remaining unidentified. Two new molecules (C$_4$H and CH$_3$CN) and seven new isotopologues (C$^{17}$O, $^{29}$SiC$_2$, $^{29}$SiO, $^{30}$SiO, $^{13}$CS, C$^{33}$S, and C$^{34}$S) are detected in this object for the first time. The column densities, excitation temperatures, and fractional abundances of the detected molecules are determined using rotation diagram analysis. Comparison of the spectra of \\object{CIT\\,6} to that of \\object{IRC+10216} suggests that the spectral properties of \\object{CIT\\,6} are  generally consistent with those of \\object{IRC+10216}. For most of the molecular species, the intensity ratios of the lines detected in the two objects are in good agreement with each other. Nevertheless, there is evidence suggesting enhanced emission from CN and HC$_3$N and depleted emission from HCN, SiS, and C$_4$H in \\object{CIT\\,6}. Based on their far-IR spectra, we find that \\object{CIT\\,6} probably has a lower dust-to-molecular gas ratio than  \\object{IRC+10216}. To investigate the chemical evolution of evolved stars, we compare the molecular abundances in the AGB envelopes \\object{CIT\\,6} and \\object{IRC+10216} and those in the bright proto-planetary nebula \\object{CRL\\,618}. The implication on the circumstellar chemistry is discussed. ",
        "introduction": "The late stages of stellar evolution from the asymptotic giant branch (AGB) to planetary nebulae (PN) are now recognized as an active period of chemical synthesis of molecules.  The detection and analysis of millimeter wave molecular emission lines are fundamental to the understanding of the physical conditions and chemical processes leading to chemical synthesis.  Due to the rapid evolution of the star, the changing physical conditions, including  dust, stellar winds, shock waves, UV emission and X-rays from the central star, etc. play different roles in circumstellar chemistry. This leads to corresponding different circumstellar chemical compositions in different evolutionary stages.  The envelopes around C-rich stars, with their enhanced carbon abundance, provide a perfect cradle for molecule formation. Hitherto, more than 60 molecular species have been detected in C-star envelopes \\citep{glassgold96,olo97,cernicharo00,ziu07}, most of which were discovered through their rotational lines at millimeter wavelengths.  Recent improvement in telescope design and receiver performance enable us to detect new molecular emission with a higher sensitivity, and thus with the possibility of shedding new light on circumstellar chemistry. The most frequently investigated C-star envelope is \\object{IRC+10216}, which is one of the richest molecular sources in the sky. Several molecular line surveys have been presented for this object \\citep[see][and the references therein]{cernicharo00,he08}, which was found to harbor extremely abundant carbon chain and metal-containing molecules. \\object{IRC+10216} has been frequently used  as a standard reference for the chemical compositions of late-type stars. This inevitably invites the issue whether \\object{IRC+10216} is a chemically unique late-type star. To settle this question, we require systematic surveys of molecular line emission from other C-star envelopes. In the present study, we report a spectral-line survey of the C-star envelope \\object{CIT\\,6} at millimeter wavelengths. This allows us to compare the similarity and difference in the chemical compositions between the two C-star envelopes.   \\object{CIT\\,6} (\\object{RW\\,LMi, GL\\,1403, IRC+30219, IRAS\\,10131+3049}) was first discovered during the Caltech 2-$\\mu$m sky survey and was among the 14 very red infrared-bright optical-faint sources found \\citep{ulrich66}. \\object{CIT\\,6} is characterized by its very low color temperature, implying that the star is surrounded by a very thick dust  envelope and has been identified as a long-period variable with a period of about 628 days  \\citep{alksnis95}. From the period-luminosity relation, \\citet{cohen96} estimated the distance of \\object{CIT\\,6} to be $400\\pm50$\\,pc, which is slightly more distant than  \\object{IRC+10216}, which has a distance of between $\\sim120$\\,pc \\citep{groen98} and $\\sim150$\\,pc \\citep{gue99}. \\object{CIT\\,6}  is believed to be more evolved and has a lower mass loss rate compared to \\object{IRC+10216} \\citep[see][e.g.]{fukasaku94}.  The large polarization found in the visible and infrared wavelengths implies that the distribution of circumstellar material around the star is asymmetric \\citep[][]{kruszewski68,dyck71}. Multi-wavelength imaging observations have been performed to study the structure of the nebula around \\object{CIT\\,6}. The optical images obtained by the Hubble Space Telescope (HST) and the near-IR images of \\object{CIT\\,6} obtained by the Keck-I telescope have revealed a bipolar dust envelope and an elongated component with time-variable asymmetry \\citep{monnier00}. Several scattering arcs were revealed by the HST-NICMOS imaging polarimetry \\citep{schmidt02}. These arcs are nearly concentric and extend to large stellar radii. Mid-IR images of \\object{CIT\\,6} were obtained by \\citet{lagadec05} using the ESO 3.6-m telescope. A cometary-like feature was revealed in their 9.7\\,$\\mu $m image.  There have been several observations of molecular lines in \\object{CIT\\,6} at millimeter wavelength. \\citet{henkel85} reported observations of a few molecular lines in \\object{CIT\\,6} and \\object{IRC+10216} between 18 and 150\\,GHz. They found that relative abundances of observed molecules in the two sources have no significant differences. Using the Nobeyama 45\\,m radio telescope, \\citet{fukasaku94} observed a few transitions in the frequency ranges between 39--47\\,GHz and 85--91GHz in a sample of evolved stars including \\object{CIT\\,6} and found that the abundance of HNC increases with the evolutionary stage of the stars. \\citet{bujarrabal94} presented observations of 10 molecular transitions in C-rich and O-rich circumstellar envelopes including \\object{CIT\\,6} with the IRAM 30\\,m radio telescope at 1.3\\,mm, 2\\,mm, and 3\\,mm  windows. A recent molecular line survey was presented by \\citet{woods03} using the SEST 15\\,m and Onsala 20\\,m telescopes. They found that \\object{CIT\\,6} stands out from the other C-rich envelopes due to its high CN/HCN ratio and low HNC/HCN ratio.  To date, the  molecular species positively detected in \\object{CIT\\,6} at millimeter wavelengths include CO, $^{13}$CO, CN, $^{13}$CN, CS, SiO, SiS, $^{29}$Si$^{32}$S, C$_{2}$H, SiC$_{2}$, HCN, H$^{13}$CN, HNC, C$_{3}$N, HC$_{3}$N, HC$^{13}$CCN, HCC$^{13}$CN, and HC$_{5}$N.  In this paper, we present the first systematical line survey of \\object{CIT\\,6} at the 2\\,mm and 1.3\\,mm windows, using the Arizona Radio Observatory (ARO) 12\\,m telescope and the Heinrich Hertz Submillimeter Telescope (SMT). The observations are described in Sec.~2. In Sect.~3 we present the identifications and abundance calculations of the detected molecular species. In Sect.~4 we discuss the implication of our findings on circumstellar chemistry. The conclusions are given in Sect.~5. ",
        "conclusions": "The presence of rich molecular species around evolved stars provides an opportunity to study the evolution of chemistry in circumstellar envelopes, which have been widely suggested as one of the main sources of organic compounds in space.  As part of our project of investigating circumstellar chemistry, this paper reports a spectral line survey of the carbon-rich envelope \\object{CIT\\,6}, covering the frequency range between 131--160, 219--244, and 252--268\\,GHz with a high sensitivity.  A total of 74 lines are reported in the survey.  We identify 69 lines belonging to 21 different molecular species and isotopologues, most of which are carbon-bearing species. The new species include two carbon-chain molecules, C$_4$H and CH$_3$CN, and seven C, O, S, and Si isotopologues. Several new transitions from known species have been detected for the first time in this object. The species with the largest number of detected emission lines in our survey is SiC$_2$, which has 19 lines. It is followed by HC$_3$N, with 7 lines.  We find that the line profiles for some molecular species have different shapes, suggesting that the chemical structure is asymmetric in the envelope.  A comprehensive 3D photochemistry model is required to account for the line intensities and profiles in \\object{CIT\\,6}.  The excitation temperatures, column densities and abundances of the detected molecules are determined through rotation-diagram analysis. The spectra of \\object{CIT\\,6} are characterized by a large CN/HCN abundance ratio. Our results suggest that there is evidence for the photodissociation of HCN and SiS and the formation of CN and HC$_3$N in the evolved AGB envelope. The strong SiO and CS emission may suggest that depletion onto dust grains and destruction by shocks are insignificant in this object. An abundance comparison with the PPN \\object{CRL\\,618} implies to a rapid chemical evolution after a star leaves the AGB stage.  In order to investigate whether the molecular environment of  \\object{IRC+10216} is intrinsically unique, we systematically compare its spectra with those of \\object{CIT\\,6}. According to the comparison, we find that the molecular species can be classified into three groups, a) for most of the species, the intensity ratios of individual lines in the two objects are in good agreement with each other; b) the emission from HC$_3$N and CN may be enhanced in \\object{CIT\\,6}; c) the emission from SiS, HCN, and C$_4$H may be depleted in \\object{CIT\\,6}. The differences of the line-intensity ratios in the two objects are probably a consequence of chemical evolution with the exception  of C$_4$H, for which a high abundance in \\object{IRC+10216} cannot be explained by photochemical models. The ISO LWS spectra show that \\object{CIT\\,6} has a lower continuum-to-line ratio than  \\object{IRC+10216}, suggesting that the latter might have a larger dust-to-molecular gas ratio.  Using the same telescope settings, we also obtained the spectra of the AGB stars \\object{IRC+10216} and \\object{CRL\\,3068}, the PPN \\object{CRL\\,2688}, and the young PN \\object{NGC\\,7207}.  A detailed study of the chemical compositions in different evolutionary stages will be published in a separate paper."
    },
    "0808/0808.1824_arXiv.txt": {
        "abstract": "The Einstein-Aether theory provides a simple, dynamical mechanism for breaking Lorentz invariance. It does so within a generally covariant context and may emerge from quantum effects in more fundamental theories. The theory leads to a preferred frame and can have distinct experimental signatures. In this letter, we perform a comprehensive study of the cosmological effects of the Einstein-Aether theory and use observational data to constrain it.  Allied to previously determined consistency and experimental constraints, we find that an Einstein-Aether universe can fit experimental data over a wide range of its parameter space, but requires a specific rescaling of the other cosmological densities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.3683_arXiv.txt": {
        "abstract": "We present a new method to estimate the average star formation rate per unit stellar mass (SSFR) of a stacked population of galaxies. We combine the spectra of $600-1000$ galaxies with similar stellar masses and parameterise the star formation history of this stacked population using a set of exponentially declining functions. The strength of the Hydrogen Balmer absorption line series in the rest-frame wavelength range $3750-4150$\\AA\\ is used to constrain the SSFR by comparing with a library of models generated using the BC03 stellar population code. Our method, based on a principal component analysis (PCA), can be applied in a consistent way to spectra drawn from local galaxy surveys and from surveys at $z \\sim 1$, and is only weakly influenced by attenuation due to dust.  We apply our method to galaxy samples drawn from SDSS and DEEP2 to study mass-dependent growth of galaxies from $z \\sim 1$ to $z \\sim 0$.  We find that, (1) high mass galaxies have lower SSFRs than low mass galaxies; (2) the average SSFR has decreased from $z=1$ to $z=0$ by a factor of $\\sim 3-4$, independent of galaxy mass.  Additionally, at $z \\sim 1$ our average SSFRs are a factor of $2-2.5$ lower than those derived from multi-wavelength photometry using similar datasets.  We then compute the average time (in units of the Hubble time, $t_{\\rm H}(z)$) needed by galaxies of a given mass to form their stars at their current rate. At both $z=0$ and at $z=1$, this timescale decreases strongly with stellar mass from values close to unity for galaxies with masses $\\sim 10^{10} M_{\\odot}$, to more than ten for galaxies more massive than $ 10^{11} M_{\\odot}$. Our results are in good agreement with models in which AGN feedback is more efficient at preventing gas from cooling and forming stars in high mass galaxies. ",
        "introduction": "\\label{sec:intro} Over the last decade, there have been a large number of  photometric and spectroscopic surveys designed to study the formation and evolution of galaxies. One major conclusion of these studies has been that the epoch when massive galaxies formed most of their stellar mass is significantly earlier than that for low mass galaxies \\citep{Heavens04, Thomas05}. This phenomenon, popularly known as ``down-sizing'', at first sight seems  at odds with the predictions of the hierarchical CDM model, in which dark matter halos of all masses grow through merging and accretion right up to the present day. The only way to reconcile the observations with the theory, is to postulate that the growth of the most massive galaxies is much slower than the growth of their surrounding halos.  The most natural way to achieve this is by invoking feedback processes, which prevent gas from cooling, condensing and forming stars in massive halos \\citep{Bower06, Croton06, DeLucia06, Guo08}.  A variety of different feedback mechanisms have been included in numerical and semi-analytic models, for example feedback from quasars and radio AGN \\citep{Silk98, Hopkins06, wang06, Wang07}, supernova heating \\citep{Cole00, Benson03, Stringer08} and heating by infalling substructure \\citep{Dekel08}.  A physical understanding of feedback is still lacking, however, and it is not understood which, if any, of the proposed mechanisms are most important in regulating the growth of galaxies.  It is likely that each mechanism will come into play at a different mass scale and cosmological epoch.  By quantifying in detail how galaxies of different masses grow as a function of time, we hope to clarify how different galaxies form their stars and how this is influenced by feedback processes.  There have been many attempts to measure the star formation rates of galaxy populations from the present day out to redshifts greater than 5 \\citep[see for example][]{Brinchmann04, Bauer05, Feulner05, Noeske07a}. For nearby galaxies, the H$\\alpha$ line provides the most reliable estimate of SFR, since it directly measures the number of ionizing photons from massive stars. A reasonably reliable  correction for dust extinction can be made provided one also measures the Balmer decrement H$\\alpha$/H$\\beta$ accounting accurately for stellar absorption. Beyond redshifts of $\\sim 0.4$, the H$\\alpha$ redshifts out of the optical part of the spectrum and is no longer accessible. [O {\\sc ii}] equivalent widths, ultraviolet (UV)-optical spectral energy distribution (SED)   fitting and infrared photometry are commonly used to estimate the star formation rate. Each of these indicators is subject to different disadvantages. The [O {\\sc ii}] equivalent width is strongly affected by dust and by  metallicity. The UV luminosities of galaxies are  also strongly affected by dust, while the infrared only provides a direct measure of SFR if one has measurements across the thermal peak of the spectrum. This has  not been  the case for many of the Spitzer surveys aimed at quantifying star formation in  high-redshift galaxies. All the indicators discussed above  may also be contaminated by AGN emission if there is an actively accreting black hole in the galaxy.  In this paper, we develop an approach to  estimate the amount of recent star formation experienced by a {\\em population} of galaxies. Our method is   based on the Balmer absorption lines  located in the rest-frame wavelength range $3750-4150$\\AA\\ of the galaxy spectrum. The advantages of our method are the following:  (1) The Balmer absorption lines are weakly influenced by dust attenuation and AGN contamination compared with the indicators discussed above.  (2) The wavelength range spanned by the Balmer absorption lines is accessible out to redshifts greater than 1, even in optical spectra. This means that the method can be applied in a consistent manner to both low and high redshift ($z \\sim 1$) samples.  (3) By stacking a large number of galaxies, we  estimate SFRs for a {\\em complete sample} of galaxies in a given stellar mass range, including those galaxies that are forming stars weakly or not at all. These galaxies are often excluded when SFRs are measured for individual objects.  We apply our method to a large sample of galaxy spectra from the Sloan Digital Sky Survey (SDSS) and the DEEP2 redshift survey to study the evolution of galaxies from $z=1$ to $z=0$. This redshift interval accounts for roughly half the age of the universe. Our study thus addresses the final stage of galaxy build-up in the Universe.  This paper is arranged as follows. In \\S2, we introduce the SDSS and DEEP2 samples used in our studies. The method to estimate the amount of recent star formation in our galaxies is developed in \\S3. We apply the method to the SDSS and DEEP2 samples, and present our results in \\S4 and \\S5, respectively. A discussion of the results is given in \\S6. \\S7 contains the summary of the paper. We use the cosmological parameters $H_0=70~{\\rm km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\Lambda}=0.7$ throughout this paper. ",
        "conclusions": "In this section, we translate SSFR into a dimensionless star formation activity parameter (see Dav\\'e 2008), defined as \\begin{equation} \\alpha_{sf} \\equiv \\frac{1}{\\rm SSFR}\\frac{1}{t_{\\rm H}(z)-1 \\rm Gyr}. \\end{equation} Physically, this can roughly represent the fraction of the Hubble time (minus a Gyr) that a galaxy needs to have formed its stars at its current rate in order to produce its current stellar mass\\footnote{We note that $\\alpha_{sf}$ is not simply the fractional Hubble time the galaxy takes to form its stars, because stellar mass is returned to the ISM, a fact which is not accounted for in this simple model.}. A Gyr is subtracted in order to take account of the fact that dark matter halos massive enough to host galaxies with reasonably high star formation rates take about 1 Gyr to assemble in a $\\Lambda$CDM Universe \\citep{Dave08}. A value of $\\alpha_{sf}$ of 1 indicates that the galaxies could feasibly have formed all their stellar mass by forming stars continuously at the rate now observed. If $\\alpha_{sf}$ is greater than 1, their past average SFR must have been greater than their current SFR for the stellar mass of the galaxy to have formed within a Hubble time.  In Figure 10, we show our estimate of this characteristic timescale $\\alpha_{sf}$ as a function of stellar mass at both high (red asterisks) and low (black diamonds) redshifts.  We remind the reader that the low-$z$, highest mass bin is uncertain due to aperture correction complications in the SDSS survey. From this plot, we find that $\\alpha_{sf}$ is a strongly increasing function of $M_*$; the timescale increases from values close to a Hubble time for galaxies with $10^{10} M_{\\odot}$ to more than an order of magnitude larger than the Hubble time for galaxies more massive than $10^{11}M_{\\odot}$.  We find that, at fixed stellar mass, $\\log \\alpha_{sf}$ evolves only slightly with redshift, decreasing by around 0.2dex between $z\\sim0$ and $z\\sim1$.  \\begin{figure} \\bc \\hspace{-0.6cm} \\resizebox{8.5cm}{!}{\\includegraphics{fig/f10.ps}}\\\\% \\caption{Star formation activity parameter $\\alpha_{sf}$ as a function of stellar mass. As in previous plots, black diamonds and red asterisks represent SDSS and DEEP2 results, respectively. Blue squares are the results derived from the DEEP2 mock catalogue, they have been displaced in the x-axis by plus 0.05dex to make the comparison clear.} \\ec \\end{figure}  \\subsection{Comparison with galaxy formation models}  The blue squares in Figure 10 (displaced in the x-axis by plus 0.05dex to make the comparison clear) show the characteristic timescale of star formation in the mock DEEP2 universe created from the Millennium run SAM as described in Section 5.2. A key result of this paper is that the model is entirely consistent with observational results at $z\\sim1$, both in amplitude and slope.  More generally, in theoretical models of galaxy formation, at a fixed stellar mass, $\\alpha_{sf}$ is predicted to remain constant out to redshifts greater than $2$ \\citep[Figure 2 of][]{Dave08}. This latter result is robust to methodology, both semi-analytic and smooth-particle-hydrodynamic (SPH) simulations agree. Specifically, between a redshift of 1 and 0 the models predict changes in $\\log\\alpha_{sf}$ of less than 0.1dex. Our observational results indicate a slightly larger increase of 0.2dex. Further studies using larger surveys will be required to confirm this small amount of evolution.  The amplitude of $\\alpha_{sf}$ and its variation with stellar mass depend on simulation methodology. The near unity $\\alpha_{sf}$ at both $z\\sim0$ and $z\\sim1$ for galaxies with $M_* \\la 10^{10.5}M_{\\odot}$ suggests that galaxy mass growth in this mass and redshift range may be dominated by smooth and steady cold mode accretion, as implemented in all current models of galaxy formation. The strong increase of $\\alpha_{sf}$ with $M_*$ that we observe at both redshift 1 and in the local Universe (Figure 10), is an expression of the phenomenon of `downsizing': massive galaxies have apparently completed most of their star formation at higher redshifts than low mass systems. In many current models of galaxy formation, the explanation of this behaviour is ``AGN feedback''.  More massive galaxies are more likely to host massive black holes which have the capability of producing more energy.  Of equal importance, these massive galaxies are hosted by larger halos, where the AGN energy can be well coupled to the material that would otherwise cool and fuel star formation in the galaxies.  As a result, AGN feedback through heating of the interstellar and intergalactic gas is more efficient in massive galaxies.  In summary, the observed amplitude and evolution of  $\\alpha_{sf}$ as presented in this paper, provide firm constraints for all galaxy formation models.  \\subsection{Comparison with previous results}  As we have shown, our results appear to be in good general agreement with cosmological galaxy formation models. We now turn to a comparison with previous results from the literature. As we shall show, there is a significant discrepancy of $0.3-0.4$dex between our $H_{\\rm Balmer}$ derived SSFRs, and those derived primarily from multiwavelength broad band photometry.  \\begin{figure} \\bc \\hspace{-0.6cm} \\resizebox{8.5cm}{!}{\\includegraphics{fig/f11.ps}}\\\\% \\caption{SSFR as a function of stellar mass. Black diamonds: aperture-corrected SDSS data. Red asterisks: SSFRs of DEEP2 galaxies. Blue crosses and green squares: SSFRs from Z07 \\& AEGIS respectively.} \\ec \\end{figure}  The results from two other studies of the SSFRs of galaxies at $z \\sim 1$ are compared with our results in Figure 11.  The blue crosses represent SSFRs derived from the COMBO-17 sample with Spitzer 24$\\mu$m and GALEX data (Zheng et al. 2007, hereafter Z07).  These authors use the measured UV+IR luminosities to derive SFRs \\citep{Bell05}. They take account of the contribution from galaxies that are not individually detected at 24$\\mu$m by stacking their images.  The green squares show results calculated by us for a sample of galaxies selected from the All-Wavelength Extended Groth Strip International Survey (AEGIS). The galaxies are a subsample of those used to create the composite spectra in this paper.  This sample is not exactly the same as presented in Noeske et~al. (2007, hereafter N07) and \\citet{Noeske07b}, as only star-forming galaxies were included in their analysis. Instead, we have averaged the SSFRs over all the galaxies in each stellar mass bin (i.e. both star-forming and quiescent ones), using the same weighting factors that we used in our spectral stacking analysis. Aside from this difference in sample, the method used to derive SFRs for the AEGIS galaxies is the same as used by N07.  The method used by N07 to derive SFRs is slightly different to that employed by Z07, in that information from emission lines in the spectra of the galaxies is utilized if it is available. For galaxies with $f_{24\\mu m}>60\\mu$Jy and strong emission lines, the total SFR is derived from a combination of the IR measurements and from DEEP2 emission lines (H$\\alpha$, H$\\beta$, or [O {\\sc ii}]$\\lambda$3727, depending on $z$) with no extinction correction. SFRs are derived from extinction-corrected emission lines only  for blue galaxies with strong emission lines and no detectable 24 micron emission. Red galaxies with weak emission lines, but no 24 micron detections, are considered star-forming and SFRs are derived from emission lines, assuming the same extinction corrections as for normal star-forming galaxies.    The dashed lines in Figure 11 are a linear fit to the data points for mass bins with $\\log M_{*}/M_\\odot \\ge 10$.  As can be seen from this plot, the slope of the relation between SSFR and mass that we derive is consistent with the results of AEGIS and Z07.  However our method yields a normalisation that is $0.3-0.4$dex lower than AEGIS and Z07.  This discrepancy is puzzling, but it does have a number of possible explanations:  1) {\\em Systematic differences in the calibration of different SFR indicators.} Our method is based on stellar absorption line indicators, while the N07 and Z07 results are based on a combination of UV, IR and emission lines. We have demonstrated that our results do agree with the SSFRs derived by B04 from emission lines at low redshift, which provides confidence that there is no significant discrepancy between our method and the standard calibration of extinction-corrected H$\\alpha$ to derive SFR.  We are unable to carry out the same test at $z \\sim 1$, because H$\\alpha$ is redshifted out of the spectral range covered by most galaxy surveys. Systematic differences in estimated stellar mass do occur when different population synthesis models are used due to differing mass-to-light ratios in the models. However, the same BC03 models have been used in the comparison to the AEGIS results.  2) {\\em Obscured AGN.} There has been no attempt to remove obscured AGN from the two $z=1$ galaxy samples with which we compare. Both the N07 and Z07 analyses make use of the 24 micron Spitzer passband as a star formation indicator.  At $z=1$, this corresponds to a rest-frame wavelength of 12 micron, which is very close in wavelength to where emission from a dusty torus would become very significant \\citep[see e.g.][]{Daddi07}. In addition, AGN emission could well be contaminating some of the optical emission lines used to estimate SFR. By contrast, the $H_{\\rm balmer}$ index originates from stellar atmospheres and is not expected to be contaminated from emission from an obscured AGN.  On the other hand, the vast majority of DEEP2 24 micron sources and  line--emitting galaxies have line ratios indicating star formation and not AGN \\citep{Weiner07}. If the AGN were highly obscured, they could contribute at 24 micron and not show up in optical lines, but to make up a difference of 0.3dex one would have to assign half of the 24$\\mu$m emission at z=1 to obscured AGN, which would appear quite extreme (B.~Weiner, private communication).  3) {\\em Evolution of IMF with redshift.} The SFR indicators used by N07 and Z07 trace O and B stars, whereas our SFR indicator is the Balmer series and is mainly influenced by A stars.  If the IMF changes with redshift such that more massive stars form at higher $z$ \\citep[e.g.][]{Dave08, Dokkum08}, then a discrepancy between the two methods may not be apparent in the analysis of the low redshift samples, but may become more pronounced at higher $z$.  4) {\\em Aperture effects.}  As we have discussed, the SSFRs we estimate for the $z=1$ galaxies may be biased somewhat low because the long-slit spectra preferentially sample the inner bulge of the galaxy. However, 95\\% of the DEEP2 galaxies (at all redshifts) with a line measurement have $r_{eff} < 0.95^{\\pp}$ as measured in the CFHT imaging \\citep{Weiner07}.  Thus the 1$^{\\pp}$ slit covers a large fraction of the galaxy. Additionally the seeing mixes the light into the slit to a much greater degree in DEEP2 than it does for the SDSS fibers. Thus the star formation gradients of $z=1$ galaxies would have to be extremely strong for this to make a factor $2-3$ difference to the SSFR estimated for the galaxy population as a whole. As can be seen from Figure 8, the correction for aperture effects in the low $z$ sample is only a factor of 2 on average, even in the case where the fibre only samples 25\\% of the total light.  In this paper, we are not able make a definitive conclusion with regard to the possibilities listed above. It is clear that there are many inherent uncertainties in estimating SFRs in galaxies and that more work is needed before the factors of $2-3$ offsets that we see between the different methods and the models can be understood in detail."
    },
    "0808/0808.2685_arXiv.txt": {
        "abstract": "In this paper un-binned statistical tools for analyzing the cosmic ray energy spectrum are developed and illustrated with a simulated data set. The methods are designed to extract accurate and precise model parameter estimators in the presence of statistical and systematic energy errors. Two robust methods are used to test for the presence of flux suppression at the highest energies: the Tail-Power statistic and a likelihood ratio test. Both tests give evidence of flux suppression in the simulated data. The tools presented can be generalized for use on any astrophysical data set where the power-law assumption is relevant and can be used to aid observational design. ",
        "introduction": "\\label{sec:Intro} The observation of suppression in the flux of the highest energy cosmic rays (CRs) has been of central interest to astro-particle physics since the prediction of the GZK-effect\\cite{refG,refZK} in 1966. Most recently both the Auger\\cite{Yamamoto:2007xj} and the HiRes\\cite{Abbasi:2007sv} detectors have released results favoring the observation of flux suppression at a $6\\sigma$ and $5\\sigma$ level of confidence, respectively.  With this in mind, we describe a set of statistical tools designed to extract the most accurate and precise information concerning the flux of the highest energy cosmic rays. By binning the data we can only lose information\\cite{refGolds} (see \\secref{sec:BvUB}) and therefore our statistical tools use an un-binned maximum likelihood approach\\cite{refPDG, refHowell, refNewm, refClauset07} to answer two related statistical questions: {\\it Is there flux suppression at the highest energies?} and, if yes, {\\it What are the characteristic cut-off energy and shape parameters?}  In detail we first generate a toy data set using the CRPropa package\\cite{refCRPropa}, as in \\secref{sec:ats}. We then fit this simulated data to the three models described in \\secref{sec:Models}. The un-binned maximum likelihood fit is outlined in \\secref{sec:ID} and methods for incorporating systematic and statistical energy errors are described in \\secref{sec:SysEE} and \\secref{sec:StatEE} respectively. In \\secref{sec:EtF} we describe several statistical tools for hypothesis testing: the Kolmogorov-Smirnov test, the tail power statistic\\cite{refPisa1, refHague, Yamamoto:2007xj}, and a likelihood ratio test\\cite{refHagueicrc1217}.  Though we cast our discussion in terms of cosmic ray energies, it is worth noting that these tools can be applied to any astrophysical data set where deviations from the power-law hypothesis are relevant, e.g. the galaxy correlation function\\cite{refZehavi} or gamma ray astronomy\\cite{refSchroedter}. ",
        "conclusions": "\\label{sec:SaC} In this paper we describe a set of statistical tools designed to extract the most accurate and precise information about the flux of the highest energy cosmic rays. We show how to use the un-binned likelihood method described in \\secref{sec:ID} to fit a data set to the three model distributions described in \\secref{sec:Models}. Techniques for incorporating the systematic and statistical errors associated with a real CR detector into the likelihood method are described in \\secref{sec:SysEE} and \\secref{sec:StatEE} respectively. In \\secref{sec:EtF} we describe $p$-values useful for extracting information about flux suppression. We show in \\secref{sec:TP} and \\secref{sec:MD} how an experimenter might use an {\\it a priori} estimate of the cut-off energy to maximize an observational setup for detecting flux suppression.  The collection of these statistical tools are the primary result of this paper. To answer the questions posed in the introduction for a given data set we suggest the following steps: \\begin{enumerate} \\item Estimate the best fit parameters $\\hat{\\theta}$ of the model; \\label{step:BF} \\begin{enumerate} \\item The estimates $\\tgH$, $\\tebH$ or $\\tehH$ and $\\tdH$ or $\\twH$ are determined via the likelihood \\equref{equ:lnQ},\\label{step:BFa} \\item The estimate of the minimum energy $\\teminH$ is that which minimizes the Kolmogorov distance $\\DKS$ (see \\secref{sec:GoF}).\\label{step:BFb} \\end{enumerate} \\item Shift the energies up and down according to the systematic uncertainty described in \\secref{sec:SysEE} and repeat step (\\ref{step:BF}). The resulting shift in parameter estimates gives the systematic uncertainty of those estimates. \\label{step:SysE} \\item Obtain the model parameter estimates using the methods in \\secref{sec:StatEE} to incorporate the statistical error of each event energy. \\label{step:StatE} \\item Test the model hypothesis; \\label{step:GoF} \\begin{enumerate} \\item The absolute goodness of fit for any of the models can be evaluated using $\\pKS$ in \\secref{sec:GoF},\\label{step:GoFa} \\item The Tail-Power statistic $\\pTP$ can be used to reject the single power-law hypothesis (nearly independently of the spectral index estimate, see \\secref{sec:TP})\\label{step:GoFb} \\item The single power-law may be rejected in favor of a specific alternative model using $\\pR$, here we study the double and Fermi power-law distributions (see \\secref{sec:MD}).\\label{step:GoFc} \\end{enumerate} \\end{enumerate} The best estimates for the {\\it characteristic cut-off energy and shape parameters}, determined via steps (\\ref{step:BF}), (\\ref{step:SysE}) and (\\ref{step:StatE}), are $\\tebH$ or $\\tehH$ and $\\tdH$ or $\\twH$ respectively. The presence of {\\it flux suppression at the highest energies} can be evaluated using step (\\ref{step:GoF}).  By applying these methods to the toy Monte-Carlo set of CRPropa events we illustrate in \\secref{sec:SumToy} how the procedure may be implemented on an actual CR detector, i.e. a detector with systematic and statistical event energies. Suppression in the tail is clear in \\figref{fig:cdf} and \\figref{fig:rescdf}; the tail power statistic is $4.6\\sigma$ and the $p$-value for the double (Fermi) power-law is $\\lg p_{\\text{DP}} = -2.7$ ($\\lg p_{\\text{FP}} = -1.9$).  The methods are sufficient and robust. Indeed, many of them have been applied by the Auger collaboration which reports suppression with $6\\sigma$ confidence\\cite{Yamamoto:2007xj}. These tools serve as a basis for further investigation of the CR spectrum such as evidence for more detailed spectral information. They can be applied to any data set, astrophysical or otherwise, to provide information both about data already collected and help to optimize future observations for detecting tail suppression.  \\newpage \\appendix"
    },
    "0808/0808.0363_arXiv.txt": {
        "abstract": "We report on the results of systematic infrared 2.5--5 $\\mu$m spectroscopy of 45 nearby ultraluminous infrared galaxies (ULIRGs) at $z <$ 0.3 using the IRC infrared spectrograph onboard the AKARI satellite. This paper investigates whether the luminosities of these ULIRGs are dominated by starburst activity, or alternatively, whether optically elusive buried active galactic nuclei (AGNs) are energetically important.  Our criteria include the strengths of the 3.3 $\\mu$m polycyclic aromatic hydrocarbon (PAH) emission features and the optical depths of absorption features at 3.1 $\\mu$m due to ice-covered dust grains and at 3.4 $\\mu$m from bare carbonaceous dust grains. Because of the AKARI IRC's spectroscopic capability in the full 2.5--5 $\\mu$m wavelength range, unaffected by Earth's atmosphere, we can apply this energy diagnostic method to ULIRGs at $z >$ 0.15. We estimate the intrinsic luminosities of extended (several kpc), modestly obscured (A$_{\\rm V}$ $<$ 15 mag) starburst activity based on the observed 3.3 $\\mu$m PAH emission luminosities measured in AKARI IRC slitless spectra, and confirm that such starbursts are energetically unimportant in nearby ULIRGs. In roughly half of the observed ULIRGs classified optically as non-Seyferts, we find signatures of luminous energy sources that produce no PAH emission and/or are more centrally concentrated than the surrounding dust. We interpret these energy sources as buried AGNs. The fraction of ULIRGs with detectable buried AGN signatures is higher in ULIRGs classified optically as LINERs than HII-regions, and increases with increasing infrared luminosity. Our overall results support the scenario that luminous buried AGNs are important in many ULIRGs at $z <$ 0.3 classified optically as non-Seyferts, and that the optical undetectability of such buried AGNs occurs merely because of a large amount of nuclear dust, which can make the sightline of even the lowest dust column density opaque to the ionizing radiation of the AGNs. ",
        "introduction": "Ultraluminous infrared galaxies (ULIRGs) have large luminosities ($L> 10^{12}L_{\\odot}$) that are radiated mostly as infrared dust emission \\citep{san88a,sam96}. This means that very luminous energy sources ($L> 10^{12}L_{\\odot}$) are present hidden behind dust, so that most of the energetic photons from these energy sources are once absorbed by the surrounding dust and the heated dust then emits strongly in the infrared. The ULIRG population dominates the bright end of the luminosity function in the local universe ($z <$ 0.3) \\citep{soi87}. The contribution from ULIRGs to the total infrared radiation density increases rapidly with increasing redshift \\citep{lef05} and becomes important at $z >$ 1.3 \\citep{per05}. Distinguishing whether the dust-obscured energy sources of ULIRGs are dominated by starburst activity (nuclear fusion inside stars), or whether active galactic nuclei (AGNs; active mass accretion onto compact supermassive black holes with M $>$ 10$^7$M$_{\\odot}$) are also energetically important, is closely related to unveiling the history of dust-obscured star formation and supermassive black hole growth in the universe.  Identifying the dust-obscured energy sources in ULIRGs is difficult because a large amount of molecular gas and dust is concentrated in the nuclear regions of ULIRGs \\citep{dow98,ink06,ima07b}, and can easily {\\it bury} (= obscure in all directions) the putative AGNs because the emitting regions of AGNs are spatially very compact. Unlike AGNs obscured by torus-shaped dust, which are classified optically as Seyfert 2s, such {\\it buried} AGNs are elusive to conventional optical spectroscopy \\citep{mai03}, and thus very difficult to detect. However, quantitative determination of the energetic importance of buried AGNs is indispensable in understanding the true nature of the ULIRG population.  To study such optically elusive buried AGNs in ULIRGs, observing them at wavelengths of low dust extinction is important. Infrared 3--4 $\\mu$m (rest-frame) spectroscopy is one such wavelength. Dust extinction at 3--4 $\\mu$m is as low as that at 5--13 $\\mu$m \\citep{lut96}. More importantly, starburst and AGN emission are distinguishable from spectral shapes based on the following two arguments. (1) A normal starburst galaxy should always show large equivalent width polycyclic aromatic hydrocarbons (PAH) emission at rest-frame 3.3 $\\mu$m, while an AGN shows a PAH-free continuum originating in AGN-heated hot dust emission \\citep{moo86,imd00,idm06}. Hence, if the equivalent width of the 3.3 $\\mu$m PAH emission is substantially smaller than that of starburst-dominated galaxies, then a significant contribution from an AGN to the observed 3--4 $\\mu$m flux is the most natural explanation \\citep{imd00,idm01,im03,idm06,ima06}. (2) In a normal starburst, the stellar energy sources and dust are spatially well mixed (Figure 1a), while in a buried AGN, the energy source (= a compact mass accreting supermassive black hole) is more centrally concentrated than the surrounding dust (Figure 1b). In a normal starburst with mixed dust/source geometry, the optical depths of dust absorption features at 3.1 $\\mu$m by ice-covered dust and at 3.4 $\\mu$m by bare carbonaceous dust cannot exceed certain thresholds, but can be arbitrarily large in a buried AGN \\citep{im03,idm06}. Therefore, detection of strong dust absorption features whose optical depths are substantially larger than the upper limits set by mixed dust/source geometry suggests the buried-AGN-like centrally concentrated energy source geometry \\citep{im03,idm06}.  Infrared 3--4 $\\mu$m slit spectroscopy of a large sample of nearby ULIRGs using infrared spectrographs attached to ground-based large telescopes was performed, and signatures of intrinsically luminous buried AGNs were found in a significant fraction of optically non-Seyfert ULIRGs \\citep{imd00,idm01,im03,idm06,ris06,ima06,ima07b,san08}. However, the observed sample was restricted to $z <$ 0.15 because above this redshift, a part of the above-mentioned spectral features pass beyond the Earth's atmospheric window ($L$-band; 2.8--4.1 $\\mu$m), making this 3--4 $\\mu$m energy diagnostic method impossible. Because of the 2.5--5 $\\mu$m spectroscopic capability of the IRC \\citep{ona07,ohy07}, which was mounted onboard the AKARI satellite \\citep{mur07} and thus unaffected by Earth's atmosphere, we can now extend this successful approach to more distant ULIRGs at $z> $ 0.15. As AKARI IRC spectroscopy from space is quite sensitive, given the lack of large background emission from Earth's atmosphere, this extension to higher redshift is feasible in terms of sensitivity. Additionally, the AKARI IRC employs slitless spectroscopy, so that the bulk of the extended starburst emission in host galaxies ($>$ several kpc scale) is covered, unless it is extended more than 1 $\\times$ 1 arcmin$^{2}$ (Onaka et al. 2007). The observed 3.3 $\\mu$m PAH emission luminosities in the AKARI IRC spectra can thus be used to roughly estimate the intrinsic luminosities of modestly obscured (A$_{\\rm V}$ $<$ 15 mag) extended starbursts in the ULIRG host galaxies, which is impossible with ground-based spectroscopy using a narrow ($<$ a few arcsec) slit. Based on observations of several sources, nearby ULIRGs are argued to be energetically dominated by highly obscured {\\it compact} ($<$1 kpc) nuclear cores, with small contributions from {\\it extended}, modestly obscured starbursts in host galaxies \\citep{soi00,fis00}. The AKARI IRC spectra can directly test whether this argument holds for the majority of nearby ULIRGs.  In this paper, we present the results of systematic 2.5--5 $\\mu$m slitless spectroscopy of a large number of nearby ULIRGs at $z <$ 0.3 using the AKARI IRC. Throughout this paper, $H_{0}$ $=$ 75 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}$ = 0.3, and $\\Omega_{\\rm \\Lambda}$ = 0.7 are adopted to be consistent with our previous papers \\citep{idm06,ima06}. The physical scale of 1$''$ is 0.74 kpc at $z =$ 0.040 (the nearest source), 2.44 kpc at $z =$ 0.15, and 3.84 kpc at $z =$ 0.268 (the farthest source). ",
        "conclusions": "We investigate energy sources of observed ULIRGs based on the 3.3 $\\mu$m PAH emission and 3.1 $\\mu$m and 3.4 $\\mu$m absorption features.  \\subsection{Modestly obscured starbursts} In most of the observed ULIRGs, the presence of starbursts is evident from detection of the 3.3 $\\mu$m PAH emission. Since dust extinction at 3.3 $\\mu$m is only $\\sim$1/15 of that in the optical $V$-band ($\\lambda$ = 0.6 $\\mu$m; Rieke \\& Lebofsky 1985; Lutz et al. 1996), the flux attenuation of 3.3 $\\mu$m PAH emission with dust extinction of A$_{\\rm V}$ $<$ 15 mag is less than 1 mag. Thus, the observed 3.3 $\\mu$m PAH emission luminosities can be used to quantitatively derive the intrinsic luminosities of modestly obscured (A$_{\\rm V}$ $<$ 15 mag) starburst activity. The observed 3.3 $\\mu$m PAH to infrared luminosity ratios (L$_{\\rm 3.3PAH}$/L$_{\\rm IR}$) are summarized in column 4 of Table 3.  Figures 4a and 4b respectively compare the observed 3.3 $\\mu$m PAH luminosity (L$_{\\rm 3.3PAH}$) and its rest-frame equivalent width (EW$_{\\rm 3.3PAH}$) measured in ground-based slit spectra (abscissa) and in AKARI IRC slitless spectra (ordinate). The abscissa and ordinate of Figure 4a trace nuclear ($<$kpc) and total starburst luminosities, respectively. If {\\it extended} (several kpc) starbursts in the host galaxies of ULIRGs are energetically much more important than modestly obscured {\\it nuclear} starbursts, then both the L$_{\\rm 3.3PAH}$ and EW$_{\\rm 3.3PAH}$ values in the ordinate are expected to be substantially larger than the abscissa. However, the difference in both L$_{\\rm 3.3PAH}$ and EW$_{\\rm 3.3PAH}$ is relatively small, a factor of a few at most. We see no evidence that the spatially {\\it extended}, modestly obscured starbursts in host galaxies are substantially (more than an order of magnitude) more luminous than the modestly obscured {\\it nuclear} starbursts in ULIRGs. Even including extended starbursts in host galaxies, the L$_{\\rm 3.3 PAH}$/L$_{\\rm IR}$ ratios in ULIRGs are factors of 2.5 to more than 10 times smaller than those found in less infrared-luminous starbursts ($\\sim$10$^{-3}$; Mouri et al. 1990; Imanishi 2002). Taken at face value, the detected modestly obscured starbursts can account for only $<$10\\% to at most 40\\% of the infrared luminosities of the observed ULIRGs. Our AKARI IRC spectra reinforce the previous arguments, based on a small sample \\citep{soi00}, that nearby ULIRGs at $z <$ 0.3 are not energetically dominated by extended modestly obscured starburst activity in host galaxies.  \\subsection{Buried AGNs with weak starbursts} The deficit in observed 3.3 $\\mu$m PAH luminosity relative to the infrared luminosity requires energy sources in addition to the modestly obscured (A$_{\\rm V}$ $<$ 15 mag) PAH-emitting starbursts. The first possibility is very highly obscured (A$_{\\rm V}$ $>>$ 15 mag) PAH-emitting starbursts, in which the 3.3 $\\mu$m PAH flux is severely attenuated, while the flux of longer wavelength infrared emission (8--1000 $\\mu$m) may not be highly attenuated. The second possibility is that an AGN exists that can produce large infrared dust emission luminosities with no PAH emission (see $\\S$1), decreasing the observed L$_{\\rm 3.3 PAH}$/L$_{\\rm IR}$ ratios. These two scenarios are difficult to differentiate based on the absolute PAH luminosities but can be distinguished by the {\\it equivalent width} of emission or absorption features.  In a normal starburst galaxy, where HII regions, molecular gas, and photodissociation regions are spatially well mixed (Figure 1a), the equivalent width of the 3.3 $\\mu$m PAH-emission feature is insensitive to dust extinction \\citep{idm06}. If PAH-free AGN emission contributes significantly to the observed 2.5--5 $\\mu$m flux, then the EW$_{\\rm 3.3PAH}$ value should decrease. The EW$_{\\rm 3.3PAH}$ values in starbursts have an average value of EW$_{\\rm 3.3PAH}$ $\\sim$ 100 nm, with some scatter, but never become lower than 40 nm \\citep{moo86}. Thus, we adopt EW$_{\\rm 3.3PAH}$ $\\lesssim$ 40 nm as a strong signature of significant AGN contribution to an observed flux.  Among ULIRGs classified optically as non-Seyferts in the IRAS 1 Jy sample, three LINER ULIRGs, IRAS 23129+2548, 08572+3915, and 17044+6720, and one HII-region ULIRG, IRAS 22088$-$1831, have EW$_{\\rm 3.3PAH}$ $<$ 20 nm (Table 3, column 5), more than a factor of 5 less than typical starburst galaxies. These ULIRGs are strong buried AGN candidates. The following non-Seyfert ULIRGs are also taken to contain luminous buried AGNs because the EW$_{\\rm 3.3PAH}$ values are $<$40 nm: the three LINER ULIRGs (IRAS 04074$-$2801, 11180+1623, and 21477+0502), two HII-region ULIRGs (IRAS 14202+2615 and 17028+5817E), and one optically unclassified ULIRG (IRAS 08591+5248).  \\subsection{Buried AGNs with coexisting strong starbursts}  Based on the EW$_{\\rm 3.3 PAH}$ values, we can easily detect buried AGNs with very weak starbursts. Even if strong starburst activity is present, {\\it weakly obscured} AGNs are detectable because weakly attenuated PAH-free continua from the AGNs can dilute the 3.3 $\\mu$m PAH emission considerably.  However, detecting {\\it deeply buried} AGNs with coexisting surrounding strong starbursts is very difficult (Figure 1c). Even if the intrinsic luminosities of a buried AGN and surrounding less-obscured starbursts are similar, the AGN flux will be more highly attenuated by dust extinction than the starburst emission. When a buried AGN is obscured by {\\it ice-covered} dust grains, the AGN flux at $\\lambda_{\\rm rest}$ = 3.3 $\\mu$m is attenuated even more severely by the strong, broad 3.1 $\\mu$m absorption feature, making the EW$_{\\rm 3.3PAH}$ values in observed spectra apparently large.  To determine whether a deeply buried AGN is present in addition to strong starbursts, we use the optical depths of dust absorption features found in the 2.5--5 $\\mu$m spectra. As described in $\\S$1 and in \\citet{idm06} in more detail, these values can be used to distinguish whether the energy sources are spatially well mixed with dust (a normal starburst), or are more centrally concentrated than the dust (a buried AGN).  For a normal starburst with the mixed dust/source geometry in a ULIRG's core, $\\tau_{3.1}$ cannot exceed 0.3, while a buried AGN can produce $\\tau_{3.1} >$ 0.3 \\citep{im03,idm06}. Therefore, detection of $\\tau_{3.1} >$ 0.3 can be used to argue for the presence of a buried AGN with a centrally concentrated energy source geometry. Considering the uncertainty of $\\tau_{3.1}$ with $\\sim$0.1 ($\\S$4), we classify ULIRGs with $\\tau_{3.1}$ $>$ 0.4 as buried AGN candidates.  Aside from the above ULIRGs with low EW$_{\\rm 3.3PAH}$, the following non-Seyfert ULIRGs in the IRAS 1 Jy sample are newly classified as buried AGNs: ten LINER ULIRGs (IRAS 05020$-$2941, 09463+8141, 11028+3130, 14121$-$0126, 16333+4630, 00482$-$2721, 09539+0857, 10494+4424, 16468+5200, 17028+5817W), four HII-region ULIRGs (IRAS 01199$-$2307, 01355$-$1814, 17068+4027, 11387+4116), and one optically unclassified ULIRG (IRAS 01494$-$1845). This large $\\tau_{3.1}$ method is sensitive to deeply buried AGNs but obviously misses weakly obscured AGNs, which are more easily detected with the above low EW$_{\\rm 3.3PAH}$ method. Hence, the large $\\tau_{3.1}$ and low EW$_{\\rm 3.3PAH}$ methods are used in complementary fashion to detect AGN signatures.  We can also investigate the dust/source geometry from the $\\tau_{3.4}$ value. ULIRGs with $\\tau_{3.4} >$ 0.2 can be used to argue for the presence of buried AGNs \\citep{im03,idm06}. Among the optically non-Seyfert ULIRGs in the IRAS 1 Jy sample, only the LINER ULIRG IRAS 08572+3915 displays $\\tau_{3.4}$ $>$ 0.2. The LINER ULIRG of interest, UGC 5101, and the Seyfert 2 ULIRG, IRAS 19254$-$7245, also show $\\tau_{3.4}$ $>$ 0.2. However, all of these three ULIRGs have already been classified as having luminous AGNs based on their low EW$_{\\rm 3.3PAH}$ values.  Imanishi et al. (2006a, 2007a) commented that exceptionally centrally concentrated starbursts (Figure 1d) and normal starbursts with mixed dust/source geometry obscured by foreground dust in edge-on host galaxies (Figure 1e) can also produce large $\\tau_{3.1}$ and $\\tau_{3.4}$ values, but argued that it is very unlikely for the bulk of ULIRGs with $\\tau_{3.1} >$ 0.3 and/or $\\tau_{3.4} >$ 0.2 to correspond to these non-AGN cases.  For the remaining non-Seyfert ULIRGs with EW$_{\\rm 3.3PAH}$ $>$ 40 nm, $\\tau_{3.1}$ $\\lesssim$ 0.4, and $\\tau_{3.4}$ $\\lesssim$ 0.2, no obvious buried AGN signatures were observed in the 2.5--5 $\\mu$m spectra. Their spectra can be explained by either of the following scenarios: normal starbursts with mixed dust/source geometry are energetically dominant, or AGNs are present, but the AGN emission is so highly attenuated that its contribution to the observed 2.5--5 $\\mu$m flux is not significant.  We have no way of distinguishing between these two scenarios. However, some examples are known (e.g., NGC 4418) in which buried AGN signatures are found only at wavelengths longer than 5 $\\mu$m \\citep{ima04,dud97,spo01}. Thus, the detected buried AGN fraction in the 2.5--5 $\\mu$m AKARI IRC spectra is only a lower limit.  \\subsection{Dust extinction and intrinsic AGN luminosities} In a buried AGN with centrally concentrated energy source geometry, dust at a temperature of 1000 K, which is close to the innermost dust sublimation radius, produces continuum emission with a peak at $\\lambda \\sim$ 3 $\\mu$m, assuming approximately blackbody emission. Since a foreground screen dust distribution model is applicable to a buried AGN (Imanishi et al. 2006a, 2007a), the $\\tau_{3.1}$ (ice-covered dust) and $\\tau_{3.4}$ (bare dust) values reflect the dust column density toward the 3 $\\mu$m continuum-emitting region, which is almost equal to the column density toward the buried AGN itself (Imanishi et al. 2006a, 2007a).  The 3.1 $\\mu$m absorption feature is detectable if ice-covered dust grains are present in front of the 3--4 $\\mu$m continuum emitting energy source. Such ice-covered dust grains are usually found deep inside molecular gas, where ambient UV radiation is sufficiently shielded \\citep{whi88,tan90,smi93,mur00}. The 3.4 $\\mu$m absorption feature should be detected if the 3--4 $\\mu$m continuum emitting energy source is obscured by bare carbonaceous dust grains \\citep{pen94,ima96,raw03}, but is undetected if the absorbing dust is ice-covered \\citep{men01}. Since absorbing dust consists of both bare and ice-covered dust grains, ULIRGs with obscured energy sources should show both features, and the true dust column density is derivable from a proper combination of $\\tau_{3.1}$ and $\\tau_{3.4}$.   However, despite the detection of the 3.1 $\\mu$m ice absorption features in many ULIRGs, only two ULIRGs in the IRAS 1 Jy sample, IRAS 08572+3915 and 17044+6720, show clear 3.4 $\\mu$m carbonaceous dust absorption features. Even including the two sources of interest, UGC 5101 and IRAS 19254$-$7245, only four ULIRGs display clearly detectable 3.4 $\\mu$m absorption features. The difference in the detection rate largely comes from the intrinsically smaller oscillator strength of the 3.4 $\\mu$m carbonaceous dust absorption feature ($\\tau_{3.4}$/A$_{\\rm V}$ = 0.004--0.007; Pendleton et al. 1994) compared to the 3.1 $\\mu$m ice absorption feature ($\\tau_{3.1}$/A$_{\\rm V}$ = 0.06; Tanaka et al. 1990; Smith et al. 1993; Murakawa et al. 2000). Even if a modestly large amount of bare carbonaceous dust grains is present in front of the continuum-emitting energy source, the $\\tau_{3.4}$ value is small, making the detection of the 3.4 $\\mu$m dust absorption feature difficult. Additionally, the 3.3 $\\mu$m PAH emission feature is often accompanied by a sub-peak at 3.4 $\\mu$m \\citep{tok91,imd00}, and this sub-peak may dilute the 3.4 $\\mu$m dust absorption feature at the same wavelength. In fact, all the four ULIRGs with detectable 3.4 $\\mu$m absorption features are limited to relatively weak PAH emitters (EW$_{\\rm 3.3PAH}$ $<$ 35 nm). Finally, when the spectrally broad 3.1 $\\mu$m absorption feature is strong, the absorption feature extends to the longer wavelength side of the 3.3 $\\mu$m PAH emission feature, making it difficult to distinguish the origin of the apparent flux depression at $\\lambda_{\\rm rest}$ $\\sim$ 3.4 $\\mu$m. Among the four sources with detectable 3.4 $\\mu$m absorption features, IRAS 08572+3915, 17044+6720, and 19254$-$7245 indeed display only weak or undetectable 3.1 $\\mu$m absorption features. The remaining source, UGC 5101, shows large $\\tau_{3.1}$, but we can recognize the 3.4 $\\mu$m absorption feature, primarily because UGC 5101 is one of the brightest sources and the signal to noise ratios are among the highest. Detection of the 3.4 $\\mu$m absorption features in the remaining ULIRGs with large EW$_{\\rm 3.3PAH}$, large $\\tau_{3.1}$, and limited signal-to-noise ratios in the continuum is basically difficult. Therefore, while the total dust column density can be estimated in a reasonably reliable way for ULIRGs with both detectable 3.1 $\\mu$m and 3.4 $\\mu$m absorption features, the estimated dust column densities are only lower limits, and should be much smaller than the actual values for ULIRGs with only detectable 3.1 $\\mu$m absorption features. The estimated dust column densities are summarized in column 4 of Table 4.  In a buried AGN, the surrounding dust has a strong temperature gradient in that inner dust, close to the central energy source, has higher temperature than outer dust. Luminosity is transferred at each temperature, and the intrinsic luminosity of inner hot dust emission at 3--4 $\\mu$m ($\\nu$F$_\\nu$) should be comparable to that of outer cool dust emission at 60 $\\mu$m, the wavelength which dominates the observed infrared emission of ULIRGs \\citep{san88a}. Thus, if the intrinsic AGN's 3--4 $\\mu$m luminosity ($\\nu$F$_\\nu$) is comparable to the observed infrared luminosities of ULIRGs, then we can argue that the buried AGN is energetically important. For ULIRGs with low EW$_{\\rm 3.3PAH}$ ($<$40 nm), we can roughly extract the AGN's PAH-free continuum at 3--4 $\\mu$m, based on the assumption that starburst activity intrinsically shows EW$_{\\rm 3.3PAH}$ $\\sim$ 100 nm. That is, for ULIRGs with EW$_{\\rm 3.3PAH}$ = 30 nm (30\\% of the typical starburst value), we consider that 70\\% of the 3--4 $\\mu$m continuum comes from AGN's PAH-free continuum emission. Thus, we can estimate the {\\it observed } 3--4 $\\mu$m flux of AGN emission. Next, for ULIRGs with both detectable 3.1 $\\mu$m and 3.4 $\\mu$m absorption features (IRAS 08572+3915, UGC 5101, and IRAS 19254$-$7245), and for IRAS 17044+6720, which displays only the 3.4 $\\mu$m absorption feature, we can estimate dust column density, or {\\it dust extinction}, toward the 3--4 $\\mu$m continuum emitting regions based on $\\tau_{3.1}$ and $\\tau_{3.4}$. When we combine the {\\it observed} AGN flux at 3--4 $\\mu$m and {\\it dust extinction} toward the 3--4 $\\mu$m continuum emitting regions, we can derive the dust-extinction-corrected {\\it intrinsic} AGN flux, and thus the intrinsic AGN luminosity. If we adopt A$_{\\rm 3-4 \\mu m}$/A$_{\\rm V}$ $\\sim$ 0.058 \\citep{rie85}, then we find $\\tau_{3.1}$/A$_{\\rm 3-4 \\mu m}$ $\\sim$ 1 and $\\tau_{3.4}$/A$_{\\rm 3-4 \\mu m}$ $\\sim$ 0.069--0.12 for the Galactic interstellar medium. Assuming this relationship, we estimate the intrinsic AGN luminosities at 3--4 $\\mu$m ($\\nu$F$_{\\nu}$) to be L $\\gtrsim$ 10$^{12}$L$_{\\odot}$ in all the 3.4$\\mu$m-absorption-detected ULIRGs. The putative AGN activity is therefore energetically sufficient to quantitatively account for the bulk of the infrared luminosities of these ULIRGs (L$_{\\rm IR}$ $\\sim$ 10$^{12}$L$_{\\odot}$).  \\subsection{Dependence of the buried AGN fraction on optical spectral type: LINER vs. HII-region} In total, AKARI IRC 2.5--5 $\\mu$m spectra of 19 LINER, 16 HII-region, and 5 optically unclassified ULIRG's nuclei in the {\\it IRAS} 1 Jy sample were obtained. The low EW$_{\\rm 3.3PAH}$ method suggests six LINER and three HII-region ULIRG's nuclei contain luminous buried AGNs ($\\S$5.2). In addition to these ULIRGs, the large $\\tau_{3.1}$ method classifies ten LINER and four HII-regions ULIRG's nuclei as sources with deeply obscured buried AGNs ($\\S$5.3). In total, the detected buried AGN fraction is 16/19 (84\\%) for LINER ULIRGs, and 7/16 (44\\%) for HII-region ULIRGs (Table 4, columns 5 and 6). Since the selection of the observed ULIRGs is based solely on the target's visibility from the AKARI satellite and should be unbiased in terms of their energy sources ($\\S$2), we argue that the detected buried AGN fraction is higher in LINER ULIRGs than in HII-region ULIRGs. The same result was found from ground-based $L$-band spectroscopy and Spitzer 5--35 $\\mu$m spectroscopy of ULIRGs at $z <$ 0.15 (Imanishi et al. 2006a, 2007a), and also from a VLA radio observational search for compact radio core emission (another good AGN indicator) \\citep{nag03}. Therefore, we confirm that a larger fraction of LINER ULIRGs possess luminous buried AGNs than HII-region ULIRGs with a probability of $\\sim$99\\%.  The higher buried AGN fraction in optically LINER ULIRGs can be explained qualitatively by a dustier starburst scenario \\citep{ima07a}. For starburst/buried AGN composite ULIRGs (Figure 1c), the optical LINER or HII-region classifications are likely largely affected by the properties of the modestly obscured starbursts at the exteriors of the buried AGN, rather than by buried AGN-related emission, as optical observations can probe only the surfaces of dusty objects. In a dusty starburst, shock-related emission can be relatively important in the optical compared to the emission from the HII-regions themselves, resulting in optical LINER classification.  When a luminous AGN is placed at the center of a {\\it less dusty} starburst classified optically as an HII-region, the AGN emission is more easily detectable in the optical, making such an object an optical Seyfert. In contrast, when a luminous AGN is placed at the center of a {\\it dusty} starburst classified optically as a LINER, the AGN emission is more elusive in the optical, so that such an object is classified as an optical non-Seyfert. Hence, this scenario can explain the observed higher fraction of optically elusive {\\it buried} AGNs in optically LINER ULIRGs compared to HII-region ULIRGs. In fact, the emission probed in the optical was found to be dustier in LINER ULIRGs than in HII-region ULIRGs (Veilleux et al. 1995, 1999).  \\subsection{The buried AGN fraction as a function of infrared luminosity}  Due to the inclusion of ULIRGs at $z >$ 0.15, we have now a large number of ULIRGs with L$_{\\rm IR}$ $\\gtrsim$ 10$^{12.3}$L$_{\\odot}$. Specifically, only 1 of 13 observed non-Seyfert ULIRGs at $z <$ 0.15 has L$_{\\rm IR}$ $\\gtrsim$ 10$^{12.3}$L$_{\\odot}$, while 16 of 26 observed non-Seyfert ULIRGs at $z >$ 0.15 have L$_{\\rm IR}$ $\\gtrsim$ 10$^{12.3}$L$_{\\odot}$ (see Table 1). We can thus investigate the buried AGN fraction, separating ULIRGs into two categories: those with L$_{\\rm IR}$ $<$ 10$^{12.3}$L$_{\\odot}$ and those with L$_{\\rm IR}$ $\\gtrsim$ 10$^{12.3}$L$_{\\odot}$.  Among the observed non-Seyfert ULIRGs in the IRAS 1 Jy sample, 22 sources have L$_{\\rm IR}$ $<$ 10$^{12.3}$L$_{\\odot}$ and the remaining 17 sources have L$_{\\rm IR}$ $\\gtrsim$ 10$^{12.3}$L$_{\\odot}$. Based on the buried AGN signatures in Table 4, the fraction of ULIRGs with detectable buried AGN signatures is 12/22 (= 55\\%) for non-Seyfert ULIRGs with L$_{\\rm IR}$ $<$ 10$^{12.3}$L$_{\\odot}$ and 12/17 (= 71\\%) for non-Seyfert ULIRGs with L$_{\\rm IR}$ $\\gtrsim$ 10$^{12.3}$L$_{\\odot}$. Although the total sample size is not large, we find a higher buried AGN fraction with increasing ULIRG infrared luminosity. \\citet{idm06} investigated, based on ground-based $L$-band (2.8--4.1 $\\mu$m) spectra, the buried AGN fraction in a larger number of ULIRGs at $z <$ 0.15 in the IRAS 1 Jy sample than in this paper. Several ULIRGs at $z <$ 0.15 studied  in this paper are also included in this ground-based study. When we divide the observed ULIRGs by \\citet{idm06} into those with L$_{\\rm IR}$ $<$ 10$^{12.3}$L$_{\\odot}$ and $\\gtrsim$ 10$^{12.3}$L$_{\\odot}$, the detected buried AGN fraction is 16/29 (= 55\\%) in non-Seyfert ULIRGs with L$_{\\rm IR}$ $<$ 10$^{12.3}$L$_{\\odot}$, and 6/9 (= 67\\%) in non-Seyfert ULIRGs with L$_{\\rm IR}$ $\\gtrsim$ 10$^{12.3}$L$_{\\odot}$. The detected buried AGN fractions in ground-based and AKARI IRC spectra are comparable for non-Seyfert ULIRGs with both L$_{\\rm IR}$ $<$ 10$^{12.3}$L$_{\\odot}$ ($\\sim$55\\%) and $\\gtrsim$ 10$^{12.3}$L$_{\\odot}$ ($\\sim$70\\%).  We therefore argue that the detected buried AGN fraction increases with ULIRG infrared luminosity. This trend parallels the higher detection rate of optical Seyfert signatures in ULIRGs with higher infrared luminosities \\citep{vei99}.  The fraction of both optical Seyferts and luminous buried AGNs is significantly smaller in galaxies with L$_{\\rm IR}$ $<$ 10$^{12}$L$_{\\odot}$ than ULIRGs \\citep{vei99,soi01}. Hence, we conclude that AGN activity becomes more important as the infrared luminosities of galaxies increase. Recently, the so-called galaxy downsizing phenomena have been found, where galaxies with currently larger stellar masses have finished their major star-formation in earlier cosmic age \\citep{cow96,bun05}. AGN feedbacks are proposed to be responsible for the galaxy downsizing phenomena \\citep{gra04,bow06,cro06}. Namely, in galaxies with currently larger stellar masses, AGN feedbacks have been stronger in the past, and gas has been expelled in a shorter time scale. Buried AGNs can have particularly strong feedbacks, because the AGNs are surrounded by a large amount of nuclear gas and dust. If we reasonably assume that galaxies with currently larger stellar masses have previously been more infrared luminous, then the detected higher buried AGN fraction in more infrared luminous galaxies may support the AGN feedback scenario as the origin of the galaxy downsizing phenomena \\footnote{ To form more stars, more star-formation should have been occurred in the past, producing stronger star-formation related infrared emission in the past. If AGN's energetic contribution is negligible in LIRGs with L$_{\\rm IR}$ $<$ 10$^{12}$L$_{\\odot}$, but important, say $\\sim$50\\%, in ULIRGs with L$_{\\rm IR}$ $>$ 10$^{12}$L$_{\\odot}$, then AGN feedbacks can be stronger in ULIRGs, and yet higher infrared luminosity ($\\sim$5 $\\times$ 10$^{11}$L$_{\\odot}$) can come from star-forming activity in ULIRGs, producing more stellar masses in ULIRGs than in LIRGs. }.  \\subsection{Comparison with Seyfert 2 ULIRGs} We compare the 2.5--5 $\\mu$m spectral properties of non-Seyfert ULIRGs showing buried AGN signatures with those of Seyfert 2 ULIRGs (i.e., known obscured-AGN-possessing ULIRGs).  The apparent main difference between them is that the ionizing radiation from the putative buried AGNs in non-Seyfert ULIRGs is obscured by the surrounding dust along virtually all lines-of-sight; in contrast, dust around the AGNs in Seyfert 2 ULIRGs is distributed in a ``torus'', and ionizing radiation from the AGNs can escape along the torus axis, allowing for optical Seyfert signature detection (Figure 5).  The four ULIRGs classified optically as Seyfert 2s in the IRAS 1 Jy sample show no clear absorption features at 2.5--5 $\\mu$m. \\citet{idm06} also found in ground-based $L$-band spectra that the fraction of optically classified Seyfert 2 ULIRGs showing strong dust absorption features is substantially smaller than buried AGNs in optically non-Seyfert ULIRGs. We thus argue that the line-of-sight dust column density toward the AGNs is lower in Seyfert 2 ULIRGs than buried AGNs in non-Seyfert ULIRGs. Thus, buried AGNs and Seyfert 2 AGNs (obscured by torus-shaped dust) differ not only in dust geometry, but also in dust column density along our line-of-sight; specifically, the dust columns toward the buried AGNs are much higher \\citep{idm06}. Since the dust covering factor around buried AGNs (almost all directions) is also larger than Seyfert 2 AGNs (torus-shaped), the total amount of nuclear dust must be larger in the former (Figure 5).  Since the gas and dust in an AGN have angular momentum with respect to the central supermassive black hole, an axisymmetric spatial distribution is more natural than a spherical geometry. In this case, the column densities can be high in certain directions but low in others (Figure 5). For a fixed angular momentum, the dust column density ratios between the highest and lowest column directions are similar among different galaxies. If the total amount of nuclear dust is modest, the direction of the lowest dust column density can be transparent to the AGN's ionizing radiation, making Seyfert signatures detectable in the optical spectra. As the total amount of nuclear dust increases, even the direction of the lowest dust column density can be opaque to the AGN's ionizing radiation, making such galaxies optically elusive buried AGNs. Thus, all of the observed spectral properties of buried AGNs and Seyfert-type AGNs are explicable given a larger amount of nuclear dust in the former. Since ULIRGs contain a large amount of nuclear gas and dust \\citep{sam96}, the high buried AGN fraction is inevitable. Understanding optically elusive buried AGNs is therefore essential if we are to unveil the true nature of the ULIRG population.  \\subsection{Dependence on far-infrared colors} Based on the {\\it IRAS} 25 $\\mu$m to 60 $\\mu$m flux ratio ($f_{\\rm 25}$/$f_{\\rm 60}$), ULIRGs are divided into cool ($<$ 0.2) and warm ($>$ 0.2) sources \\citep{san88b}. Many of the non-Seyfert ULIRGs with detectable buried AGN signatures show cool far-infrared colors (Table 1). Although AGNs classified optically as Seyferts usually show warm far-infrared colors \\citep{deg87,kee05},  cool far-infrared colors of buried AGNs are the natural consequence of a large amount of nuclear dust, where contributions from the outer, cooler dust components to the infrared radiation become more important than for optical Seyferts.  Figure 6(a) compares {\\it IRAS} 25 $\\mu$m to 60 $\\mu$m flux ratios (i.e., far-infrared color) with the observed 3.3 $\\mu$m PAH to infrared luminosity ratios (L$_{\\rm 3.3PAH}$/L$_{\\rm IR}$). Seyfert ULIRGs tend to appear in the warmer far-infrared color range than non-Seyfert ULIRGs, as expected from the decreased amount of nuclear dust in the former ($\\S$5.7). No systematic difference in the L$_{\\rm 3.3PAH}$/L$_{\\rm IR}$ ratios between non-Seyfert and Seyfert ULIRGs exists.  Figure 6(b) compares the far-infrared colors and EW$_{\\rm 3.3PAH}$ for non-Seyfert ULIRGs.  Buried AGNs appear in both the warm and cool ranges. Although it is sometimes argued that ULIRGs with cool far-infrared colors must be starburst-dominated, simply because Seyfert-type AGNs show warm far-infrared colors (e.g., Downes \\& Solomon 1998), we do not confirm this to be true for the heavily buried AGNs."
    },
    "0808/0808.0539_arXiv.txt": {
        "abstract": "Using a large  galaxy group  catalogue constructed  from the Sloan  Digital  Sky  Survey Data  Release  4  (SDSS  DR4) with  an  adaptive halo-based  group finder,  we investigate  the luminosity  and  stellar mass functions for  different populations of galaxies  (central versus satellite; red versus blue; and galaxies in  groups of different masses) and for groups themselves. The  conditional stellar  mass function (CSMF),  which describes the stellar  distribution of galaxies in  halos of a given  mass for central and satellite galaxies  can be well modeled with  a log-normal distribution and a modified  Schechter form, respectively. On average,  there are about 3 times  as  many  central   galaxies  as  satellites.   Among  the  satellite population, there are  in general more red galaxies than  blue ones. For the central population, the luminosity function  is dominated by red galaxies at the massive  end, and by  blue galaxies  at the low  mass end.  At  the very low-mass  end ($M_\\ast  \\la 10^9  h^{-2}\\Msun$), however, there is a  marked increase in  the number of red  centrals.  We speculate  that these galaxies are located close  to large halos so that their  star formation is truncated by the large-scale environments.  The stellar-mass function of galaxy groups is well described by a double  power law, with a characteristic stellar mass at $\\sim 4\\times 10^{10}h^{-2}\\Msun$.   Finally, we use the observed stellar mass function of central galaxies to  constrain the stellar mass - halo mass relation for low  mass halos, and obtain $M_{\\ast,  c}\\propto M_h^{4.9}$ for $M_h \\ll 10^{11} \\msunh$. ",
        "introduction": "In recent years,  great progress has been made in  our understanding about how galaxies form and evolve in dark matter halos owing to the development of halo models  and   the  related  halo  occupation  models.    The  halo  occupation distribution (hereafter HOD), $P(N \\vert M_h)$, which gives the probability of finding $N$ galaxies (with some specified properties) in a halo of mass $M_h$, has  been extensively used  to study  the galaxy  distribution in  dark matter halos and galaxy clustering on large  scales (e.g.  Jing, Mo \\& B\\\"orner 1998; Peacock \\& Smith 2000; Seljak  2000; Scoccimarro \\etal 2001; Jing, B\\\"orner \\& Suto 2002;  Berlind \\&  Weinberg 2002; Bullock,  Wechsler \\&  Somerville 2002; Scranton 2002; Zehavi  \\etal 2004, 2005; Zheng \\etal  2005; Tinker \\etal 2005; Skibba et al. 2007; Brown  et al.  2008).  The conditional luminosity function (hereafter CLF), $\\Phi(L \\vert M_h) {\\rm d}L$, which refines the HOD statistic by  considering the average  number of  galaxies with  luminosity $L  \\pm {\\rm d}L/2$  that reside  in  a halo  of  mass $M_h$,  has  also been  investigated extensively (Yang, Mo \\&  van den Bosch 2003; van den Bosch,  Yang \\& Mo 2003; Vale \\& Ostriker 2004, 2008; Cooray 2006;  van den Bosch et al.  2007) and has been applied  to various redshift surveys,  such as the  2-degree Field Galaxy Redshift Survey (2dFGRS), the Sloan  Digital Sky Survey (SDSS) and DEEP2 (e.g. Yan, Madgwick \\& White 2003; Yang \\etal 2004; Mo et al. 2004; Wang \\etal 2004; Zehavi \\etal 2005; Yan, White \\& Coil 2004).  These investigations demonstrate that the HOD/CLF statistics are  very powerful tools to establish and describe the  connection between galaxies  and dark  matter halos,  providing important constraints  on  various physical  processes  that  govern  the formation  and evolution of  galaxies, such as  gravitational instability, gas  cooling, star formation, merging,  tidal stripping  and heating, and  a variety  of feedback processes, and how their efficiencies scale with halo mass.  Furthermore, they also  indicate that  the  galaxy/dark halo  connection  can provide  important constraints  on cosmology  (e.g.,van  den Bosch,  Mo  \\& Yang  2003; Zheng  \\& Weinberg 2007).  However, as pointed out in Yang et al.  (2005c; hereafter Y05c), a shortcoming of  the  HOD/CLF  models  is   that  the  results  are  not  completely  model independent.  Typically, assumptions have  to be made regarding the functional form of  either $P(N  \\vert M_h)$  or $\\Phi(L \\vert  M_h)$.  Moreover,  in all HOD/CLF studies to date, the  occupation distributions have been determined in an  indirect way:  the  free parameters  of  the assumed  functional form  are constrained  using {\\it  statistical}  data on  the  abundance and  clustering properties of the galaxy population.   An alternative method that can directly probe the galaxy - dark halo connection (e.g. HOD/CLF models) is to use galaxy groups as  a representation of dark matter  halos and to study  how the galaxy population changes with  the properties of the groups  (e.g., Y05c; Zandivarez et al. 2006; Robotham  et al. 2006; Hansen \\etal 2007; Yang  et al. 2008). For such a  purpose, one has to properly  find the galaxy groups  that are closely connected to  the dark matter halos.   In recent studies, Yang  et al. (2005a; 2007) developed  an adaptive  halo-based group finder  that has  such features \\footnote{In this paper, we refer  to systems of galaxies as groups regardless of their richness, including isolated  galaxies (i.e., systems with a single member) and rich clusters of galaxies.}.  This group finder has been applied to the  2dFGRS (Yang et al.   2005a) and to  the SDSS (Weinmann et  al. 2006a; Yang et al. 2007).  Detailed tests with mock galaxy catalogues have shown that this  group finder  is very  successful in  associating galaxies  according to their  common dark  matter halos.   In particular,  the group  finder performs reliably  not only  for rich  systems, but  also for  poor  systems, including isolated central galaxies in low mass  halos.  This makes it possible to study the  galaxy-halo connection  for systems  covering  a large  dynamic range  in masses.  With  a well-defined  galaxy group catalogue,  one can then  not only study the properties  of galaxies in different groups  (e.g.  Y05c; Yang \\etal 2005d; Collister \\& Lahav 2005; van den Bosch \\etal 2005; Robotham \\etal 2006; Zandivarez  \\etal 2006;  Weinmann \\etal  2006a,b;  van den  Bosch \\etal  2008; McIntosh \\etal 2007; Yang et al.   2008), but also probe how dark matter halos trace the large-scale structure of the Universe (e.g.  Yang \\etal 2005b, 2006; Coil \\etal 2006; Berlind \\etal 2007; Wang et al.  2008a).  Recently, this group finder has been applied to the Sloan Digital Sky Survey Data Release 4 (SDSS DR4), and the group catalogues constructed are described in detail in Yang \\etal (2007; Paper I hereafter).  In these catalogues various observational selection effects are taken into account, and each of the groups is assigned a reliable halo mass. The group catalogues including the membership of the groups are available at these links \\footnote{http://gax.shao.ac.cn/data/Group.html} \\footnote{http://www.astro.umass.edu/$^\\sim$xhyang/Group.html}.  In Yang et al.  (2008; Paper II hereafter) we have used these group catalogues to obtain various halo occupation statistics and to measure the CLFs for different populations of galaxies. In this paper, the third in the series, we will focus on the conditional stellar mass functions (CSMFs) for different populations of galaxies. In addition, we will also examine the general luminosity and stellar mass functions for different populations of galaxies and for groups themselves.  Finally, we will demonstrate how to use the observed luminosity and stellar mass functions for central galaxies to constrain the HOD in small halos.  This paper is organized as  follows: In Section~\\ref{sec_data} we describe the data (galaxy and group catalogues) used in this paper. Section~\\ref{sec_CSMFs} presents  our  measurement of  the  CSMFs  for all,  red  and  blue galaxies. Sections~\\ref{sec_LF_gax} and ~\\ref{sec_LF_grp} present our measurement of the luminosity and stellar  mass functions for galaxies and  groups, respectively. In Section  \\ref{sec_small}, we probe  the properties of the  central galaxies that can be formed in those small halos.  Finally, we summarize our results in Section~\\ref{sec_summary}.   Throughout  this  paper,  we use  a  $\\Lambda$CDM `concordance' cosmology  whose parameters  are consistent with  the three-year data    release    of    the     WMAP    mission:    $\\Omega_m    =    0.238$, $\\Omega_{\\Lambda}=0.762$,  $n_s=0.951$, $h=0.73$ and  $\\sigma_8=0.75$ (Spergel et al. 2007).  If not quoted, the units of luminosity, stellar and halo masses are in terms of  $h^{-2}\\Lsun$, $h^{-2}\\Msun$ and $h^{-1}\\Msun$, respectively. Finally, unless  noted differently, the luminosity functions  and stellar mass functions are  presented in units  of $h^3{\\rm Mpc}^{-3}  {\\rm d} \\log  L$ and $h^3{\\rm Mpc}^{-3} {\\rm  d} \\log M_{\\ast}$, respectively, where  $\\log$ is the 10 based logarithm. ",
        "conclusions": "\\label{sec_summary}  In this paper, we have derived the luminosity and stellar mass functions for different populations of galaxies (central versus satellite; red versus blue; and galaxies in halos of different masses), and for groups themselves, using a large galaxy group catalogue constructed from the SDSS Data Release 4 (DR4). Our main results can be summarized as follows: \\begin{enumerate} \\item  For central  galaxies, the  conditional stellar  mass  function (CSMF), which  describes the stellar  mass distribution  of galaxies  in halos  of a given mass can be well described  by a log-normal distribution, with a width $\\sigma_{\\rm log M_\\ast}\\sim 0.17$, quite independent of the host halo mass. The  median   central  stellar  mass  increases  rapidly   with  halo  mass, $M_\\ast\\propto M_h^{4.9}$, for halos with masses $M_h\\ll 10^{11}\\msunh$, but only   slowly,   $M_\\ast\\propto   M_h^{0.3}$,   for   halos   with   $M_h\\gg 10^{13}\\msunh$. \\item For satellite  galaxies, the conditional stellar mass  function in halos of different masses can be described reasonably well by a modified Schechter form than breaks away faster than the Schechter function at the massive end. The  faint end  slope appears  to  be steeper  for more  massive halos.   On average, there are about 3 times as many central galaxies as satellites. \\item  When stellar  mass functions  are measured  separately for  galaxies of different colors,  we find that the  central population is  dominated by red galaxies  at the massive  end, and  by blue  galaxies at  the low-mass  end. Among the satellite population, there  are in general more red galaxies than blue ones.  At the very  low-mass end ($M_\\ast \\la 10^9 h^{-2}\\Msun$), there is a marked increase in the  number of red centrals. We speculate that these galaxies are located  close to large halos so that  their star formation has been affected by their environments. \\item The stellar-mass  function of galaxy groups, which  describes the number density of  galaxy groups as a function  of the total stellar  mass of group member  galaxies,  is  well  described   by  a  double  power  law,  with  a characteristic stellar mass at  $\\sim 4\\times 10^{10}h^{-2}\\Msun$. This form is very different  from that of the halo mass  function, indicating that the efficiencies  of  star formation  in  halos  of  different masses  are  very different. \\item  The stellar  mass function  for  the central  galaxies can  be used  to provide stringent constraint  on the mean $M_{\\ast,c}$ -  $M_h$ relation for low-mass halos. \\end{enumerate}  We anticipate that a comparison of these results with predictions of numerical simulations and/or  semi-analytical models will  provide stringent constraints on how galaxies form and evolve in dark matter halos."
    },
    "0808/0808.0822_arXiv.txt": {
        "abstract": "{ Though pick-up ions (PUIs) are a well known phenomenon in the inner heliosphere, their phase-space distribution nevertheless is a theoretically unsettled problem. Especially the question of how pick-up ions form their suprathermal tails, extending to far above their injection energies, still now is unsatistactorily answered. Though Fermi-2 velocity diffusion theories have revealed that such tails are populated, they nevertheless show that resulting population densities are much less than seen in observations showing power-laws with a velocity index of ``-5''. We first investigate here, whether or not observationally suggested power-laws can be the result of a quasi-equilibrium state between suprathermal ions and magnetohydrodynamic turbulences in energy exchange with eachother. We demonstrate that such an equilibrium cannot be established, since it would require too high pick-up ion pressures enforcing a shock-free deceleration of the solar wind. We furthermore show that Fermi-2 type energy diffusion in the outer heliosphere is too inefficient to determine the shape of the distribution function there. As we can show, however, power-laws beyond the injection threshold can be established, if the injection takes place at higher energies of the order of 100 keV. As we demonstrate here, such an injection is connected with modulated anomalous cosmic ray (ACR) particles at the lower end of their spectrum when they again start being convected outwards with the solar wind. Therefore, we refer to these particles as ACR-PUIs. In our quantitative calculation of the pick-up ion spectrum resulting under such conditions we in fact find again power-laws, however with a velocity power index of ``-4'' and fairly distance-independent spectral intensities. As it seems these facts are observationally well supported by VOYAGER measurements in the lowest energy channels. ",
        "introduction": "Suprathermal ions, picked-up by the supersonic solar wind flow as ionized neutral atoms, have become known as pick-up ions (PUIs) and are produced all over the inner heliosphere with a typical upwind-downwind asymmetry with respect to the inflow direction of the neutral ISM inflow vector \\citep{rucinski93,fahr99}. In the case of PUI protons, their production is due to photoionization and charge exchange of interstellar H-atoms \\citep[see][]{rucinski91,fahr99,rucinski03,bzowski07}. Their spatial distribution seems well understood, while the PUI phase-space transport is a much less settled subject. Especially it exists an ongoing debate of how efficiently pick-up ions just after the pick-up process are accelerated to higher energies due to nonlinear wave-particle interactions \\citep[see e.g.][]{isenberg87,bogdan91,fichtner96,fichtner01,chalov96,chalov98,chalov04} and whether at all energy diffusion plays a relevant role in this transport.  Some hint is given by the solar wind proton temperature behavior with distance. The observed non-adiabatic temperature behavior namely proves that a specific solar wind proton heating must operate in the outer heliosphere which can only be due to energy absorption from pick-up ion generated turbulence, since convected turbulence amplitudes quickly die out with distance \\citep[see][]{smith01,chashei02}.  Freshly injected PUIs represent keV-energetic protons in the supersonic solar wind frame and may be called here: ``primary pick-up ions'' (or: PUIs$^{\\ast }$). The velocity distribution of these newly produced PUIs$^{\\ast }$ is toroidal and unstable \\citep[see][]{winske84,winske85,lee87,fahr88}. With the free energy of this unstable distribution PUIs$^{\\ast }$ drive Alfv\\'{e}nic wave power. The latter enforces pitchangle isotropization of the initial velocity distribution \\citep[see][]{chalov98,chalov99}. Due to wave-wave coupling, the wave energy generated by PUIs$^{\\ast }$ at the injection wavelength $\\lambda _{i}=U_{s}/\\Omega _{p} $ is diffusively transported in wavevector space both to smaller wavelengths where it can be absorbed by solar wind protons and to larger wavelengths where it is reabsorbed by all PUIs. This effect is seen as the main reason of solar wind proton heating occuring in the outer heliosphere \\citep{smith01,fahr02,chashei03,stawicki04}. Only a small fraction of about 5 percent of the PUI-generated wave energy reappears in the observed proton temperatures. Freshly injected PUIs excite turbulences that can organize a power-law distribution. From this distribution, both the  solar wind ions and the PUIs themselves can absorb energy as shown by \\citet{chashei03}. Also the approach by \\citet{isenberg03} where energy diffusion of pick-up ions is not taken into account shows that only a low degree (2-5 percent) of the pick-up ion driven wave energy is absorbed by solar wind protons in form of thermal energy.  This raises the question where the major portion of the wave energy produced during the primary pick-up process goes to. To clarify the energy redistributions, kinetic and spectral details of the relevant processes have to be investigated. A detailed numerical study of the PUI velocity distribution and the spectral Alfv\\'{e}nic/Magnetosonic wave power evolution has meanwhile been carried out \\citep{chalov04,chalov06a} and presents a simultaneous solution of a coupled system of equations consistently describing the isotropic velocity distribution function of PUIs and the spectral wave power intensity.  As one can see from this study, the largest portion of the self-generated wave energy is reabsorbed by PUIs themselves as a result of the cyclotron resonant interaction and leads to PUI-acceleration. It could perhaps be hoped that this energization of pick-up protons due to Fermi-2 stochastic acceleration processes eventually leads to the ubiquitous power-law PUI-tails pointed out by \\citet{fisk06,fisk07}. To the opposite, however, as reflected in the results presented by \\citet{chalov04,chalov06a} it is evident that this is not the case: even at larger distances close to the termination shock (100 AU) the PUI distributions show a rapid cut-off at energies higher than the injection energy. The question thus is raised here why power-laws have been seen at all. An explanation that we are favoring here is a new injection source to the PUI regime from high energies connected with modulated anomalous cosmic ray particles. These protons are primary ACR particles that occur with a spectrum down to the typical energy of the usually assumed PUIs. At this part of the spectrum, both particle species cannot be distinguished. Therefore, we refer to them as ACR-PUIs.  In Sect.~\\ref{sect2} we investigate the physical possibility of power-law ions in the outer heliosphere as they are recently proposed by several authors and we find that they cannot occur with a power-index of $-5$. As we show in Sect.~\\ref{sect3}, the proposed processes are not effective enough to produce the desired ion tails which means that another mechanism has to lead to the observed spectrum. In Sect.~\\ref{sect4}, we show how a high-energy source can be derived by taking a modulated ACR spectrum upstream of the solar wind termination shock. This injection mechanism is discussed in Sect.~\\ref{sect5} where we show that these high-energy ions can lead to power-law ion tails, however with a power-index of $-4$. The results are discussed and compared with observations in Sect.~\\ref{sect6} ",
        "conclusions": ""
    },
    "0808/0808.2016_arXiv.txt": {
        "abstract": "We review the present status of three-flavour neutrino oscillations, taking into account the latest available neutrino oscillation data presented at the {\\em Neutrino 2008\\/} Conference. This includes the data released this summer by the MINOS collaboration, the data of the neutral current counter phase of the SNO solar neutrino experiment, as well as the latest KamLAND and Borexino data. We give the updated determinations of the leading 'solar' and 'atmospheric' oscillation parameters. We find from global data that the mixing angle $\\theta_{13}$ is consistent with zero within $0.9\\sigma$ and we derive an upper bound of $\\sin^2\\theta_{13} \\leq 0.035 \\, (0.056)$ at 90\\%~CL (3$\\sigma$). \\vskip 1cm \\noindent Keywords: Neutrino mass and mixing; solar and atmospheric neutrinos; reactor and accelerator neutrinos ",
        "introduction": "\\label{sec:introduction}   Thanks to the synergy amongst a variety of experiments involving solar and atmospheric neutrinos, as well as man-made neutrinos at nuclear power plants and accelerators~\\cite{exp-talks-nu08} neutrino physics has undergone a revolution over the last decade or so. The long-sought-for phenomenon of neutrino oscillations has been finally established, demonstrating that neutrino flavor states $(\\nu_e,\\nu_\\mu,\\nu_\\tau)$ are indeed quantum superpositions of states $(\\nu_1,\\nu_2,\\nu_3)$ with definite masses $m_i$~\\cite{Amsler:2008zz}. The simplest unitary form of the lepton mixing matrix relating flavor and mass eigenstate neutrinos is given in terms of three mixing angles $(\\theta_{12},\\theta_{13},\\theta_{23})$ and three CP-violating phases, only one of which affects (conventional) neutrino oscillations~\\cite{schechter:1980gr}. Here we consider only the effect of the mixing angles in current oscillation experiments, the sensitivity to CP violation effects remains an open challenge for future experiments~\\cite{Bandyopadhyay:2007kx,Nunokawa:2007qh}. Together with the mass splitting parameters $\\Dms \\equiv m^2_2-m^2_1$ and $\\Dma \\equiv m^2_3- m^2_1$ the angles $\\theta_{12}, \\theta_{23}$ are rather well determined by the oscillation data. In contrast, so far only upper bounds can be placed upon $\\theta_{13}$, mainly following from the null results of the short-baseline CHOOZ reactor experiment \\cite{Apollonio:2002gd} with some effect also from solar and KamLAND data, especially at low $\\Dma$ values~\\cite{Maltoni:2003da}.  Here we present an update of the three-flavour oscillation analyses of Refs.~\\cite{Maltoni:2003da} and \\cite{Maltoni:2004ei}.  This new analysis includes the data released this summer by the MINOS collaboration~\\cite{Adamson:2008zt}, the data from the neutral current counter phase of the SNO experiment (SNO-NCD)~\\cite{Aharmim:2008kc}, the latest KamLAND~\\cite{:2008ee} and Borexino~\\cite{Collaboration:2008mt} data, as well as the results of a recent re-analysis of the Gallex/GNO solar neutrino data presented at the Neutrino 2008 conference~\\cite{gallex-nu08:07}. In Section~\\ref{sec:lead-solar-atmosph} we discuss the status of the parameters relevant for the leading oscillation modes in solar and atmospheric neutrinos. In Section~\\ref{sec:th13} we present the updated limits on $\\theta_{13}$ and discuss the recent claims for possible hints in favour of a non-zero value made in Refs.~\\cite{Balantekin:2008zm,Escamilla:2008vq,Fogli:2008jx}. We summarize in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this work we have provided an update on the status of three-flavour neutrino oscillations, taking into account the latest available world neutrino oscillation data presented at the {\\em Neutrino 2008\\/} Conference. Our results are summarized in Figures~\\ref{fig:dominant-12}, \\ref{fig:dominant-23} and \\ref{fig:th13}.  Table~\\ref{tab:summary} provides best fit points, $1\\sigma$ errors, and the allowed intervals at 2 and 3$\\sigma$ for the three-flavour oscillation parameters.  \\begin{table}[ht]\\centering \\catcode`?=\\active \\def?{\\hphantom{0}}  \\begin{tabular}{|@{\\quad}>{\\rule[-2mm]{0pt}{6mm}}l@{\\quad}|@{\\quad}c@{\\quad}|@{\\quad}c@{\\quad}|@{\\quad}c@{\\quad}|} \\hline parameter & best fit & 2$\\sigma$ & 3$\\sigma$ \\\\ \\hline $\\Delta m^2_{21}\\: [10^{-5}\\eVq]$ & $7.65^{+0.23}_{-0.20}$  & 7.25--8.11 & 7.05--8.34 \\\\[2mm] $|\\Delta m^2_{31}|\\: [10^{-3}\\eVq]$ & $2.40^{+0.12}_{-0.11}$  & 2.18--2.64 & 2.07--2.75 \\\\[2mm] $\\sin^2\\theta_{12}$ & $0.304^{+0.022}_{-0.016}$ & 0.27--0.35 & 0.25--0.37\\\\[2mm] $\\sin^2\\theta_{23}$ & $0.50^{+0.07}_{-0.06}$ & 0.39--0.63 & 0.36--0.67\\\\[2mm] $\\sin^2\\theta_{13}$ & $0.01^{+0.016}_{-0.011}$  & $\\leq$ 0.040 & $\\leq$ 0.056 \\\\ \\hline \\end{tabular} \\caption{ \\label{tab:summary} Best-fit values with 1$\\sigma$ errors, and 2$\\sigma$ and 3$\\sigma$ intervals (1 \\dof) for the three--flavour neutrino oscillation parameters from global data including solar, atmospheric, reactor (KamLAND and CHOOZ) and accelerator (K2K and MINOS) experiments.} \\end{table}   \\paragraph{Acknowledgments.} This work was supported by MEC grant FPA2005-01269, by EC Contracts RTN network MRTN-CT-2004-503369 and ILIAS/N6 RII3-CT-2004-506222.  We thank Michele Maltoni for collaboration on global fits to neutrino oscillation data.    \\appendix"
    },
    "0808/0808.0225_arXiv.txt": {
        "abstract": "{ Standard cosmology has many successes on large scales, but faces some fundamental difficulties on small, galactic scales. One such difficulty is the cusp/core problem. High resolution observations of the rotation curves for dark matter dominated low surface brightness (LSB) galaxies imply that galactic dark matter halos have a density profile with a flat central core, whereas N-body structure formation simulations predict a divergent (cuspy) density profile at the center. It has been proposed that this problem can be resolved by stellar feedback driving turbulent gas motion that erases the initial cusp. However, strong gravitational lensing prefers a cuspy density profile for galactic halos. In this paper, we use the most recent high resolution observations of the rotation curves of LSB galaxies to fit the core size as a function of halo mass, and compare the resultant lensing probability to the observational results for the well defined combined sample of the Cosmic Lens All-Sky Survey (CLASS) and Jodrell Bank/Very Large Array Astrometric Survey (JVAS). The lensing probabilities based on such density profiles are too low to match the observed lensing in CLASS/JVAS.  High baryon densities in the galaxies that dominate the lensing statistics can reconcile this discrepancy, but only if they steepen the mass profile rather than making it more shallow. This places contradictory demands upon the effects of baryons on the central mass profiles of galaxies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.0932_arXiv.txt": {
        "abstract": "We present the first resolved images of the eclipsing  binary $\\beta$ Lyrae, obtained with the CHARA Array interferometer and the MIRC combiner in the $H$ band. The images clearly show the mass donor and the thick disk surrounding the mass gainer at all six epochs of observation. The donor is brighter and generally appears elongated in the images, the first direct detection of photospheric tidal distortion due to Roche-lobe filling.  We also confirm expectations that the disk component is more elongated than the donor and is relatively fainter at this wavelength. Image analysis and model fitting for each epoch were used for calculating the first astrometric orbital solution for $\\beta$~Lyrae, yielding precise values for the orbital inclination and position angle. The derived semi-major axis  also allows us to estimate the distance of $\\beta$ Lyrae; however, systematic differences between the models and the images limit the accuracy of our distance estimate to about 15\\%. To address these issues, we will need a more physical, self-consistent model to account for all epochs as well as the multi-wavelength information from the eclipsing light curves. ",
        "introduction": "Interacting binaries are unique testbeds for many important astrophysical processes, such as mass and momentum transfer, accretion, tidal interaction, etc. These processes provide information on the evolution and properties of many types of objects, including low-mass black holes and neutron stars (in low-mass X-ray binaries), symbiotic binaries, cataclysmic variables, novae, etc.  Although these types of objects are widely studied by indirect methods such as spectroscopy, radial velocity, and sometimes eclipse mapping,  very few of them have been directly resolved because they are very close to each other and far away from us. Thus, directly imaging  interacting binaries, although very challenging, will greatly help us to improve our understanding of these objects.  The star $\\beta$ Lyrae (Sheliak, HD 174638, HR 7106, $V$ = 3.52, $H$=3.35) is a well known interacting and eclipsing binary that has been widely studied since its discovery in 1784 \\citep{Goodricke1785}.  According to the current picture \\citep{Harmanec2002}, the system consists of a B6-8~II Roche-lobe filling mass-losing star, which is generally denoted as the donor or the primary,  and an early B type mass-gaining star that is  generally denoted as the gainer or the secondary. The donor, which was initially more massive than the gainer, has a current mass of about 3 \\msun, while the gainer has a mass of about 13 \\msun. It is thought that the gainer is completely embedded in a thick accretion disk with bipolar jet-like structures perpendicular to the disk, which creates a light-scattering halo above its poles. The orbit of the system is highly circular \\citep{Harmanec1993}, and is very close to edge-on \\citep{Linnell2000}. Recent RV study on the  ephemeris of the system gives a period of $12.^d94$ \\citep{Ak2007}. The period is increasing at a rate of $\\sim19$ sec  per year due to the high mass transfer rate, $2\\times10^{-5}\\msun$ yr$^{-1}$, of the system.  The primary eclipse of the light curve (i.e., at phase 0) corresponds to the eclipse of the donor. In the $UBV$ bands, the surface of the donor is brighter than that of the gainer, and  therefore the primary minimum is deeper than the secondary minimum. At longer wavelengths, however, the studies of  \\citet{Jameson1976} and \\citet{Zeilik1982} suggest that the relative depth of the secondary minimum in the light curve gradually deepens and becomes deeper than the primary minimum at wavelengths longer than  $3.6\\mu m$.    Light curve studies and theoretical models have shown that, at the distance of 296pc \\citep{van-Leeuwen2007}, the estimated separation of the binary is only 0.92 milli-arcsecond (hereafter $mas$, $58.5\\rsun$). The angular diameter of the donor is  $\\sim$0.46 mas (29.4\\rsun), and the disk surrounding the gainer is only $\\sim1$ mas across \\citep[e.g.,][]{Linnell2000, Harmanec2002}. The goal of directly imaging $\\beta$ Lyr, therefore,  requires the angular resolution only achievable by today's long-baseline interferometers. Recently, \\citet{Schmitt2008} used the NPOI interferometer to image successfully the H$\\alpha$ emission of $\\beta$ Lyr, an update to the pioneering work of \\citet{Harmanec1996}.  Also, radio work using MERLIN found a nebula  surrounding the secondary but could not resolve its bipolar shape \\citep{Umana2000}. Despite recent progress, the individual objects of the system have not been resolved yet, putting even a simple astrometric orbit beyond our reach.  In this study, we present the first resolved images of the \\betlyr system at multiple phases, obtained with the CHARA Array and the MIRC combiner. We give a brief introduction to our observations and data reduction in \\S\\ref{observations}. We present our aperture synthesis images with simple models in \\S\\ref{imaging}. In \\S\\ref{orbit} we discuss our astrometric orbit of \\betlyr and we give the outlook for future work in \\S\\ref{summary}. ",
        "conclusions": ""
    },
    "0808/0808.0103_arXiv.txt": {
        "abstract": "% In this paper we present a number of metrics for usage of the SAO/NASA Astrophysics Data System (ADS). Since the ADS is used by the entire astronomical community, these are indicative of how the astronomical literature is used. We will show how the use of the ADS has changed both quantitatively and qualitatively. We will also show that different types of users access the system in different ways. Finally, we show how use of the ADS has evolved over the years in various regions of the world.  The ADS is funded by NASA Grant NNG06GG68G. ",
        "introduction": "The SAO/NASA Astrophysics Data System (hereafter ADS), is a digital library and a vital source for bibliographic information in astronomy. The vast majority of astronomical researchers in the world use the ADS on a daily or near-daily basis. The use of the ADS has not only changed quantitatively but also qualitatively. Initially almost exclusively used by professional astronomers, the ADS now also has become a public service through external, general search engines (like Google, Yahoo, Microsoft Live Search and Ask.com, to name a few). In~\\citet{henneken07} we observed that up to the middle of 2004, the number of ADS users doubled on a bi-yearly basis. Since the ADS started to be indexed by general search engines, the number of incidental users has increased dramatically. However, the number of typical users (more than 10 visits per month) has continued to follow the same growth pattern.  With different types of users come different types of use. A professional astronomer has different interests than an occasional user. One way of illustrating this is to look at the distribution of publication years for the literature people are interested in. We will also look at the diversity of ADS users from a geographical point of view. This will indicate whether increased Internet access actually results in an increase of ADS usage. This is particularly interesting with respect to aspects of the ``Digital Divide'' (see e.g.~\\citet{ITU07}). In the next section, we will describe the character of the data we are working with. The following section will show the results, which will then be discussed in section 4. ",
        "conclusions": "In terms of its audience, the ADS has not only changed quantitatively, but also qualitatively. Besides a steady growth of the ADS regular users, we observe a dramatic increase in incidental users. The ADS is per definition the gateway to online literature for scientists, used by virtually all professional astronomers on a daily basis. Since 2005 there is a growing role as a source of science education of the general public.  Comparing the group of ``ADS regulars'' with the group visiting the ADS via Google Scholar shows that the obsolescence curve for the latter is fairly flat, corresponding with reading behavior by people acquainting themselves with a subject. This means Google Scholar is not the right tool for staying up-to-date with the latest events in a field. Looking at how professional astronomers use the ADS shows that the obsolescence function for them closely follows the citation rate (as a function of paper age).  Although ADS usage increased in regions like the EU and the USA, the percentage of world usage has decreased for these regions. This is because the growth in World usage is mainly driven by regions with the biggest potential for growth. The density of Internet users reaches a saturation point in middle- and high-income regions at which point ADS usage increases at a slower rate. It is encouraging to see the rapid increase in Internet user density in low-income regions and a similar increase in the number of ADS users in those regions. It indicates that increased access to electronic information is being used and in this sense there is a narrowing of the ``Digital Divide'' for these regions. Whether this increased access also resulted in an increased scientific output needs further bibliometric research."
    },
    "0808/0808.2937_arXiv.txt": {
        "abstract": "A framework is outlined to assess Cepheids as potential cluster members from readily available photometric observations. A relationship is derived to estimate colour excess and distance for individual Cepheids through a calibration involving recently published HST parallaxes and a cleaned sample of established cluster Cepheids. Photometric {\\it V--J} colour is found to be a viable parameter for approximating a Cepheid's reddening. The non-universal nature of the slope of the Cepheid PL relation for {\\it BV} photometry is confirmed. By comparison, the slopes of the {\\it VJ} and {\\it VI} relations seem relatively unaffected by metallicity. A new Galactic Cepheid confirmed here, GSC 03729-01127 (F6-G1 Ib), is sufficiently coincident with the coronal regions of Tombaugh 5 to warrant follow-up radial velocity measures to assess membership. CCD photometry and O--C diagrams are presented for GSC 03729-01127 and the suspected cluster Cepheids AB Cam and BD Cas. Fourier analysis of the photometry for BD Cas and recent estimates of its metallicity constrain it to be a Population I overtone pulsator rather than a Type II s-Cepheid. AB Cam and BD Cas are not physically associated with the spatially-adjacent open clusters Tombaugh 5 and King 13, respectively, the latter being much older ($\\log \\tau \\simeq 9$) than believed previously. Rates of period change are determined for the three Cepheids from archival and published data. GSC 03729-01127 and AB Cam exhibit period increases, implying fifth and third crossings of the instability strip, respectively, while BD Cas exhibits a period decrease, indicating a second crossing, with possible superposed trends unrelated to binarity.  More importantly, the observed rates of period change confirm theoretical predictions. The challenges and prospects for future work in this area of research are discussed. ",
        "introduction": "The {\\it All Sky Automated Survey} \\citep[ASAS,][]{po00}, the {\\it Northern Sky Variability Survey} \\citep[NSVS,][]{wo04}, and {\\it The Amateur Sky Survey} \\citep[TASS,][]{dr06} have detected many new Cepheid variables through their photometric signatures, resulting in a valuable expansion of the Galactic Cepheid sample \\citep{sd04} once confirmed by spectroscopic observation. In the case of GSC 03729-01127 (TASSIV 6349369), a suspected Cepheid studied here, the variable may be an open cluster member and a potentially valuable calibrator for the Cepheid period-luminosity (PL) relation. Cepheids continue to provide the foundation for the universal distance scale, and such variables could serve as an efficient means of quantifying the extinction to Galactic and extragalactic targets.  The results of the seminal Hubble Space Telescope (HST) Key Project yielded a Hubble constant of $H_0=72\\pm8$ km s$^{-1}$ Mpc$^{-1}$ \\citep{fr01}, a value supported by cosmological constraints inferred from WMAP observations \\citep{sp07}. The HST results are tied to Large Magellanic Cloud (LMC) Cepheids, which advantageously provide a large sample of common distance. However, distance estimates for the LMC exhibit an unsatisfactorly large scatter \\citep{fr01,be02}, resulting in an uncertain zero-point. Moreover there exists a difference in metallicity between LMC Cepheid variables relative to both Galactic Cepheids and those in galaxies used for calibrating secondary distance candles, the effects of which remain actively debated. \\citet{ta03}, for example, suggest that the LMC Cepheid PL relation appears to characterize short-period Cepheids as too bright relative to their Galactic counterparts, and long-period Cepheids as too faint. Conversely, \\citet{vl07} and \\citet{be07} suggest, on the basis of revised Hipparcos and newly-derived HST parallaxes, that the slopes of the PL relations ($V,I$) for Cepheids in the Galaxy and the LMC are consistent to within their cited uncertainties. Indeed, the results presented in section \\ref{extragal}, complementing in part those of \\citet{fo07}, appear to confirm ideas put forth in each of the above studies, namely that the slope of the PL relation is not universal in certain passbands, and for {\\it VJ} and {\\it VI} constructed relations, any putative difference in slope arising from metallicity effects appears negligible in comparison with other concerns and uncertainties related to extragalactic observations. Nevertheless, a consensus has yet to emerge and a resolution to the above debate may be assisted by renewed efforts towards establishing Galactic Cepheids as cluster members, a connection that provides direct constraints on Cepheid luminosities, intrinsic colours, masses, metallicities, and pulsation modes.  \\citet{tu02} compiled an extensive list of suspected cluster Cepheids based upon preliminary analyses, but only a few cluster/Cepheid pairs have been studied with the necessary detail to determine the parameters of the associated clusters accurately, or to obtain the necessary radial velocity measures needed to establish membership in cases where reliable proper motions are unavailable. Efforts to discover new Galactic open clusters \\citep{al03,mo03,kro06} and Cepheid variables \\citep{po00,wo04,dr06} have resulted in a welcome increase to the number of suspected cluster Cepheids. This paper outlines a framework to assess the viability of such cases efficiently, with an intent to highlight cases requiring further attention and focus. Section \\ref{framework} develops a relationship to estimate colour excesses and distances for individual Cepheids from several photometric parameters. Section \\ref{cepheids} presents CCD photometry, spectroscopic results, and O--C analyses for the suspected cluster Cepheids BD Cas \\citep{ts66,tu02}, AB Cam \\citep{vb57,ts66}, and a new Galactic Cepheid confirmed here, GSC 03729-01127. Distances, colour excesses, and ages are also derived for the associated open clusters Tombaugh 5 and King 13 from 2MASS photometry \\citep{cu03}. ",
        "conclusions": "The relationships highlighted in section 3 yield reliable parameters when investigating short period Cepheids ($P\\le11^{\\rm d}$). Parameters determined for longer period Cepheids from such relations are less certain, primarily because of an absence of mid-to-long period calibrators needed to identify a unique set of co-efficients consistent over a broad period baseline. At present $\\ell$ Car is the only established long-period calibrator (parallax) used in deriving the co-efficients. A further drawback of the analysis rests in the adopted parameters for the calibrating clusters, which exhibit an unsatisfactory amount of scatter in the literature and more recent analyses \\citep[e.g.,][]{an07,ho03}. Refining distance estimates to the calibrating set of clusters by means of deep CCD photometry, analagous to the impressive results from the CFHT Open Cluster Survey \\citep{ka01a,ka01b}, is a priority in moving forward.  The framework is also tied to the HST sample of Cepheids with parallaxes and field reddenings established by \\citet{be07}.  It is noted that the parallax measures for RT Aur and Y Sge differ significantly between HST and Hipparcos \\citep[][see their table 1]{vl07}.  The framework outlined here should permit an efficient investigation of suspected cluster Cepheids \\citep{tu02}, including objects uncovered by cross correlations between newly discovered Cepheids in the ASAS and TASS with open cluster databases (i.e., WebDA). A potential goal is an expansion of the sample of cluster Cepheids, with particular emphasis on long period cluster Cepheids. Of equal importance, however, is the task of purging line-of-sight coincidences from current lists promulgating the literature, something that is particularly acute given that high-precision data are needed to address the question of the universality of the PL relation.  Four longer term objectives exist. First is to use the new relations to determine Galactic parameters and map interstellar extinction. Second, is to establish mean photometry for an entire calibrating set. Third, with regard to the universality of the PL relation and establishing long-period Cepheid calibrators, realistically it will be the highly anticipated results from the GAIA mission \\citep{ct06}, a next generation follow-up to the Hipparcos mission, that will provide the large and unbiased sample of Cepheid parallaxes needed to advance our knowledge of the field. In conjunction with a cleaned sample of cluster Cepheids, it should lead to a proper refinement of the relations outlined in section \\ref{framework}, and, consequently, the realization of the outlined objectives.  Fourth is the longer term prospect of conducting extragalactic surveys using JWST \\citep{ga06} to determine the distances and extinction to, and within [equation (\\ref{eqn4})], higher redshift galaxies. The aperture size and infrared sensitivity of the telescope will permit deeper sampling of extragalactic Cepheids, especially since the variables are substantially brighter in the infrared than the optical and the diminishment in flux from reddening is comparitively less.  \\subsection*{ACKNOWLEDGEMENTS} We are indebted to the following individuals and groups who helped facilitate the research: Alison Doane and the staff of the Harvard College Observatory Photographic Plate Stacks, Charles Bonatto for useful discussions on taking advantage of data from the 2MASS survey, Pascal Fouqu\\'{e}, Laszlo Szabados and Leonid Berdnikov, whose comprehensive work on evolutionary trends in Cepheid variables was invaluable in our analysis, Arne Henden and the staff at the AAVSO, Dmitry Monin, Les Saddelmeyer, and the rest of the staff of the Dominion Astrophysical Observatory, Doug Welch who maintains the McMaster Cepheid Photometry and Radial Velocity Archive, the staff at la Centre de Donn\\'{e}es astronomiques de Strasbourg, and Carolyn Stern Grant and the staff at the Astrophysics Data System (ADS). Reviews on Cepheids by Michael Feast, Donald Fernie, and Nick Allen were useful in the preparation of this work.   Lastly, we extend a special thanks to Sandra Hewitt for her exceptional kindness in accommodating visiting astronomers to the Harvard Plate Stacks."
    },
    "0808/0808.0759_arXiv.txt": {
        "abstract": "The comparison of the black hole mass function (BHMF) of active galactic nuclei (AGN) relics with the measured mass function of the massive black holes in galaxies provides strong evidence for the growth of massive black holes being dominated by mass accretion. We derive the Eddington ratio distributions as functions of black hole mass and redshift from a large AGN sample with measured Eddington ratios given by Kollmeier et al. We find that, even at the low mass end, most black holes are accreting at Eddington ratio $\\lambda\\sim0.2$, which implies that the objects accreting at extremely high rates should be rare or such phases are very short. Using the derived Eddington ratios, we explore the cosmological evolution of massive black holes with an AGN bolometric luminosity function (LF). It is found that the resulted BHMF of AGN relics is unable to match the measured local BHMF of galaxies for any value of (constant) radiative efficiency $\\eta_{\\rm rad}$. Motivated by Volonteri, Sikora \\& Lasota's study on the spin evolution of massive black holes, we assume the radiative efficiency to be dependent on black hole mass, i.e., $\\eta_{\\rm rad}$ is low for $M_{\\rm bh}<10^8{\\rm M}_\\odot$ and it increases with black hole mass for $M_{\\rm bh}\\ge 10^8{\\rm M}_\\odot$. We find that the BHMF of AGN relics can roughly reproduce the local BHMF of galaxies if $\\eta_{\\rm rad}\\simeq0.08$ for $M_{\\rm bh}<10^8{\\rm M}_\\odot$ and it increases to $\\ga 0.18$ for $M_{\\rm bh}\\ga10^9{\\rm M}_\\odot$, which implies that most massive black holes ($\\ga 10^9{\\rm M}_\\odot$) are spinning very rapidly. ",
        "introduction": "It is believed that quasars are powered by accretion on to massive black holes, and the growth of the massive black holes could be governed by mass accretion in quasars. The massive black holes (AGN relics) should be present in the centres of galaxies. Thus, the luminosity functions (LF) of active galactic nuclei (AGN) provide important clues on the growth of massive black holes. It is indeed found that most nearby galaxies contain massive black holes at their centres, and a tight correlation is revealed between central massive black hole mass and the velocity dispersion of the galaxy \\citep{fm00,g00}. The black hole mass is also found to be tightly correlated with the luminosity of the spheroid component of its host galaxy \\citep*[e.g.,][]{m98,mh03}. These correlations of the black hole mass with velocity dispersion/host galaxy luminosity were used to derive the mass functions of the central massive black holes in galaxies \\citep*[e.g.,][]{yt02,m04,t06,g07}. On the other hand, the black hole mass function (BHMF) of AGN relics can also be calculated by integrating the continuity equation of massive black hole number density on the assumption of the growth of massive black holes being dominated by mass accretion, in which the activity of massive black holes is described by a LF of AGN \\citep*[e.g.,][]{c71,soltan82,ct92,sb92,m04,s04,t06}. Such calculations on the cosmological evolution of massive black holes were usually carried out by adopting two free parameters: the radiative efficiency $\\eta_{\\rm rad}$ and the Eddington ratio $\\lambda$ for AGN. The derived BHMF of AGN relics in this way is required to match that estimated either from velocity dispersion or the luminosity of the spheroid of its host galaxy, which always leads to $\\lambda\\sim 1$, i.e., almost all AGN are required to be accreting close to the Eddington limit \\citep*[e.g.,][]{yt02,m04,s04,t06}.  In the last decade, several approaches for measuring the masses of the central black holes in AGN have been developed, in which the reverberation mapping may be the most effective one \\citep{p93,k00}. Using the tight correlation between the size of the broad-line region and the optical luminosity established with the reverberation mapping method for a sample of AGN, the black hole masses of AGN can be easily estimated from their optical luminosity and width of broad emission line. The Eddington ratios for thousands of AGN were estimated with the analyses of the Sloan Digital Sky Survey (SDSS) by \\citet{md04}, which indicate that the mean Eddington ratio $L_{\\rm bol}/L_{\\rm Edd}\\simeq 0.1$ at $z\\sim 0.2$ to $\\simeq0.4$ at $z\\sim 2$. \\citet{w04} also derived the Eddington ratio distribution for a sample of $\\sim$~500 AGN with redshifts $0\\la z\\la 5$. As pointed by \\citet{k06}, both these derived Eddington ratios are heavily weighted towards high-luminosity objects due to the limited sensitivity of SDSS. \\citet{k06} estimated the Eddington ratios of AGN discovered in the AGN and Galaxy Evolution Survey (AGES), which is more sensitive than the SDSS \\citep{k04}. The derived Eddington ratio distribution at {\\it fixed luminosity} is well described by a single lognormal distribution peaked at $\\sim 0.25$ independent of redshift and luminosity \\citep*[see][for the details]{k06}.  In this work, the Eddington ratio distribution at {\\it fixed luminosity} given by \\citet{k06} is converted to that at {\\rm fixed black hole mass} by using an AGN LF. We integrate the continuity equation for black hole number density adopting the derived Eddington ratio distributions of AGN to calculate the BHMF of AGN relics at different redshifts $z$, which is different from a free parameter $\\lambda$ adopted in most previous works. The resultant BHMF of AGN relics is constrained by those estimated from the galaxy LFs \\citep{m04,t06}. The conventional cosmological parameters $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70~ {\\rm km~s^{-1}~Mpc^{-1}}$ have been adopted in this work. ",
        "conclusions": "As in the most previous works, we implicitly assume that the black hole growth is dominated by mass accretion in bright AGN, while some inactive black holes may still be accreting gases, though their mass accretion rates are very low. If the duration of the accretion in these objects is as long as the Hubble timescale, they can accrete sufficient mass comparable with that accumulated in bright AGN phases, as the AGN phase is much shorter than the Hubble timescale. It is believed that the advection dominated accretion flows (ADAFs) are present in those objects, which are very hot and radiate mostly in hard X-ray bands \\citep{ny94}. They are very difficult to be detected due to low luminosity, unless those in the nearby Universe. \\citet{cao05} suggested that the accretion of such low-luminosity objects can be constrained by the hard X-ray background, though the emission from most of these individuals cannot be detected by any facilities now. It was found that less than $\\sim 5$ per cent of the local black hole mass density was accreted during the ADAF phases, which will be even lower if the Compton-thick AGN are included \\citep*[see][for the details]{cao07}. \\citet{h06} considered the distribution of local supermassive black hole Eddington ratios and accretion rates, accounting for the dependence of radiative efficiency and bolometric corrections on the accretion rate. They also found that black hole mass growth was dominated by AGN phase, and not by the radiatively inefficient low accretion rate phase in which most local supermassive black holes are currently observed.    The main difference of this work from the previous works is that the Eddington ratio distributions are derived from an AGN sample with measured Eddington ratios \\citep{k06}. The Eddington ratio distributions for fixed black hole mass derived in our work approximate to the lognormal distribution (see Fig. \\ref{fig1}), and the mean Eddington ratios are in the range of $\\sim 0.1-0.3$ varying with black hole mass and redshift (see Fig. \\ref{fig2}). For most cases, the mean Eddington ratios peak at $\\sim 10^8{\\rm M}_\\odot$, and then decline with increasing black hole mass. Even at the low mass end, most black holes are accreting at $\\lambda\\sim0.2$, which implies that the objects accreting at extremely high rates should be rare or such phases are very short. It was suggested that the radiative efficiency $\\eta_{\\rm rad}$ declines for a slim accretion disc provided the mass accretion rate is sufficiently high due to the photon trapping effort \\citep[e.g.,][]{a88,b78,w99}. \\citet{w00}'s calculations on the slim discs showed that the radiative efficiency will not deviate significantly from that for standard thin discs if $L_{\\rm bol}/L_{\\rm Edd}\\la 2$, which implies that the present adopted radiative efficiency independent of Eddington ratio $\\lambda$ is indeed a good assumption.  There is only one free parameter $\\eta_{\\rm rad}$ in our calculations for the cosmological evolution of massive black holes. We find that the resulted BHMF of AGN relics is unable to reproduce the measured local BHMF for any value of $\\eta_{\\rm rad}$ adopted, provided the radiative efficiency $\\eta_{\\rm rad}$ is independent of black hole mass, as treated in previous works \\citep*[e.g.,][]{yt02,m04,t06}. The mean Eddington ratios adopted in our calculations are derived from an AGN sample, which are in the range of $\\sim 0.1-0.3$. Thus, it is not surprising that the local BHMFs cannot be reproduced by our calculations with any constant radiative efficiency, because the Eddington ratio $\\lambda\\sim 1$ is usually required in order to let the resulted BHMF match the local one in those works. In this work, we use two different corrections (either luminosity-independent or luminosity-dependent) for the Compton-thick AGN (see Sect. 3 for the details), and find that the final results are quite similar (see Figs. \\ref{fig3} and \\ref{fig4}). We also use the hard X-ray LF derived from an AGN sample at high redshifts by \\citet{s08} in stead of the bolometric LF of \\citet{h07} in the calculations. It is found that the main results of this work change very little and the main conclusion is not altered.   \\citet{v07} studied on how the accretion from a warped disc influences the evolution of black hole spins and concluded that within the cosmological framework, one indeed expects most supermassive black holes in elliptical galaxies to have on average higher spin than black holes in spiral galaxies, where random, small accretion episodes (e.g., tidally disrupted stars, accretion of molecular clouds) might have played a more important role. Thus, we tentatively adopt a $M_{\\rm bh}$-dependent radiative efficiency (see Eq. \\ref{etarad}), in which $\\eta_{\\rm rad}$ remains constant for $M_{\\rm bh}\\le 10^8$ and increases with black hole mass for $M_{\\rm bh}> 10^8$. This $M_{\\rm bh}$-dependent radiative efficiency is qualitatively consistent with the results of \\citet{v07}. It is found that the measured BHMFs can be fairly well reproduced by our model calculations with this $M_{\\rm bh}$-dependent radiative efficiency (see Figs. \\ref{fig3} and \\ref{fig4}), which require $\\eta_{\\rm rad}\\ga 0.18$ for $M_{\\rm bh}\\ga 10^9{\\rm M}_\\odot$. This provides evidence for most massive black holes being spinning very rapidly. It is interesting to find that $a\\simeq 0.9$ is also required by the fitting of the residual hard X-ray background with the emission from the ADAFs in the low-luminosity objects \\citep{cao07}. Our calculations can be improved if the mean spin parameter $a$ as a function of black hole mass is available from the work within the cosmological framework, which is beyond the scope of present work.  {  In our present calculations of the black hole evolution, the black hole mergers have been neglected. \\citet{s07} assessed the importance of the black hole mergers on the evolution of the BHMF. They found that the impact of black hole mergers on the cosmological evolution of BHMF may probably be small compared with black hole accretion processes, while its impact on the black hole spin evolution may be important \\citep*[e.g.,][]{wc95,hb03,v05,v07}. The effect of black hole mergers increases the number density of very massive black holes \\citep{s07}, which implies that the radiative efficiencies for very massive active black holes should be higher than the present values if black hole mergers are included in our calculations. This strengthens our conclusion that most massive black holes are spinning very rapidly. }"
    },
    "0808/0808.3793_arXiv.txt": {
        "abstract": "The apparent spectral evolution observed in the steep decay phase of many GRB early afterglows raises a great concern of the high-latitude ``curvature effect'' interpretation of this phase. However, previous curvature effect models only invoked a simple power law spectrum upon the cessation of the prompt internal emission. We investigate a model that invokes the ``curvature effect'' of a more general non-power-law spectrum and test this model with the Swift/XRT data of some GRBs. We show that one can reproduce both the observed lightcurve and the apparent spectral evolution of several GRBs using a model invoking a power-law spectrum with an exponential cut off. GRB 050814 is presented as an example. ",
        "introduction": "} Most of the early X-Ray afterglows detected by Swift (Gehrels et al. 2004) show a steep decay phase around 100$\\sim$1000 seconds after the burst trigger (Tagliaferri et al. 2005). The main characteristics of this steep decay phase include the following. (1) It connects smoothly to the prompt $\\gamma$-ray light curve extrapolated to the X-ray band, suggesting that it is the ``tail'' of the prompt emission (Barthelmy et al. 2005, O'Brien et 2006, Liang et al 2006).  (2) The decay slope is typically $3 \\sim 5$ when choosing the GRB trigger time as the zero time point $t_0$ (Tagliaferri et al. 2005; Nousek et al. 2006; Zhang et al 2006). (3) The time-averaged spectral index of the steep decay phase is much different from that of the later shallow decay phase, indicating that it is a distinct new component that is unrelated to the conventional afterglow components (Zhang et al 2006; Liang et al. 2007). (4) Strong spectral evolution exists in about one third of the bursts that have a steep decay phase (Zhang et al. 2007, hereafter ZLZ07; Butler \\& Kocevski 2007; Starling et al. 2008). All these features suggest that the steep decay phase holds the key to understand the connection between the prompt emission (internal) phase and the traditional afterglow (external) phase. Any proposed model (see M\\'esz\\'aros 2006; Zhang 2007 for reviews) should be able to explain these features.  The so called ``curvature effect'', which accounts for the delayed photon emission from high latitudes with respect to the line of sight upon the abrupt cessation of emission in the prompt emission region (Fenimore et al. 1996; Kumar \\& Panaitescu 2000; Dermer 2004; Dyks et al. 2005; Qin 2008a), has been suggested to play an important role in shaping the sharp flux decline in GRB tails (Zhang et al. 2006; Liang et al. 2006; Wu et al. 2006; Yamazaki et al. 2006). In the simplest model, it is assumed that the instantaneous spectrum at the end of the prompt emission is a simple power law with a spectral index $\\beta$. The predicted temporal decay index of the emission is (with the convention $F_{\\nu} \\propto t^{-\\alpha} \\nu^{-\\beta}$) \\begin{equation} \\alpha=2+\\beta~, \\label{curvature-1} \\end{equation} if the time origin to define the $\\log-\\log$ light curve, $t_0$, is taken as the beginning of the last emission episode before the cessation of emission. Adopting a time-averaged $\\beta$ in the tails, Liang et al. (2006) found that Eq.(\\ref{curvature-1}) is generally valid. The strong spectral evolution identified in a group of GRB tails (ZLZ07) apparently violates Eq.(\\ref{curvature-1}), which is valid only for a constant $\\beta$. ZLZ07 then investigated a curvature effect model by assuming a structured jet with varying $\\beta$ at different latitudes and that the line of sight is near the jet axis\\footnote{Notice that this structured jet model is different from the traditional one that invokes an angle-dependent energy/Lorentz factor, but not the spectral index (Zhang \\& M\\'esz\\'aros 2002; Rossi et al. 2002).}. One would then expect that Eq.(\\ref{curvature-1}) is roughly satisfied, with both $\\alpha$ and $\\beta$ being time-dependent. ZLZ07 found that this model does not fit the data well.  These facts do not rule out the curvature effect interpretation of GRB tails, however. This is because the instantaneous spectrum upon the cessation of prompt emission may not be a simple power law. If the spectrum has a curvature, as the emission from progressively higher latitudes reach the observer, the XRT band is sampling different segments of the intrinsic curved spectrum (Fig.1). This would introduce an apparent spectral evolution in the decaying tail. The main goal of this paper is to test this more general curvature effect model using the available Swift XRT data. ",
        "conclusions": "We have successfully modeled the lightcurve and spectral evolution of the X-ray tail of GRB050814 using the curvature effect model of a cutoff power law spectrum with an exponential cutoff ($k=1$). It has been discussed in the literature (e.g. Fan \\& Wei 2005; Barniol-Duran \\& Kumar 2008) that the GRB central engine may not die abruptly, and that the observed X-ray tails may reflect the dying history of the central engine. If this is indeed the case, the strong spectral evolution in the X-ray tails would demand a time-dependent particle acceleration mechanism that gives a progressively soft particle spectrum. Such a behavior has not been predicted by particle acceleration theories. Our results suggest that at least for some tails, the spectral evolution is simply a consequence of the curvature effect: the observer views emission from the progressively higher latitudes from the line of sight, so that the XRT band is sampling the different segments of a curved spectrum. This is a simpler interpretation.  The phenomenology of the X-ray tails are different from case to case (ZLZ07). We have applied our model to some other clean X-Ray tails, such as GRB050724, GRB080523, and find that they can be also interpreted by this model. Some other tails have superposed X-ray flares, making a robust test of the model difficult. A systematic survey of all the data sample is needed to address what fraction of the bursts can be interpreted in this way or they demand other physically distinct models (e.g. Barniol-Duran \\& Kumar 2008; Dado et al. 2008). This is beyond the scope of this Letter."
    },
    "0808/0808.3046_arXiv.txt": {
        "abstract": "s{We present a powerful method for exploring various processes in the presence of strong external fields and matter. The method implies utilization of the exact solutions of the modified Dirac equations which contain the effective potentials accounting for the influences of external electromagnetic fields and matter on particles. We briefly discuss the basics of the method and its applications to studies of different processes, including a recently proposed new mechanism of radiation by neutrinos and electrons moving in matter (the spin light of the neutrino and electron). In view of a recent ``prediction'' of an order-of-magnitude change of the muon lifetime under the influence of an electromagnetic field of a CO$_2$ laser, we revisit the issue and show that such claims are nonrealistic.} ",
        "introduction": "The problem of particles' interactions under the influence of external electromagnetic fields and matter is one of the important topics in particle physics. Besides the possibility for better visualization of fundamental properties of particles and their interactions when they are influenced by external conditions, the interest to this problem is also stimulated by astrophysical and cosmological applications, where strong electromagnetic fields and dense matter may play important roles. There are well established methods for such kind of investigations that have a long-standing history.  In particular, the method of exact solutions of quantum equations, which is based on a Furry representation \\cite{FurPR51} of QED, is widely used in studies of particles' interactions in external electromagnetic fields. In this technique, the evolution operator $U_{F}(t_1, t_2)$, which determines the matrix element of the process, is presented in the usual form \\begin{equation} U_{F} (t_1, t_2)= T exp \\left[-i \\int \\limits_{t_1}^{t_2}j_{\\mu}(x) A^{\\mu}{d}x \\right], \\end{equation} where $A_{\\mu}(x)$ is the quantized part of the potential corresponding to the radiation field, which is accounted for within the perturbation-series techniques. At the same time, the electron (a charged particle) current is presented as \\begin{equation} j_{\\mu}(x)={e \\over 2}\\left[\\overline \\Psi_e \\gamma _{\\mu}, \\Psi_e \\right], \\end{equation} where $\\Psi_e$ are the exact solutions of the Dirac equation for an electron in the presence of an external electromagnetic field given by the classical non-quantized potential $A_{\\mu}^{ext}(x)$: \\begin{equation}\\label{D_eq_QED} \\left\\{ \\gamma^{\\mu}\\left(i\\partial_{\\mu} -eA_{\\mu}^{cl}(x)\\right) - m_e \\right\\}\\Psi_e (x)=0. \\end{equation} Note that within this approach the interaction of charged particles with an external electromagnetic field is taken into account exactly while the radiation field is allowed for by perturbation-series expansion techniques. That is why the method discussed is often referred to as the ``method of exact solutions''. The detailed discussion on foundations of this method and its applications to different processes such as, for instance, the synchrotron radiation by an electron in magnetic fields can be found in~\\cite{SokTerSynRad68,ritus87,nikishov87}. Many results, which ate very important for astrophysical applications, have been obtained within the discussed method when considering the neutron beta-decay in a constant magnetic field. These studies have been started in~\\cite{Kor64TerLysKor65}. Short reviews on the studies of beta-decay and the related cross symmetric processes in strong magnetic fields can be found in~\\cite{ShiStuP05KouStuPRC05}. ",
        "conclusions": ""
    },
    "0808/0808.1872_arXiv.txt": {
        "abstract": "Planetary nebulae (PN) are an essential tool in the study of the chemical evolution of the Milky Way and galaxies of the Local Group, particularly the Magellanic Clouds. In this work, we present some recent results on the determination of chemical abundances from PN in the Large and Small Magellanic Clouds, and compare these results with data from our own Galaxy and other galaxies in the Local Group. As a result of our continuing long term program, we have a large database comprising about 300 objects for which reliable abundances of several elements from He to Ar have been obtained. Such data can be used to derive constraints to the nucleosynthesis processes in the progenitor stars in galaxies of different metallicities. We also investigate the time evolution of the oxygen abundances in the SMC by deriving the properties of the PN progenitor stars, which include their  masses and ages. We have then obtained an age-metallicity relation taking into account both oxygen and [Fe/H] abundances. We show that these results have an important consequence on the star formation rate of the SMC, in particular by suggesting a star formation  burst in the last 2-3 Gyr. ",
        "introduction": "Planetary nebulae (PN) are an essential tool in the study of the chemical evolution of the Milky Way and galaxies of the Local Group, particularly the Magellanic Clouds (see for example \\cite[Maciel et al. 2006a]{maciel2006a}, \\cite[Richer \\& McCall 2006] {richer}, \\cite[Buzzoni et al. 2006] {buzzoni}, \\cite[Ciardullo 2006)] {ciardullo}. As the offspring of stars within a reasonably large mass bracket (0.8 to about 8 solar masses), PN encompass an equally large age spread, as well as different spatial and kinematic distributions. They usually present bright emission lines and can be easily distinguished from other emission line objects, so that their chemical composition and spatiokinematical properties are relatively well determined. In this work, we present some recent results on the determination of chemical abundances from PN in the Large and Small Magellanic Clouds, and compare these results with data from our own Galaxy and other galaxies in the Local Group. We also investigate the time evolution of the oxygen abundances in the  SMC by deriving an age-metallicity relation for this object. ",
        "conclusions": ""
    },
    "0808/0808.3336_arXiv.txt": {
        "abstract": "We present the atlas of protoplanetary disks in the Orion Nebula based on the ACS/WFC images obtained for the \\emph{HST Treasury Program on the Orion Nebula Cluster}. The observations have been carried out in 5 photometric filters nearly equivalent to the standard $B, V, H\\alpha, I,$ and $ z$ passbands. Our master catalog lists 178 externally ionized proto-planetary disks (\\emph{proplyds}), 28 disks seen only in absorption against the bright nebular background (\\emph{silhouette disks}, 8 disks seen only as dark lanes at the midplane of extended polar emission (\\emph{bipolar nebulae} or \\emph{reflection nebulae}) and 5 sources showing jet emission with no evidence of neither external ionized gas emission nor dark silhouette disks. Many of these disks are associated with jets seen in $H\\alpha$ and circumstellar material detected through reflection emission in our broad-band filters; approximately 2/3 have identified counterparts in x-rays. A total of 47 objects (29 proplyds, 7 silhouette disks, 6 bipolar nebulae, 5 jets with no evidence of proplyd emission or silhouette disk) are new detections with HST. We include in our list 4 objects previously reported as circumstellar disks which have not been detected in our HST/ACS images either because they are hidden by the bleeding trails of a nearby saturated bright star or because of their location  out of the HST/ACS Treasury Program field. Other 31 sources previously reported as extended objects do not harbor a stellar source in our HST/ACS images. We also report on the detection  of 16 red, elongated sources. Their location at the edges of the field, far from the Trapezium Cluster core ($\\gtrsim 10'$), suggests that these are probably background galaxies observed through low extinction regions of the Orion Molecular Cloud OMC-1. ",
        "introduction": "The Orion Nebula (M42, NGC 1976) is a unique laboratory for studying the physical processes related to star and planet formation. It harbors one of the richest and youngest clusters (Orion Nebula Cluster, ONC)\\ in the solar neighborhood, spanning the full spectrum of stellar and sub-stellar masses down to a few Jupiter masses \\citep{Luca00}. In 1979 several compact photoionized knots were firstly detected in the central region of the Orion Nebula as emission-line sources \\citep{Laqu79}, and then important follow-up studies were made in radio \\citep{Gara87, Chur87} and via emission-line spectroscopy \\citep{Meab88, Meab93, Mass93}. Since the early 1990's, \\emph{Hubble Space Telescope} (HST) observations of the ONC have been fundamental for clarifying the main characteristics of these young stellar objects (YSO) and their accretion disks. After the pioneering surveys of \\citet{Odel93} and \\citet{Pros94}, performed with the spherically-aberrated WF/PC, \\citet{Odel94} used WFPC2 to discover several externally ionized proto-planetary disks (\\emph{proplyds}), as well as a number of disks seen only in absorption against the bright nebular background (\\emph{silhouette disks}), both rendered visible by their location in or near the core of the \\ion{H}{2} region. Following this discovery, other HST programs have increased the number of known objects \\citep{Odel96, McCa96, Ball98, Ball00}. \\citet{Odel01} and \\citet{Smit05}, targeting areas out of the core, showed that these systems are ubiquitous across the Great Orion Nebula.  So far a total of $\\sim$ 200 silhouette disks and bright proplyds has been revealed by the HST observations of the Orion Nebula, the large majority through narrow-band filters centered on the H$\\alpha$ $\\lambda 6563$ emission lines, and occasionally through filters centered on the [N II] $\\lambda 6583$, [O I] $\\lambda 6300$, [O III] $\\lambda 5007$ and [S II] $\\lambda 6717+6731$ lines.  In this paper we present an atlas of multi-color observations of circumstellar disks and resolved circumstellar emission obtained with the Wide Field Channel of the Advanced Camera for Surveys (ACS/WFC). These images are part of the \\emph{HST Treasury Program on the Orion Nebula Cluster} (Cycle 13, GO Program 10246, P.I. M. Robberto), aimed at measuring with high precision the main stellar parameters of the cluster members. For this reason, the Treasury Program used broad-band filters to obtain the most accurate photometry of each source, together with H$\\alpha$ narrow-band images to address the presence of circumstellar emission that may contaminate the photometry and the point spread function of the broad-band data. The combination of broad-band and narrow-band images opens a new window on the study of disks in the OMC. It also makes it possible to detect disks where the nebular background is too faint, thanks to the light of the central stars reflected by the circumstellar material at the disk's polar regions (reflection nebulae).  After a brief description of the observations (\\S 2), we present the new ACS/WFC images of all circumstellar disks  (\\S 3). We then provide a complete catalogue of circumstellar disks in the Orion Nebula, including also the few disks that were not detected in our programs for a variety of reasons (\\S 4). Finally, after a brief description of the new proplyds, silhouette disks, and bipolar nebulae (\\S 5, \\S 6) we present the images of 16 red, elongated, and diffuse objects which most probably represent galaxies seen through the background curtain provided by the Orion Molecular Cloud (OMC-1, \\S 7). A few remarkable objects are being investigated and will be discussed in separate papers (see e.g. \\citet{Robb08b} on 124-132). ",
        "conclusions": "In this paper we have shown the HST/ACS images of the 178 proplyds, 28 disks seen only in silhouette, 8 reflection nebulae without external ionized plasma, 5 jets without neither external ionized plasma nor silhouette disk and 16 other extended objects observed by the HST Treasury Program on the Orion Nebula Cluster. For every object we have reported all the images taken through the 5 photometric filters used by the HST Treasury Program (F435W, F555W, F658N, F775W, F850LP).  The fact that most of these objects are associated to X-rays sources observed by the Chandra Orion Ultradeep Project is particularly interesting, since high energy photons could play an important role in the star and planet formation processes.  Among all the objects reported, 63 have been discovered by these images: 29 proplyds, 7 silhouette disks, 6 reflection nebulae with no external ionized plasma, 5 jets with no external ionized plasma or silhouette disk, and 16 other elongated object.  Searching in the literature we found that 4 objects previously reported as circumstellar disks have not been detected by HST/ACS images either because hidden by the saturation bleeding trails of a close bright star or because located out of HST/ACS Treasury Program field of view. For other 30 sources previously reported as extended objects HST/ACS images reveal no circumstellar emission around them.  A brief description of all the newly discovered proplyds, disks seen only in silhouette and reflection nebulae with no external ionized plasma has been carried out in \\S 5 and \\S 6.  Finally, we have discussed possible interpretations for the nature of the 16 extended objects. Because of their location far from the Trapezium Cluster ($\\gtrsim 10'$) and because of their red color, they are probably background galaxies reddened by the Orion Molecular Cloud OMC-1, but the alternative hypothesis of reflection nebulae turned on by red pre-main-sequence stars cannot be ruled out by our observations only. \\\\ \\\\ We wish to thank the referee, William Henney, for his useful comments that greatly improved the manuscript. Support for program 10246 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. The work of LR at STScI was done under the auspices of the STScI Summer Student Program."
    },
    "0808/0808.0515_arXiv.txt": {
        "abstract": "{Using observed GALEX far-ultraviolet (FUV) fluxes and VLA images of the 21-cm \\HI\\ column densities, along with estimates of the local dust abundances, we measure the volume densities of a sample of actively star-forming giant molecular clouds (GMCs) in the nearby spiral galaxy M83 on a typical resolution scale of 170 pc.\\\\ Our approach is based on an equilibrium model for the cycle of molecular hydrogen formation on dust grains and photodissociation under the influence of the FUV radiation on the cloud surfaces of GMCs.\\\\ We find a range of total volume densities on the surface of GMCs in M83, namely 0.1 - 400 cm$^{-3}$ inside $R_{25}$, 0.5 - 50 cm$^{-3}$ outside $R_{25}$. Our data include a number of GMCs in the \\HI\\ ring surrounding this galaxy. Finally, we discuss the effects of observational selection, which may bias our results.} ",
        "introduction": "This paper aims to measure the total gas densities in a sample of giant molecular clouds (GMCs) in the nearby spiral M83 (NGC 5236), using a method initially proposed by \\citet{1997ApJ...487..171A}. This method is based on the simple (but unavoidable) fact that ultraviolet photons from young, newly-formed stars will react back on the surrounding parent GMCs, dissociating the molecular gas on their surfaces and turning the (virtually invisible) \\Htwo\\ into its easily-detected atomic form. The motivation for this approach was provided by the discovery of \\citet{1986Natur.319..296A} that the \\HI\\ delineating the spiral arms in a nearby galaxy showed a large-scale morphology that was more consistent with photodissociation near the \\HII\\ regions than it was with compression of the \\HI\\ farther upstream in the spiral shock. The presence of atomic hydrogen is then indicative of a photodissociation region (PDR). The method was first applied in some detail to M101 by \\citet{2000ApJ...538..608S} and more recently to M81 by \\citet{2008ApJ...673..798H}.  The large, nearby spiral M83 has been a frequent target of searches for molecular gas using the CO(1-0) spectral line, which is easily detectable in M83. However, such studies have only recently begun to be carried out with sufficient linear (spatial) resolution to discern differences in the location of molecular gas and the emissions from young, hot stars. For example, \\citet{1991ApJ...381..130L} describe the presence of molecular (CO) emission ~300 pc downstream from M83's eastern spiral arm dust lane, detected with a $5\\arcsec\\ \\times 12\\arcsec$ beam. \\citet{1999ApJ...513..720R} find that the CO emission is spatially separated from the dust lane as well as the young stars on scales of a few hundred parsec. They suggest UV heating, cosmic-ray heating or a two-component molecular phase to explain this morphology. Their observations show features that show similarity to the largest GMCs in the Milky Way with masses on the order of several millions of solar masses. \\citet{1993ApJ...414...98W} explore the CO content in the central kiloparsec of M83 and infer a low-density component ($n_{H_2} \\lesssim 10^3-10^4 \\rm{cm}^{-3}$) and a warm (above 50 K), high-density component ($n_{H_2} \\gtrsim 10^4-10^5 \\rm{cm}^{-3}$). \\citet{2002AJ....123.1892C} study CO(1-0), CO(2-1) and neutral gas in M83, observing a strong truncation of the molecular disk at 6\\arcmin\\ accompanied by a warped atomic outer disk. Using the conventional assumptions about how to convert CO surface brightness to \\Htwo\\ column density, they conclude that roughly 80\\% of the total gas mass in M83 is \\Htwo. \\correction{The} \\HI\\ 21 cm and CO surface brightness are found to be correlated, but with a complex pattern of offsets. They speculate that the temperature of the CO gas is $>$ 20K in the nucleus and $<$ 7K in the outer disk. \\citet{2004A&A...413..505L} give a full overview of previous observations and present full CO (1-0) and CO (2-1) maps based on thousands of telescope pointings, but with a modest spatial resolution of $\\sim 1$ kpc. They conclude that the molecular gas spiral arms mostly trace the dust lanes. They expect the \\Htwo\\ mass to dominate that of \\HI\\ within 7.3 kpc of the center. At their linear resolution, CO and \\HI\\ emissions are correlated strongly within the optical disk. \\citet{2005A&A...441..491V} look into various tracers of star formation, including Polycyclic Aromatic Hydrocarbon (PAH) emission lines, another potential indicator of PDRs. Various combinations of these lines prove to be good tracers of star forming regions in M83. They are found predominantly in the spiral arms. We studied the occurrence of PAHs near PDRs in M81 \\citep{2008ApJ...673..798H} previously and confirmed that they are found near the star forming regions in almost all cases. Similar M83 data was not available at this time.  \\begin{figure*}[t!] \\centering \\includegraphics[width=17cm]{locplot} \\caption{The location of the candidate PDRs are shown with the GALEX FUV image in the background. Not all FUV sources are visible on this image, but all sources are sources of FUV radiation. The inner circle signifies a 1 kpc galactocentric radius, the outer one signifies $R_{25}$.} \\label{fig:locplot} \\end{figure*}  \\citet{2005ApJ...619L..79T} find and discuss sites of recent star formation in the extreme outer disk of M83, associated with the warped \\HI\\ disk of M83, which raises questions about star formation efficiency and modes of star formation. The nature of this outer disk is investigated further in \\citet{2007ApJ...661..115G}, who find evidence that individual young stars are responsible for the UV emission in the outer regions. \\citet{2007AJ....134..135Z} find that these UV sources are quite common out to $2 \\times R_{25}$. In this paper we include candidate PDRs outside the main optical disk of M83 in an effort to probe the amount of available gas in this environment.  The suspected high densities of molecular gas in M83, the morphology of the gas in the spiral arms, and the evidence of recent star formation in the outer regions, make this galaxy an excellent target for a study of photodissociated atomic hydrogen and the volume densities of M83's GMCs using our method.  In \\S \\ref{sec:method}, we explain our approach and the data we used. In \\S \\ref{sec:results} we present our results. We discuss and summarize these results in \\S \\ref{sec:discussionconlusions}. ",
        "conclusions": "\\label{sec:discussionconlusions}  In this paragraph we compare our results to \\Htwo\\ densities based on CO measurements. Then we discuss background issues, extinction and dust issues. We end with a summary of our findings.  \\subsection{CO results comparison} \\label{sub:CO}  The study of \\citet{1999ApJ...513..720R} is detailed enough to find typical sizes and masses of GMCs in the eastern arm of M83. These results can be compared to ours with some additional assumptions. Using a CO(1-0) map with kinematic information, they find CO masses and virial masses onindicatingindicating the order of $1 \\times 10^6 M_{\\odot}$. For example, their source 9 (their Table 4) yields a mass derived from CO emission of 7.3 $\\times 10^6 M_\\odot$ and 3.6 $\\times 10^6 M_\\odot$ for the virial masses, at a 50-80\\% uncertainty. This source corresponds to our FUV source no. 13. Since our method does not yield any information on cloud sizes (we only observe HI on the surface of GMCs), nor on kinetic temperatures in these clouds, we will take typical values of our measured volume densities and compare densities using a typical GMC radius of 75 pc for a spherical cloud. $n = 2 n_{H_2}$ inside the GMC and a typical density near our source no. 13 is $20~\\rm{cm}^{-3}$, or $n_{H_2} = 10~\\rm{cm^{-3}}$. 7.3 $\\times 10^6 M_\\odot$ is equivalent to about $n_{H_2} = 85 \\rm{cm^{-3}}$, which is in reasonable agreement with our results considering the uncertainties in both results (70-80\\%). More recent unpublished data (Lord, private communication, 2008) may be quantitatively different and more detailed, but does not yield substantially different results at the resolution of the data that were used in this paper.  One of the key differences between M81 and M83 is the abundance of CO emission. While CO is extremely faint and hard to detect in M81 \\citep[e.g.][]{2007A&A...473..771C}, M83 displays bright CO features. The method we used here to find hydrogen densities (and molecular hydrogen densities in GMCs) does not yield any morphological information, since we are observing \\HI\\ on the surface of PDRs. We can note, however, that the size and scale we assume for our candidate PDRs (the range of values of $\\rho_{\\HI}$) is consistent with findings using CO. Since we did not find suitable candidate PDRs in the inner kiloparsec of M83, we cannot compare our results with for example the results of \\citet{1993ApJ...414...98W}, although the densities we find match their low-density component. As to the truncation of the molecular disk that was mentioned by \\citet{2002AJ....123.1892C}, we do see a difference between the inner and outer disk of M83, but it does not seem to be that significant in our results. No differences are seen in candidate PDRs in arm or inter-arm regions, insofar as the large scatter in our results allows us to draw this conclusion. We do not find more than a handful of regions with the densities indicated by \\citet{2007ApJ...664..363H}, although the expected sizes of our candidate PDRs are the same. Finally, the hydrogen densities presented here are similar to the M81 results but extend to higher maximum values. This is consistent with the brighter CO emission in M83, since the higher \\Htwo\\ densities will lead to a greater degree of excitation of the CO molecules. Alternatively, this could also indicate a higher fraction of CO molecules in M83's clouds \\citep[e.g.][]{2007A&A...473..771C}.  \\subsection{Background levels, extinction and dust} \\label{sub:FUV}  The FUV fluxes seem to be higher in the inner parts of M83 and decreasing (on average) going outward, as is shown in Fig. \\ref{fig:F}. These higher fluxes in the inner parts could be caused by brighter sources, and/or a larger number of sources. Since the individual sources are unresolved, no firm conclusions can be drawn here. Outside $R_{25}$ the measured fluxes are roughly constant. At the same time, the abundances of PDR-produced \\HI\\ surrounding these FUV sources (Fig. \\ref{fig:HI} - note that these are individual measurements, not annular averages) seem to follow a similar trend, consistent with its connection to the photodissociating UV radiation. The FUV source contrast does not vary with galactocentric radius, so while both source and background radiation field decrease in intensity towards the outer regions of M83, their relative strength does not change.  Another factor that could influence the FUV radiation is galactic foreground extinction. The higher the applied foreground extinction, the higher the resulting total hydrogen volume density. The 0.52 mag extinction \\citep{1998ApJ...500..525S} towards M83 is significantly less than the 1.37 mag we used towards M81 in \\citet{2008ApJ...673..798H}. The \\citet{1998ApJ...500..525S} extinction correction for M81 would be 0.58, significantly lowering the total hydrogen volume densities. We find a \\correction{wider} range of values in M83, including higher gas densities. We therefore expect M83 to harbor more gas than M81. The M81 and M83 results are hard to compare directly because the extinction corrections were based on different sources in the literature. We intend to compare the two galaxies, together with similar M33 results, more quantitatively in a future paper, with consistent extinction corrections.  Our results are also sensitive to the dust-to-gas ratio. The slope of M83's metallicity is relatively shallow. Our preferred dust model assumes that metallicities in the outer parts of M83 remain high. Lower metallicities would result in higher gas densities. Our results in these regions are therefore most likely lower limits and much more gas could be present. The recent results by \\citet{2007ApJ...661..115G} are ambiguous in the sense that the authors provide high metallicity and low metallicity results, depending on the adopted model. In the high metallicity scenario, the slope of M83's metallicity remains shallow out to large galactocentric radius. The low metallicity scenario is accompanied by a sharp drop in metallicity starting at approximately $R_{25}$.  \\citet{2000ApJ...538..608S} found a similar range of gas densities in M101 using basically the same method with additional extinction corrections. Their candidate PDRs were all within M101's $R_{25}$, but the FUV luminosities within that range are similar. As in our previous M81 results, no internal extinction correction was applied. If any such correction were applied, it should scale with the dust-to-gas ratio. This means that the extinction would decrease going outward. The FUV fluxes would be even higher in the center of M83 (leading to higher total hydrogen volume densities), while the fluxes in the outer regions of M83 would not be affected much.  \\subsection{Summary} \\label{sub:summary}  In summary, we have investigated atomic hydrogen found in candidate PDRs in M83 and used the physics of these PDRs to derive total hydrogen volume densities. We carefully considered the contribution of observational selection effects to our results. We find a range of densities: 0.1 - 400 cm$^{-3}$ inside $R_{25}$, 0.5 - 50 cm$^{-3}$ outside $R_{25}$, based on measurements of \\HI\\ believed to be produced in large scale PDRs in M83. The higher GMC volume densities which we find within $R_{25}$ correlate well with the presence of bright CO(1-0) emission there. This points to enhanced collisional excitation as one reason for the CO emission in these GMCs. We note that this is also consistent with our results in M81, where we find little evidence for high GMC volume densities, and for which the CO(1-0) emission is faint.  Our measurements go out to a galactocentric radius of 13 kpc (deprojected). Our study used the tilted ring model from \\citet{1974ApJ...193..309R} to get proper galactocentric radii outside the M83 optical disk. Our results are notably sensitive to the local dust-to-gas ratio, especially since the metallicity in the outer regions remains uncertain."
    },
    "0808/0808.3385_arXiv.txt": {
        "abstract": "% RCW~38 is a uniquely young ($<$1 Myr), embedded ($A_V \\sim 10$) stellar cluster surrounding a pair of early O stars ($\\sim$O5.5) and is one of the few regions within 2~kpc other than Orion to contain over 1000 members. X-ray and deep near-infrared observations reveal a dense cluster with over 200 X-ray sources and 400 infrared sources embedded in a diffuse hot plasma within a 1~pc diameter.  The central O star has evacuated its immediate surroundings of dust, creating a wind bubble $\\sim$0.1 pc in radius that is confined by the surrounding molecular cloud, as traced by millimeter continuum and molecular line emission. The interface between the bubble and cloud is a region of warm dust and ionized gas, which shows evidence for ongoing star formation.  Extended warm dust is found throughout a  2--3 pc region and coincides with extended X-ray plasma.  This is evidence that the influence of the massive stars reaches beyond the confines of the O star bubble. RCW~38 appears similar in structure to RCW~49 and M~20 but is at an earlier evolutionary phase.  RCW~38 appears to be a blister compact H{\\small II} region lying just inside the edge of a giant molecular cloud. ",
        "introduction": "%  The evolution of high mass clustered star forming regions is complex and poorly understood.  Only the nearby ($\\sim$400 pc), optically revealed, Orion Nebula Cluster (ONC) is well studied (see Muench et al.\\ and  O'Dell et al.\\ in this Handbook).  Yet, a wide variety of high mass embedded clusters is found within 2 kpc of the Sun.  Within this limit, the young cluster RCW~38 (08$^h$59$^m$47.2$^s$ -47$^\\circ$31$'$57$''$ (J2000), {\\it l,b} = $268.03^\\circ,-0.98^\\circ$) is one of the few regions other than the ONC to contain over 1000 members (Lada \\& Lada 2003; Wolk et al.\\ 2006).  RCW~38 has an embedded and dense stellar population, comparable to other regions that have been studied recently with Spitzer (e.g., M~20 and RCW~49; Rho et al.\\ 2004, Whitney et al.\\ 2004, Churchwell et al.\\ 2004).  RCW~38 provides a unique opportunity to study the evolution of a rich cluster during the phase where its most massive members, a pair of O5.5 stars (DeRose et al.\\ 2008) have just completed their ultracompact {H}{\\small II}\\ region (UC{H}{\\small II}) phase and are now greatly influencing its natal environment and the evolution of its low mass members. ",
        "conclusions": "RCW~38 appears to be a blister compact HII region with winds originating from the central O-star IRS~2.  These winds are excavating the mass in the immediate vicinity of IRS~2 creating a shell--like structure, detected as a radio continuum ring and also evident at infrared and millimeter wavelengths.  A few hundred young low mass stars are found in the immediate vicinty of IRS~2, and may be directly exposed to its winds and ionizing radiation.  The region as a whole is estimated to contain around 2000 young stars, with 30 OB star candidates.  The region of the ring to the west of IRS~2, containing IRS~1, appears to be particularly active and is likely the site of ongoing (triggered) star formation.  The IRS~1 ridge appears to be the interface between the IRS~2 wind and either a similar wind from a high mass star further to the west, or a dense clump of gas that is being ablated by IRS~2.  Further observations are needed to test these scenarios. In many ways, RCW~38 appears to be a younger more embedded version of the ONC and deserves further study at higher angular resolution and across the spectrum."
    },
    "0808/0808.1349_arXiv.txt": {
        "abstract": "After the discovery that supermassive black holes (SMBHs) are ubiquitous at the center of stellar spheroids and that their mass $\\Mbh$, in the range $10^6\\Msun -10^9\\Msun$, is tightly related to global properties of the host stellar system, the idea of the co-evolution of elliptical galaxies and of their SMBHs has become a central topic of modern astrophysics. Here, I summarize some consequences that can be derived from the galaxy scaling laws and present a coherent scenario for the formation and evolution of elliptical galaxies and their central SMBHs, focusing in particular on the establishment and maintenance of their scaling laws.  In particular, after a first observationally based part, the discussion focuses on the physical interpretation of the Fundamental Plane. Then, two important processes in principle able to destroy the galaxy and SMBH scaling laws, namely galaxy merging and cooling flows, are analyzed. Arguments supporting the necessity to clearly distinguish between the origin and maintenance of the different Saling Laws, and the unavoidable occurrence of SMBH feedback on the galaxy Interstellar Medium in the late stages of galaxy evolution (when elliptical galaxies are sometimes considered as ``dead, red objects''), are then presented.  At the end of the paper I will discuss some implications of the recent discovery of super-dense ellipticals in the distant Universe.  In particular, I will argue that, if confirmed, these new observations would lead to the conclusion that at early epochs a relation between the stellar mass of the galaxy and the mass of the central SMBH should hold, consistent with the present day Magorrian relation, while the proportionality coefficient between $\\Mbh$ and the scale of velocity dispersion of the hosting spheroids should be significantly smaller than that at the present epoch. ",
        "introduction": "The mutual interplay between supermassive black holes (hereafter SMBHs) at the center of stellar spheroids\\footnote{The term {\\it early-type galaxies} is generically used for galaxies belonging to the family of {\\it elliptical} galaxies (Es), $S0$ galaxies, {\\it dwarf ellipticals} (dE), and {\\it dwarf spheroidals} (dSph); the class of {\\it stellar spheroids} is made of early-type galaxies and {\\it bulges} of spiral galaxies. For a detailed account of the observational properties of these classes see, e.g.,~\\cite{ref:bm98,ref:bert00}.} and their host systems is now established beyond any reasonable doubt, as indicated by the remarkable correlations found between host galaxy properties and the masses of their SMBHs (e.g.~\\cite{ref:mag98}-\\cite{ref:nfd06}). More specifically, it is now believed that all early-type galaxies with $M_{\\rm B}\\lsim -18$ mag (\\cite{ref:ferrA06}) host a central SMBH (e.g.~\\cite{ref:kr95}-\\cite{ref:dez01}), whose mass $\\Mbh$ scales {\\it linearly} with the spheroid stellar mass $\\Mstar$; the correlation of $\\Mbh$ with the central stellar velocity dispersion $\\sigz$ of the host galaxy is even tighter.  It is then natural to argue (e.g.~\\cite{ref:sr98}-\\cite{ref:crEA06}) that the central SMBHs have played an important role in the processes of galaxy formation and evolution, the imprint of which is represented by the Scaling Laws (hereafter SLs) mentioned above. As an additional supporting argument, several groups have noted the link between the cosmological evolution of QSOs and the formation history of galaxies (e.g.~\\cite{ref:kh00}-\\cite{ref:meEA03}, see also~\\cite{ref:pei95}-\\cite{ref:hco04}).  In addition to the SLs of their central SMBHs, early-type galaxies are also known to follow well defined empirical SLs relating their global observational properties, such as total luminosity $L$, effective radius $\\re$, and central velocity dispersion $\\sigz$.  Among others we recall the Faber \\& Jackson (\\cite{ref:fj76}, hereafter FJ), the Kormendy (\\cite{ref:korm77}), the Fundamental Plane (\\cite{ref:djd87,ref:dreA87}, hereafter FP), the color-$\\sigz$ (\\cite{ref:ble92}), and the Mg$_2$-$\\sigz$ (e.g.~\\cite{ref:buEA88}-\\cite{ref:ber03c}) relations.  Clearly, all together these scaling relations reveal the {\\it remarkable homogeneity} of early-type galaxies, provide invaluable information about their formation and evolution, and set stringent requirements that must be taken into account by any proposed galaxy-SMBH formation scenario. We recall that two alternative scenarios for the formation of Es have been proposed. In the {\\it monolithic collapse} picture, ellipticals are formed at early times by dissipative processes (e.g.,~\\cite{ref:els62}-\\cite{ref:lar75}; see also~\\cite{ref:bin77,ref:ro77}), while in the {\\it hierarchical merging scenario} spheroidal systems are the end--products of several merging processes of smaller galaxies, the last major merger taking place in relatively recent times, i.e. at $z \\lsim 1$ (e.g.~\\cite{ref:too77}-\\cite{ref:coEA00}).  Each of the two scenarios scores observational and theoretical successes and drawbacks (e.g.~\\cite{ref:jpo80,ref:ren06}). For example, the merging picture (in its {\\it dry} flavor, i.e. neglecting the role of the dissipative gas), could be supported by some observational data suggesting that a fraction of red galaxies in clusters at intermediate redshift are undergoing merging processes: these galaxies could be the progenitors of present-day early-type galaxies (e.g.~\\cite{ref:vdEA99}). At the same time, it is not clear how repeated merging events can produce a class of objects following striking scaling laws involving their global structure, dynamics, and stellar population properties, while theoretical investigations showed that the dynamical processes expected to follow the strongly dissipative phases of monolithic collapse apparently lead to systems surprisingly similar to real Es (e.g., see~\\cite{ ref:va82,ref:bs84}).  In this paper I will summarize some selected topics of research in this field and present a possible coherent scenario for the co-evolution of SMBHs and their host spheroids. Because of the enormous body of dedicated literature now available, I will focus my attention on a few selected topics. Very important arguments, such as the dynamics of binary SMBHs and their effects on the central regions of galaxies (e.g.~\\cite{ref:mer06a}-\\cite{ref:mer07} and references therein), or the growth of SMBHs via accretion of stars (e.g.~\\cite{ref:cb94} and references therein) are just mentioned here but not discussed.  The paper is organized as follows. After a first, observationally based part (Sect.~2), where the main SLs followed by early-type galaxies and by their central SMBHs are described, in Sect.~3 I will focus on the physical interpretation of the FP: this phenomenological part is a prerequisite to the construction of a possible formation scenario.  Then (Sect.~4) I discuss two important physical mechanisms in principle able to destroy the galaxy and SMBH SLs, namely {\\it galaxy merging} and {\\it cooling flows}. In this context, I present some arguments that require that we distinguish between the {\\it origin} and {\\it maintenance} of the SLs and the unavoidable occurrence of SMBH {\\it feedback} on the galaxy Interstellar Medium in the late stages (when Es are sometimes considered as ``dead, red objects''). This last point is strictly related to the solution of the long-standing problem of cooling flows in Es (and clusters), and to the quiescence of Active Galactic Nuclei (AGN) in low-redshift systems.  Section 5 addresses the problem of the {\\it origin} of the SLs, while Sect.~6 provides a summary of the main results, with a short discussion of some very recent observational findings, i.e. the surprisingly high stellar density of Es at high redshift. ",
        "conclusions": "We are finally in the position to connect the different pieces of information described in the previous Sections, to see if it is possible to form a plausible scenario in which the existence of the galaxy and central black holes Scaling Laws described in Sect. 2 can be traced back to the process of galaxy formation.  The {\\it first} important clue is that stellar spheroids are a remarkably regular family of stellar systems: the regularity is apparent in terms not only of density profiles, but also of orbit composition and of stellar populations. All these indications point towards a {\\it common} formation mechanism, where the galaxy mass has been a {\\it major} parameter, because many galaxy Scaling Laws involve the galaxy luminosity.  In Sect.~3 a review of the different proposed interpretations of the galaxy Scaling Laws has been presented. While a definite answer is not reached yet, it is generally acknowledged that the galaxy SLs are mainly due to a {\\it systematic variation with the galaxy luminosity} of the dark matter amount and distribution, of the light distribution (the so-called weak homology), and finally of the metallicity of the bulk of the stellar mass.  The {\\it second} piece of information about galaxy Scaling Laws is indirect and comes from cosmological simulations: by studying simulations of structure formation on the scale of clusters of galaxies, it is found that well defined Scaling Laws are naturally (i.e., by initial conditions) imprinted in the resulting dark matter halos.  Thus, it is tempting to put the two points above together and to suggest that the formation of stellar spheroids proceeded mainly in a way similar to the {\\it monolithic collapse} scenario. This would explain the galaxy Scaling Laws just as the imprint of the dark matter halos Scaling Laws (at the mass scale of galaxies) on the baryons. Numerical simulations of fast (dissipationless) collapse in pre-existing dark matter halos can reproduce Sersic profiles similar to those observed, from the outer parts of the models down to their central regions.  Cold dissipationless collapse is a process which is expected to dominate the late stages of an initially dissipative process, in which the gas cooling time of the forming galaxy is shorter than its dynamical time, so that stars form ``in flight'', and the subsequent dynamical evolution is dissipationless. Observational evidence supports this argument. In fact, the observed color-magnitude and Mg$_2$-$\\sigz$ relations, and the increase of the $[\\alpha/{\\rm Fe}]$ ratio with $\\sigz$ in the stellar population of Es (e.g.~\\cite{ref:ber03c},\\ref{ref:tgb99}-\\cite{ref:saEA00}), suggest that star formation in massive ellipticals was not only more efficient than in low-mass galaxies, but also that it was faster (i.e., completed before SNIa explosions took place), with the time-scales of gas consumption and ejection shorter or comparable to the galaxy dynamical time (e.g.~\\cite{ref:matt94, ref:pmat04}) and decreasing with increasing galaxy mass.  The {\\it third} piece of information is related to the effects of {\\it dry and wet merging} on the Scaling Laws: in fact we know that ellipticals cannot be originated by parabolic merging of low mass spheroids only, even in the presence of substantial gas dissipation (which, at variance with dry merging, is able to increase the galaxy central velocity dispersion, see also \\cite{ref:bnj07})).  However, it is also known that SLs such as the FJ, Kormendy, FP, and the $\\Mbh$-$\\sigz$ relations, when considered over the whole mass range spanned by ellipticals in the local Universe, are robust against merging (see also~\\cite{ref:peng07}).  Thus the galaxy Scaling Laws, possibly established at high redshift by the fast collapse in pre-existing dark matter halos of gas rich and clumpy stellar distributions (e.g.~\\cite{ref:mgsc07}), can persist even in the presence of a (small) number of dry mergers at lower redshift (\\cite{ref:mhb07}).  If this is the case, then monolithic-like collapse at early times and subsequent merging could just represent the different phases of galaxy formation (collapse) and evolution (merging, in addition to the aging of the stellar populations and related phenomena).  The possibility that monolithic collapse and successive merging events are just the leading physical processes at different times in galaxy evolution, and that they are both important for galaxy formation, is perhaps indicated also by a ``contradictory'' and often overlooked peculiarity of massive ellipticals. In fact, while the Kormendy relation dictates that the mean stellar density of galaxies decreases with increasing galaxy mass (a natural result of parabolic dry merging), the normalized light profiles of Es becomes steeper and their metallicity increases at increasing galaxy mass (as expected in case of significant gas dissipation).  Thus, the present-day light profiles of ellipticals could represent the fossil evidence of the impact of both processes.  If the above scenario is correct, then one expects that the star formation history in the Universe and the QSO activity should be roughly parallel.  It remains to be clarified if QSO activity brought the process of galaxy formation to an end, or the star formation feedback by ejecting the remaining gas from the galaxy brought the epoch of vigorous QSO activity to a conclusion, or finally if a combination of stellar and AGN feedback was the key factor.  My personal view is that we currently have more indications supporting the idea that galaxy formation was stopped more by stellar feedback (i.e., supernova heating) than by the AGN feedback (but see~\\cite{ref:scha07}). In any case, after the end of the fast star formation epoch, {\\it necessarily} a new evolutionary phase begins for the galaxies and their SMBHs. This obvious fact is curiously neglected quite often in the current literature, but it is unavoidable. In fact, over a cosmological time in a passively evolving galaxy the {\\it stellar mass losses} amount, over a cosmological time, to a considerable fraction ($\\gsim 30\\,\\%$) of the total stellar mass, i.e., $\\sim$ 2 orders of magnitude larger than the observed SMBHs masses. If only a minor fraction of this recycled gas (the basic ingredient of the galaxy ``cooling flow'' model!) were accreted on the central SMBH, the SLs involving the BH masses (such as the Magorrian and the $\\Mbh$-$\\sigz$) would be completely different. Thus, the need of an {\\it extremely efficient} feedback from SMBHs is {\\it not} required by complex physical arguments, but just by the {\\it mass budget} of the SMBHs. In addition, this feedback {\\it must} be active over the whole galaxy life, and cannot be temporally concentrated just at the end of the star formation epoch. {\\it However, a moderate accretion from the recycled gas by the evolving stars will not destroy an already established SL, such as the Magorrian relation, as the available ``fuel'' for accretion is naturally proportional to the stellar mass of the host system.}  I conclude this Review with a brief comment on a recent and very interesting observational finding, i.e. the fact that apparently stellar spheroids were much {\\it denser} than today at redshift $1\\lsim z\\lsim 2$ (e.g.~\\cite{ref:dise05,ref:truj07,ref:cima08}, and references therein). The obvious question is what mechanism could make a galaxy ``expand''. Of course, internal dynamical processes cannot be invoked, as their time-scales are measured {\\it either} by the galaxy dynamical time (very short compared to the age of the system), {\\it or} by the 2-body relaxation time (which is orders of magnitude longer than the age of the Universe). Thus, the only obvious possibility is to postulate that a few events of dry merging are common in the life of early-type galaxies.  This would also help to explain the ``central density-slope paradox'' discussed above.  In addition, if dry merging (through a small number of events) is the solution to the problem of superdense galaxies then, in order for present-day galaxies to obey the Magorrian relation, the SMBHs at the center of the superdense progenitors should also follow the same SL, because no significant amount of gas can be accreted on the center in a dry merging. Then, the superdense galaxies cannot follow the $\\Mbh$-$\\sigz$ relation observed in the local Universe because their velocity dispersion is higher than in local galaxies of the same mass. In practice, if superdense galaxies are the progenitors of the nearby Es, and if their expansion was caused by dry merging, they should obey the Magorrian relation observed in the local Universe, but they should fail the local $\\Mbh$-$\\sigz$ relation, by exhibiting a systematically lower zero-point.  Testing observationally this conjecture would be an interesting goal for the future."
    },
    "0808/0808.1441_arXiv.txt": {
        "abstract": "Several short-lived radionuclides (SLRs) were present in the early solar system, some of which should have formed just prior to or soon after the solar system formation.  Stellar nucleosynthesis has been proposed as the mechanism for production of SLRs in the solar system, but no appropriate stellar source has been found to explain the abundances of all solar system SLRs.  In this study, we propose a faint supernova with mixing and fallback as a stellar source of SLRs with mean lives of $<$5 Myr ($^{26}$Al, $^{41}$Ca, $^{53}$Mn, and $^{60}$Fe) in the solar system.  In such a supernova, the inner region of the exploding star experiences mixing, a small fraction of mixed materials is ejected, and the rest undergoes fallback onto the core.  The modeled SLR abundances agree well with their solar system abundances if mixing-fallback occurs within the C/O-burning layer.  In some cases, the initial solar system abundances of the SLRs can be reproduced within a factor of 2.  The dilution factor of supernova ejecta to the solar system materials is $\\sim$10$^{-4}$ and the time interval between the supernova explosion and the formation of oldest solid materials in the solar system is $\\sim$1 Myr.  If the dilution occurred due to spherically symmetric expansion, a faint supernova should have occurred nearby the solar system forming region in a star cluster. ",
        "introduction": "The former presence of short-lived radionuclides (SLRs) in the early solar system ($^{10}$Be, $^{26}$Al, $^{36}$Cl, $^{41}$Ca, $^{53}$Mn, $^{60}$Fe, $^{107}$Pd, $^{129}$I, and $^{182}$Hf) has been inferred from excesses in the abundances of their daughter nuclides in meteorites, linearly correlated with the abundance of a parent element \\citep[e.g.,][]{McD03,Ki05}.  These now-extinct SLRs may provide high-resolution ($<$0.1 million years (Myr)) chronometers for events that occurred during the first several million years of solar system evolution and may also be potential heat sources for asteroidal metamorphism and/or differentiation.  The SLRs with relatively long mean-lives such as $^{107}$Pd, $^{129}$I, $^{182}$Hf, and perhaps $^{53}$Mn may have been products of steady-state nucleosynthesis in the galaxy \\citep{Ja05}, while those with mean-lives ($\\tau$) of $<$5 Myr, $^{10}$Be ($\\tau$=2.2 Myr), $^{26}$Al ($\\tau$=1.03 Myr), $^{36}$Cl ($\\tau$=0.43 Myr), $^{41}$Ca ($\\tau$=0.15 Myr), $^{60}$Fe ($\\tau$=2.2 Myr), and possibly $^{53}$Mn ($\\tau$=5.3 Myr), should have been produced either by energetic-particle irradiation in the early solar system or by stellar nucleosynthesis just prior to or shortly after the birth of the solar system.  Beryllium-10, which was found in CAIs (calcium-aluminum-rich inclusions) (McKeegan et al., 2000), is not produced by stellar nucleosynthesis, but can be efficiently formed by energetic particle irradiation.  On the other hand, $^{60}$Fe, the abundance of which in the early solar system also appears to require a late addition \\citep{TH03,Mo05,Ta06}, can be efficiently produced only in stars.  Thus, the presence of $^{10}$Be and $^{60}$Fe in the early solar system suggests that both energetic particle irradiation and stellar nucleosynthesis make contributions to the inventories of solar system SLRs.  However, it is not clear yet which process contributed more significantly to the inventory of the SLRs that could be synthesized by both processes, such as $^{26}$Al, $^{36}$Cl, $^{41}$Ca, and $^{53}$Mn.  There have been several attempts to find a plausible stellar source(s) for the abundances of SLRs in the early solar system.  A low-mass (1.5 M$_{\\sun}$) thermally-pulsing asymptotic-giant-branch (TP-AGB) star cannot produce enough $^{60}$Fe to match the initial abundance of $^{60}$Fe in the solar system \\citep[e.g.,][]{Bu03,Wa06,SS06}.  Models for intermediate-mass AGB stars (5 M$_{\\sun}$ with the solar metallicity or 3 M$_{\\sun}$ with 1/3 $\\times$ the solar metallicity) could explain the inferred solar system abundance of $^{60}$Fe as well as the abundances of $^{26}$Al, $^{41}$Ca and $^{107}$Pd \\citep{Wa06}.  However, the probability of encounters between molecular clouds and AGB stars seems to be extremely low \\citep{KM94}, implying that an AGB star was an unlikely source of SLRs in the solar system.  Mass-loss winds from Wolf-Rayet stars may have been a source of the solar system $^{26}$Al, $^{36}$Cl, $^{41}$Ca, and $^{107}$Pd if the nuclides were incorporated into the solar system within a time-interval of $\\sim$10$^{5}$ year after production \\citep{Ar06}.  However, a WR wind alone would not contain enough $^{60}$Fe and $^{53}$Mn to explain their solar system initial abundances \\citep{Ar06}.  Type II core-collapse supernovae have also been considered as plausible sources for SLRs.  However, most supernova models imply that if a supernova provided $^{26}$Al and $^{41}$Ca into the solar system, it would also supply 10-100 times more $^{53}$Mn than its estimated initial abundance in the solar system \\citep[e.g.,][]{GV00,Ou05,SS06}.  This discrepancy could be explained by fallback of the innermost layers containing most of the $^{53}$Mn onto a collapsing stellar core \\citep{MC00,Me05}.  In such a case, $^{53}$Mn could primarily be derived from a different source such as the interstellar medium.  \\cite{Wa06} proposed that $^{53}$Mn and possibly $^{60}$Fe were injected into the solar system as supernova ejecta with a time-interval of $\\sim$10$^{7}$ years after production, long enough for $^{41}$Ca and $^{26}$Al to decay completely, and that $^{26}$Al and $^{41}$Ca in the solar system may have been produced either by energetic particle irradiation or by an AGB star.  Another problem with supernovae as sources of SLRs is overproduction of $^{60}$Fe if all the $^{26}$Al in the solar system was derived from supernovae.  Although the yield of $^{60}$Fe depends the mass loss and initial mass \\citep[e.g.,][]{LC06}, the expected amount of $^{60}$Fe injected from a supernova would be, in general, a few times higher than its highest estimate for the early solar system, as we will show later.  This problem may still remain even in models for fallback supernovae.  In this study, we propose a supernova with mixing and fallback, with a kinetic energy of explosion slightly less than that for a typical supernova ($\\sim$10$^{51}$ erg) as a potential source of $^{26}$Al, $^{41}$Ca, $^{53}$Mn, and $^{60}$Fe in the early solar system.  Faint supernovae such as SN1997D and SN1999br have such kinetic energies \\citep[e.g.,][]{No06}.  In models for supernovae with mixing-fallback, the inner region of the exploding star experiences mixing, some fraction of mixed materials is ejected, and the rest undergoes fallback onto the core \\citep[e.g.,][]{UN02,UN05,No06,To07}.  Nucleosynthesis in a faint supernova with mixing-fallback successfully reproduces the elemental abundance patterns of hyper metal-poor stars \\citep{UN03,Iw05,No06}. ",
        "conclusions": ""
    },
    "0808/0808.3358_arXiv.txt": {
        "abstract": "{We present large scale 870\\,\\um\\ maps { of} the nearby starburst galaxies NGC\\,253 and NGC\\,4945 as well as the nearest giant elliptical radio galaxy Centaurus A (NGC\\,5128) obtained with the newly commissioned {Large Apex Bolometer Camera (LABOCA)} operated at the { Atacama Pathfinder Experiment} telescope. Our continuum images reveal for the first time the distribution of cold dust at a angular resolution of { 20$''$} across the entire optical disks of NGC\\,253 and NGC\\,4945 out to a radial distance of $10'$ (7.5 kpc). In NGC\\,5128 our LABOCA image also shows, for the first time { at submillimeter wavelengths}, the synchrotron emission associated with the radio jet and the inner radio lobes. From an analysis of the 870\\,\\um\\ emission in conjunction with ISO-LWS, IRAS and long wavelengths radio data we find temperatures for the cold dust in the disks of all three galaxies of 17--20\\,K, comparable to the dust temperatures in the disk of the Milky Way. The total gas mass in the three galaxies is determined to be 2.1, 4.2 and $2.8\\times10^{9}\\,\\msol$ for NGC\\,253, NGC\\,4945 and NGC\\,5128, respectively. The mass of the warmer (30--40\\,K) gas associated with the central starburst regions in NGC\\,253 and NGC\\,4945 only accounts for $\\sim10\\%$ of the total gas mass. A detailed comparison between the gas masses derived from the dust continuum and the integrated CO(1--0) intensity in NGC\\,253 suggests that changes of the CO luminosity to molecular mass conversion factor are mainly driven by a metallicity gradient and only to a lesser degree by variations of the CO excitation. An analysis of the synchrotron spectrum in the northern radio lobe of NGC\\,5128 shows that the synchrotron emission from radio to the ultraviolet (UV) wavelengths is well described by a broken power law and that the break frequency is a function of the distance from the radio core as expected for aging electrons. We derive an outflow speed of $\\sim0.5$\\,{\\it c} at a distance of 2.6\\,kpc from the center, consistent with the speed derived in the vicinity of the nucleus. ",
        "introduction": "Nearly half the bolometric luminosity in the local universe is emitted at mid- and far-infrared (IR) wavelengths. The IR radiation is produced by warm interstellar dust grains that are heated by UV photons from hot high mass stars. This thermal emission of dust grains therefore carries valuable information on feedback processes from star formation, the chemical composition of the { interstellar medium (ISM)} and also on the total amount of interstellar matter in galaxies, the gas surface density and its relation to star-forming regions. The high spatial resolution of the Spitzer Space Telescope has greatly improved our view in this important { wavelength} range and allows for the first time to study spatially resolved spectral energy distributions (SEDs) across galaxies and { the relation of IR emission} to other tracers of star formation. \\\\ In order to provide a complete measurement of the dust SED it is highly desirable to combine the IR observations with data on the long wavelength (Rayleigh-Jeans) tail of the SED in the submm regime (see, e.g., Draine \\etal\\ \\cite{draine07}). This is because the submm observations are particularly sensitive to cold dust which dominates the dust mass in galaxies. \\\\ Ground based observations of the submm emission of nearby galaxies, however, remain a challenging task because they largely suffer from limitations due to the Earth's atmosphere. Furthermore, existing submm cameras have been limited in their field of view (FoV) to a few arc minutes which makes it difficult to survey large areas on the sky in a reasonable amount of time.\\\\ With the commissioning of the Large APEX Bolometer Camera (LABOCA, Siringo \\etal\\ \\cite{siringo07} and in prep.) at the APEX telescope  (G\\\"usten \\etal\\ \\cite{guesten06}) at 5100\\,m altitude on Chajnantor this situation has largely improved. With its large field of view and large number of detectors in combination with the extremely dry conditions at the site, LABOCA provides the mapping-speed and sensitivity required to survey large areas on the sky down to mJy noise levels.\\\\ In this paper we present the first large scale 870\\,\\um\\ maps { by LABOCA} towards two nearby starburst galaxies NGC\\,253 and NGC\\,4945, and towards the nearest giant elliptical radio galaxy Cen\\,A (NGC\\,5128). In Sect.\\,2 we describe the LABOCA observations and the data reduction, Sect.\\,3 focuses on the dust SEDs and gas masses. In Sect.\\,4 we discuss the thermal and non-thermal emission processes in our target galaxies and Sect.\\,5 summarizes our results. ",
        "conclusions": "\\subsection{Dust temperatures} Our dust temperature estimate for the coldest, mass carrying dust component of 17\\,K in the disk of NGC\\,253 is in excellent agreement with estimates based on lower resolution ISOPHOT and IRAS data (Melo \\etal\\ \\cite{melo02}, Radovich \\etal\\ \\cite{radovich01}). The dust SEDs in the disks of NGC\\,4945 as well as NGC\\,5128 show similarly low dust temperatures, which suggests that the temperatures of the dust in the disks of all three active galaxies are comparable to those found in the Milky Way (MW, e.g. Cox \\& Mezger \\cite{cox89}). None of our target shows evidence for dust with temperatures below 10\\,K. This finding is also in agreement with results based on SCUBA and Spitzer observations of the SINGS galaxy sample (Draine \\etal\\ \\cite{draine07}). { Our dust SEDs, however, do not rule out the presence of such cold gas. Its effect on the total dust/gas mass estimate, on the other hand, is not dramatic because a cold (e.g. a 8\\,K) dust component in the disks cannot contribute to more than $\\sim$40\\% of the observed 870\\microns\\ flux in order to be consistent with the overall SED shape.}\\\\ For the nuclear starburst regions our analysis suggests that NGC\\,253 is free of cold (T$<$30\\,K) dust while the central SEDs for NGC\\,4945 and NGC\\,5128 both show a clear signature of dust at temperatures comparable to those found in the disks. This, however, does not necessarily imply that the nuclear starburst itself contains cold gas because the 80$''$ aperture analyzed here is still too large to separate spatially the central active region from the surrounding disk. For NGC\\,5128 the situation is further complicated by the strong non-thermal contribution from the central AGN and { the jets} to the 870\\,\\um\\ flux, which makes it difficult to determine the submm flux related to thermal dust emission alone. Therefore, the suggested decrease of the cold { dust component's temperature} from the disk towards the center in NGC\\,5128 may simply reflect an imperfect separation between the thermal and non-thermal flux contributions.  \\subsection{Gas masses} Due to the large areas our target galaxies project on the sky, only few mm/submm observations have been published so far which allow to determine the total molecular gas content across the entire optical disks. { For NGC\\,253, Houghton \\etal\\ (\\cite{houghton97}) have presented CO(1--0) observations having a similar extent} than our LABOCA maps. They find a total molecular gas mass of $2.4\\times 10^{9}\\,\\msol$ using a CO conversion factor of $3\\times10^{20}\\,{\\rm cm}^{-2}\\,({\\rm K\\,\\kms})^{-1}$. This value includes the contribution of heavier elements.  The \\hi\\ mass of NGC\\,253 is $1.8\\times 10^{9}\\,\\msol$ (Koribalski \\etal\\ \\cite{kori04}) which yields a total mass of $4.2\\times 10^{9}\\,\\msol$, a factor of 2 higher than our estimate based on the dust continuum. This may argue for a smaller CO conversion factor or a higher gas-to-dust mass ratio. { Several studies have suggested that the average gas-to-dust mass ratio of a galaxy is a function of its enrichment (e.g. Draine \\etal\\ \\cite{draine07}, Engelbracht \\etal\\ \\cite{engel08}). Indeed the characteristic oxygen abundance (a measure for the average metallicity of a galaxy) in NGC\\,253 is lower by  $\\sim 0.1$\\,dex  than the MW's metallicity (Pilyugin \\etal\\ \\cite{pil04}). But such a small difference in metallicity is not expected to cause the higher gas-to-dust mass ratio in NGC\\,253 illustrated in Fig.\\,\\ref{dust_metall} where we show the location of our target galaxies on the dust-to-gas mass ratio vs. characteristic oxygen abundance plot for the SINGS galaxies (Draine \\etal\\ \\cite{draine07}, their Fig\\,16). NGC\\,253 lies well within the factor 2 scatter band on this metallicity relation observed for the SINGS galaxies. Given the simplifying assumptions made in our mass estimate (e.g. the constant dust absorption coefficient across the disk of NGC\\,253 and the use of constant CO conversion factor), the mass estimates for the dust continuum and the CO and \\hi\\ are in reasonable agreement.}\\\\    \\begin{figure} \\centering \\includegraphics[width=8.5cm,angle=0]{9909fig8.eps} \\caption{Dust-to-gas mass ratio vs. characteristic oxygen abundance plot for the SINGS galaxies adopted from Draine \\etal\\ (\\cite{draine07}, their Fig.\\,16). Our target galaxies are shown as filled squares. The CO intensities have been converted to \\hh\\ masses using a conversion factor of $2.0\\times10^{20}\\,{\\rm cm}^{-2}\\,({\\rm K\\,\\kms})^{-1}$.} \\label{dust_metall} \\end{figure}  \\noindent CO(1--0) observations covering the optical disk of NGC\\,4945 have been published by Dahlem \\etal\\ (\\cite{dahlem93}). They give an \\hh\\ mass of $2.2\\times 10^{9}\\,\\msol$ (scaled to D\\,=\\,3.8\\,Mpc and using a CO conversion factor of $2\\times10^{20}\\,{\\rm cm}^{-2}\\, ({\\rm K\\,\\kms})^{-1}$). Correcting for heavier elements and taking  the \\hi\\ mass of $1\\times 10^{9}\\,\\msol$ (Ott \\etal\\, \\cite{ott01}, scaled to D\\,=\\,3.8\\,Mpc) into account yields a total gas mass of $4\\times 10^{9}\\,\\msol$, in good agreement with our results. { Our dust mass is also in agreement with the expected metallicity dependence for MW like dust (see Fig\\,\\ref{dust_metall}). For the determination of the characteristic oxygen abundance we have used here the O/H --$M_{B}$ relation for spiral galaxies from Pilyugin \\etal\\ (\\cite{pil04}).}\\\\ \\\\ For NGC\\,5128 the situation is different as our estimated gas mass based on the dust emission exceeds those based on CO and \\hi\\ by more than a factor of 4: Eckhart \\etal\\ (\\cite{eckhart90}) derive a total molecular gas mass based on CO of $2.7\\times 10^{8}\\,\\msol$ (scaled to D\\,=\\,3.5\\,Mpc). The \\hi\\ mass accounts for $2.5\\times 10^{8}\\,\\msol$ (Gardner \\& Whiteoak \\cite{gardner76}, scaled to D\\,=\\,3.5\\,Mpc). Correcting for heavier elements this yields a total gas mass of  $6\\times 10^{8}\\,\\msol$ compared to our estimate of $2.8\\times 10^{9}\\,\\msol$. We note that this discrepancy still holds if we would assume that the cold dust towards the center has a temperature of 20\\,K instead of the 14\\,K derived  from our SED fitting (see discussion above).\\\\ Based on SCUBA 850 and 450\\,\\um\\ observations, Leeuw \\etal\\ (\\cite{leeuw02}) find a dust mass of $2.2\\times 10^{6}\\,\\msol$ for the central 4.5\\,kpc which translates into a gas mass of $3.3\\times 10^{8}\\,\\msol$ assuming a gas-to-dust mass ratio of 150 -- a factor of two lower than the estimate based on CO and a factor of 8 lower than our estimate. This discrepancy is difficult to understand because our assumed dust absorption coefficient, $\\kappa_\\nu$, differs at most by $\\sim30\\%$ from the values used by  Leeuw \\etal\\ and our LABOCA flux is only slightly higher than the SCUBA 850\\,\\um\\ flux (7.8 Jy compared to 6.4 Jy adding the inner and outer disk as defined in Leeuw \\etal). Furthermore, our dust temperature is somewhat higher than the 12\\,K derived by Leeuw \\etal, so that we would expect only a small difference between the two dust mass estimates. Using the dust temperatures, emissivities and 850\\,\\um\\ fluxes given in  Leeuw \\etal\\ we indeed estimate dust and gas masses of $1.1\\times 10^{7}\\,\\msol$ and  $1.7\\times 10^{9}\\,\\msol$ for NGC\\,5128 in better agreement with our results. Our higher dust masses are also in agreement with the results of Mirabel \\etal\\ (\\cite{mirabel99}). Their 850\\,\\um\\ fluxes, however, are about a factor of 10 higher than those published by Leeuw \\etal\\ (\\cite{leeuw02}) and from our LABOCA maps and their dust mass should therefore be treated with some caution. Independent estimates of the dust mass in NGC\\,5128 have been derived by Block \\& Sauvage (\\cite{block00}) using the V-15\\,\\um\\ ratio which results in $M_{\\rm d}=2.3\\times10^{6}\\,\\msol$ (scaled to D\\,=\\,3.5\\,Mpc). This study, however, only addresses the dust content of the central 90$''$ and may miss diffuse dust which could significantly increase the estimated dust mass. \\\\ Considering these uncertainties in the dust mass in Cen\\,A it is impossible to obtain strong conclusions by comparing the gas mass derived from CO and the dust continuum. We note, however, that the molecular gas mass derived by  Eckart \\etal\\ (\\cite{eckhart90}) is based on a standard galactic conversion factor and therefore unlikely to underestimate the gas mass by more than a factor of 2. Therefore, the high dust mass derived from our LABOCA data may suggest that the dust properties or the gas-to-dust mass ratio in NGC\\,5128 are different from those in NGC\\,253 and NGC\\,4945.\\\\ { To our knowledge no oxygen abundance estimates across the disk of NGC\\,5128 has been published which would allow secure derivation of the characteristic oxygen abundance to check for metallicity effects as a possible reason for the low gas-to-dust mass ratio. Estimates based on the globular cluster systems of NGC5128 indicate metalicities similar to or somewhat higher than their MW counterparts (see Israel \\etal\\ \\cite{israel98} and references therein). Using, as for NGC\\,4945, $M_B$ as indicator for O/H leads to a similar result (12+log(O/H)\\,=\\,--8.5) although it is not clear if this relation holds for elliptical galaxies. We therefore assume a characteristic oxygen abundance similar to the MW which places NGC5128 well above the expected metallicity dependence in Fig.\\,\\ref{dust_metall}.}  \\subsection{CO intensity vs dust column density in NGC\\,253} Among our target galaxies the CO emission in NGC\\,253 has been studied in most detail and large scale CO maps exist in the literature with similar spatial resolution than our LABOCA data. Sorai \\etal\\ (\\cite{sorai00}) imaged large parts of the optical disk in CO(1--0) at a spatial resolution of $16''$ using the  Nobeyama 45m telescope. { To investigate variations of the mass tracing capabilities of CO in comparison to the dust continuum we have compared their integrated CO intensities to the \\hh\\ column densities derived from our LABOCA maps within the central 6\\,Kpc of NGC\\,253 taking the neutral gas mass fraction measured from \\hi\\ into account. The gas mass (\\hi\\ \\& \\hh) distribution from the dust was computed using a dust temperature distribution as suggested from our dust SED analysis: 35\\,K within the central  $30''\\times 16''$, 19\\,K out to a radial distance of $40''$ (ISO aperture) and 17\\,K for emission further out and a constant gas-to-dust mass ratio of 150. This gas mass distribution was corrected for \\hi\\ using the radial dependence of molecular gas mass fraction from Sorai \\etal\\ (\\cite{sorai00}, their Fig.\\,9\\,a) and finally converted to a \\hh\\ column density map.}  \\noindent The comparison was done for { 20$''$ sized pixels after smoothing both data sets to that resolution} and is shown in Fig.\\,\\ref{xco}. From this figure there is no obvious difference between the CO conversion factor in the central starburst region and positions in the bar/disk region of NGC\\,253. A linear fit yields conversion { factors of $1.7\\times10^{20}\\,{\\rm cm}^{-2}\\,({\\rm K\\,\\kms})^{-1}$ (2.7\\,\\msol\\,(K\\,\\kms\\,pc$^2$)$^{-1}$) in both regions with errors of 3\\% and 7\\% for the central and the bar/disk region}. { The larger error for the slope outside the nuclear starburst region is mainly due to the smaller dynamic range of the data points. The result is little dependent on the neutral gas correction. For example, if we assume that all gas is in molecular form and ignore the \\hi\\ (which maximizes the conversion factor), we find the same value for the nuclear region and an increase of only 50\\% ($2.5\\times10^{20}\\,{\\rm cm}^{-2}\\,({\\rm K\\,\\kms})^{-1}$) for the bar/disk region.}  \\noindent This result is surprising at the first glance, as several detailed CO studies have suggested that the CO conversion factor is { several times} lower in the central starburst region { (e.g. Mauersberger \\etal\\ \\cite{mauers96}, Harrison \\etal\\ \\cite{harrison99} for the center of NGC\\,253)} than the galactic conversion factor, which { in turn} is expected to be a good approximation for the regions outside the starburst. { In contrast Fig.\\,\\ref{xco} hints at a constant conversion factor over the entire galaxy, including its nucleus. One must note, however, that in the strict sense, Fig.~\\ref{xco} implies a constant ratio only for the measured CO/dust masses (or fluxes) across the galaxy, while the conversion of these into total gas mass may both also depend on other characteristics. Specifically, our result is consistent with the possibility of both the gas/CO and the gas/dust ratios varying similarly with metalicities. For example, Mauersberger \\etal\\ (\\cite{mauers96}) assumed a metallicity of 3 times the solar value for their determination of the CO conversion factor in NGC\\,253. Matching this, a gas-to-dust mass ratio that scales inversely with metalicity, i.e.~50 (if a galactic value of 150 is assumed) for the nucleus of NGC\\,253, is implied by our results. The apparent constancy of the CO/dust ratio across NGC253 also suggests that changes in the gas/CO conversion are unlikely to be dominated by the variation of CO excitation, since that would be apparent in deviations from the strict linear relation, which we do not see. }   \\begin{figure*} \\centering \\includegraphics[width=16.0cm,angle=0]{9909fig9.eps} \\caption{Observed integrated CO(1--0) intensities (Sorai \\etal\\ \\cite{sorai00}) compared to the \\hh\\ gas column densities derived from our dust continuum analysis in NGC\\,253 { corrected for the \\hi\\ contribution}. Filled squares show positions towards the central starburst region, open circles positions in the bar/disk of NGC\\,253. The dashed line is a linear fit to data in the starburst region only, the dotted line a fit to all other data points. The right part of the figure is a zoom to the full diagram shown on the left.} \\label{xco} \\end{figure*}  \\subsection{Radio lobes in Cen\\,A\\label{radiolobes}}  The 870\\,\\um\\ emission associated with the radio lobes is very similar in morphology to the 1.4GHz continuum emission (Fig.\\,\\ref{cena_image}\\,right \\& Fig.\\,\\ref{cena_lobes}). As the synchrotron emission from the northern radio lobe has also been detected at 15--4.5\\,\\um\\, in the UV and at X-ray energies (Brookes \\etal\\ \\cite{brookes06}, Hardcastle, Kraft \\& Worrall \\cite{hardcastle06}) it is clear that the 870\\,\\um\\ emission is synchrotron in nature and not associated with dust driven out from the disk by the radio jet. The southern radio lobe is detected for the first time at wavelengths shorter than 3.5\\,cm. We show a high contrast signal to noise map of our LABOCA observations in Fig.\\,\\ref{cena_lobes}. From this image it is evident that the northern radio jet is also detected at 870\\,\\um. The jet, however, is located in a low level negative bowl surrounding the strong central emission. Such negative artifacts are unavoidable byproducts of the skynoise removal which results in the filtering of low spatial frequencies. As a consequence the true 870\\,\\um\\ flux of the jet is difficult to determine. The 870\\,\\um\\ flux of the northern and southern lobes are 2.5 and 0.9\\,Jy respectively.\\\\ The variation of the spectral index of the northern lobe with wavelength has recently been investigated by Hardcastle, Kraft \\& Worrall (\\cite{hardcastle06}) and Brookes \\etal\\ (\\cite{brookes06}). Based on { Spitzer/IRAC}, GALEX and Chandra data Hardcastle \\etal\\ fit the synchrotron emission by a broken power law. Their model with a break frequency at $\\sim 300\\,\\rm{GHz}$, however, underestimates the observed UV flux and also provides a poor fit to the observed slope between the IRAC bands. We have included our LABOCA fluxes in the analysis of the spectral index in the three regions in the northern lobe defined by Hardcastle \\etal\\ (see Fig.\\,\\ref{cena_lobes}). These regions cover the radio jet (inner), the southern part of the radio lobe (middle) and the northern part of the lobe (outer region). The LABOCA flux within the three regions is $40\\pm20$\\,mJy, $365\\pm40$\\,mJy and $1.0\\pm0.1$\\,Jy for the inner, middle and outer region respectively. The synchrotron spectra of the three regions are shown in Fig.\\,\\ref{power}. We fit the radio to X-ray data of all three regions  with a non-standard ($\\Delta\\alpha\\not=0.5$), broken power law. For the inner region all data can well be fit with a break frequency of  $\\rm{log}_{10}(\\nu_{\\rm b/Hz,inner})=14.5\\pm0.65$ (315\\,THz) and a change of the power index of $\\Delta\\alpha_{\\rm inner}=0.8\\pm0.3$. { Our LABOCA flux falls somewhat below the expected value from this model (68\\,mJy) which we attribute to the difficulties to determine the 870\\,\\microns\\ flux density in this particular region as described above. Because the break frequency in the inner region falls into the NIR regime this data point, however, is not critical for the fitting.} For the middle and the outer regions the break frequency shifts, as expected, to lower frequency yielding $\\rm{log}_{10}(\\nu_{\\rm b/Hz,middle})=12.5\\pm0.3$ (2.9\\,THz) and $\\rm{log}_{10}(\\nu_{\\rm b/Hz,outer})=11.7\\pm0.6$ (500\\,GHz). The change of the power index for both regions is close to that predicted by a continuous injection model (KP model) with $\\Delta\\alpha_{\\rm middle}=0.6\\pm0.1$ and $\\Delta\\alpha_{\\rm outer}=0.6\\pm0.15$. For both regions the broken power law model, however, overpredicts the X-ray flux, which indicates that the energy distribution of the electrons is more complex than a distribution yielding a simple broken power-law. A detailed modeling of the energy distribution is beyond the scope of this paper but we note, that also a single injection (SI) model with its exponential decrease at high energies does not provide a good fit to the data. From Fig.\\,\\ref{power} it can be seen that a model intermediate between the SI and KP, e.g. a broken power law with a high energy exponential decrease, presumably gives a better description of the observed fluxes between the radio and X-ray.\\\\ For the inner jet region, where the broken power law fits all observed fluxes, we can also estimate the speed of the electrons using our fit to the break frequency: assuming a magnetic field strength of 3\\,nT, the minimum energy value (Brookes \\etal\\ \\cite{brookes06}), we get an age of the electrons of $\\sim\\,17000$\\,yrs which yields a speed of 0.5\\,{\\it c} using the measured projected distance of 2.6\\,kpc from the nucleus and a jet view angle of $\\sim20^{\\circ}$ (Hardcastle \\etal\\ \\cite{hardcastle03}). This speed is consistent with the electron speed derived close to the central source. (Hardcastle \\etal\\ \\cite{hardcastle03}).   \\begin{figure} \\centering \\includegraphics[width=9.0cm,angle=0]{9909fig10.eps} \\caption{High contrast representation of the signal to noise map of our 870\\,\\um\\ observations towards NGC\\,5128 smoothed to an effective spatial resolution of 25$''$ (beam smoothed). { The negative parts of the image are highlighted by two grey contours with signal to noise ratios of $-1$ and $-2$.} This presentation also shows the faint submm emission from the northern radio jet. The white and black contours are the 1.4\\,GHz VLA map from Condon \\etal\\ (\\cite{condon96}). The { white} boxes indicate those regions which have been used to analyze the spectral index.} \\label{cena_lobes} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=8.5cm,angle=0]{9909fig11.eps} \\caption{Synchrotron spectra from radio to X-rays in three apertures in the northern radio lobe. The apertures are shown in Fig.\\,\\ref{cena_lobes} and correspond to the \"outer\" (shown as dashed line and square symbols), \"middle\" (solid, circles) and \"inner\" (dotted, triangles) region defined in Hardcastle, Kraft \\& Worrall \\cite{hardcastle06}. The filled symbols represent the 870\\,\\um\\ flux densities. For all other data see Table\\,1 in Hardcastle, Kraft \\& Worrall (\\cite{hardcastle06}). The lines show broken power law fits to the data. The presentation for the outer/inner region have been up/down scaled by 1.5 dex.} \\label{power} \\end{figure}"
    },
    "0808/0808.3391_arXiv.txt": {
        "abstract": "Correlation studies of prompt and afterglow emissions from gamma-ray bursts (GRBs) between different spectral bands has been difficult to do in the past because few bursts had comprehensive and intercomparable afterglow measurements.  In this paper\\footnote{to appear in The Astrophysical Journal, Dec 20, 2008, v. 689, no. 2} we present a large and uniform data set for correlation analysis based on bursts detected by the {\\it Swift} mission.  For the first time, short and long bursts can be analyzed and compared.  It is found for both classes that the optical, X-ray and gamma-ray emissions are linearly correlated, but with a large spread about the correlation line; stronger bursts tend to have brighter afterglows, and bursts with brighter X-ray afterglow tend to have brighter optical afterglow.  Short bursts are, on average, weaker in both prompt and afterglow emissions. No short bursts are seen with extremely low optical to X-ray ratio as occurs for ``dark'' long bursts. Although statistics are still poor for short bursts, there is no evidence yet for a subgroup of short bursts with high extinction as there is for long bursts.  Long bursts are detected in the dark category at the same fraction as for pre-{\\it Swift} bursts.  Interesting cases are discovered of long bursts that are detected in the optical, and yet have low enough optical to X-ray ratio to be classified as dark. For the prompt emission, short and long bursts have different average tracks on flux {\\it vs} fluence plots.  In {\\it Swift}, GRB detections tend to be fluence limited for short bursts and flux limited for long events. ",
        "introduction": "One of the longest enduring Gamma Ray Burst (GRB) classification schemes is based on their distributions in duration and spectral hardness.  Both quantities seem to cluster into two separate classes with the longer events (those above $\\sim2$ s; Kouveliotou et al. 1993) being predominantly softer while the shorter ones are harder. The mechanism for the origin of the GRB explosions (the central engine) appears to be quite different for the two types. % Long bursts are ascribed to the core collapse to a black hole of a massive, young, rapidly rotating star in the ``collapsar'' model (Woosley 1993; MacFadyen \\& Woosley 1999; Woosley \\& Bloom 2006) which is supported by observations such as the coincidence of SNe with well-observed nearby GRBs (Galama et al. 1998; Bloom et al. 1999; Staneck et al. 2003; Hjorth et al. 2003; Pian et al. 2006). The prevalent model for short bursts  has  them        caused by the coalescence of a binary pair of compact old stars (Lattimer \\& Schramm 1974; Paczy\\'nski 1986; Eichler et al. 1989; Mochkovitch et al. 1993; Rosswog, Ramirez-Ruiz, \\& Davies 2003; Oechslin, Janka, \\& Marek  2007) which is supported by recent observations of progenitor sites with low star formation activity (Gehrels et al. 2005; Bloom et al. 2006; Fox et al. 2005; Villasenor et al. 2005, Hjorth et al. 2005; Barthelmy et al. 2005b; Berger et al. 2005).  In both scenarios, a highly-relativistic collimated outflow of particles and radiation       occurs producing prompt gamma-ray emission from shock accelerated electrons, which evolves into a long-lasting afterglow from shock interactions with the circumburst medium (e.g., M\\'esz\\'aros \\& Rees 1997). For short bursts there are also models for the afterglow in which a radioactive wind causes emission in the first day or so (Li \\& Paczy\\'nski 1998, Kulkarni 2005).  Correlation studies of prompt and afterglow emission are crucial for understanding their production mechanisms and environmental effects.  For example, Jakobsson et al. (2004) developed a criterion for determining which GRBs are ``dark'' bursts, by comparing the relative intensity of their X-ray and optical afterglows to find what fraction of bursts have high column densities. Stratta et al. (2004) studied the X-ray and optical absorption properties of 13 GRBs studied by {\\it BeppoSax}. Roming et al. (2006) and Fynbo et al. (2007) expanded on previous work to include (long) bursts from the {\\it Swift} satellite.  A more detailed work on dark bursts using a broad-band spectral analysis is given by Rol et al. (2005, 2007).  Zhang et al. (2007) present a study comparing radiative efficiencies for short and long bursts as derived from a correlation analysis.  Using {\\it Swift} short bursts, Berger (2007) compared their X-ray afterglow to their gamma-ray prompt emission, and found that 20\\% have anomalously low X-ray to gamma ray ratios indicating very low density burst sites, possibly in globular clusters, for that subpopulation (see also Berger et al 2007). Other correlation studies have been undertaken by Salmonson \\& Galama (2002), Firmani et al. (2006), Nava et al. (2006), Butler (2007), and Nysewander, Fruchter, \\& Pe'er (2008). An early study of X-ray afterglow properties at $t=11$ hr was carried out by Piran et al. (2001).  In this study we perform correlation studies using the extensive data set from {\\it Swift}.  Sections 2 and 3 cover observations and results, respectively, while in Section 4, we discuss the implications of the results and in Section 5 the conclusions and future prospects. ",
        "conclusions": "\\subsection{Correlations \\& Short/Long Distributions}   We show in this work that correlations exist between prompt and afterglow fluxes of GRBs and between different wavelength bands in the afterglow.  The highest significance correlation is between the prompt emission gamma-ray fluence and the X-ray afterglow flux at a significance level of 99.9999996\\% for long bursts and 69\\% for short bursts. The correlation between the optical afterglow and X-ray afterglow fluxes is less significant at 99\\% significance for long bursts and only $\\sim30$\\% for short bursts (for a small sample, however).  %    It is important to note that there is a wide spread in the data for all of the correlations. The correlations are real and significant, but the fraction of the observed variations due to the correlations between the above parameters accounts for only a portion of the data spread.  The correlation can only be used to predict a flux level to within approximately an order of magnitude.  The fraction of the variation due to the correlations is given by the square of the correlation parameter, $r$, which, as shown in Table 4, varies from a few percent to 50\\%.  The rest of the data spread is due to other factors such as correlations with additional unknown parameters.  An example of an additional parameter is extinction in the optical afterglow.   Short bursts are weaker on average than long bursts in afterglow fluxes.  There is overlap with the dimmer long bursts, but the short bursts extend to lower intensities than seen for long bursts.  The average X-ray flux density at 3 keV at 11 hr for the short bursts is $<F_x({\\rm short})> = 9.6\\times10^{-3}\\mu$Jy, which is more than an order of magnitude less than the average for long bursts of    $<F_x({\\rm long})> = 0.10\\mu$Jy.       The X-ray to gamma-ray correlation in Figure 2 has a positive correlation with a slope of roughly unity. This suggests that brighter bursts have more kinetic energy  in the afterglow phase to power the afterglow. This is a manifestation of similar radiative efficiency among different bursts and between long and short GRBs. Such a point was made by Zhang et al. (2007) based on an analysis of a smaller sample of early {\\it Swift} GRBs.  Except for the bursts below the ``dark'' line, most bursts in Figure 1 are confined between lines with $\\beta_{OX} =$ 0.5 and 1.0. This is consistent with a general interpretation that the optical and X-ray emission belong to the same spectral component with an index close to 0.75. Within the standard model for emission via synchrotron radiation, for slow cooling which is generally relevant at $t = 11$ hr, one expects $\\beta_{OX} \\sim (p-1)/2$ for $\\nu_m < \\nu_O < \\nu_X < \\nu_c$, which has a typical value of 0.75 for electron distribution power law $p = 2.5$. (An equivalent statement is that for this spectrum, the predicted ratio $F_R/F_X \\approx 350$ yields a line intermediate between % $\\beta_{OX}=0.5$ and $1.0$ in % Fig. 1.) % This suggests that on average, the cooling frequency is above or not much below the X-ray band at 11 hr.    \\subsection{Dark GRBs}  Another comparison of short and long GRBs relates to dark bursts. Jakobsson et al. (2004, see also De Pasquale et al. 2003) used the simple criterion to define dark bursts as those with extremely low optical to X-ray afterglow ratio, falling below the line of optical to X-ray spectral index, $\\beta_{\\rm OX}$, equal to 0.5. It may seem counterintuitive that there can be dark bursts with optical detections and bursts not detected in the optical that are not ``dark'', but the important criterion is how optically faint the burst is relative to its X-ray flux. For the pre-{\\it Swift} sample there were 5 bursts with upper limits below the dark-burst line (restricting the Jakobsson et al. sample to include only those with upper limits fainter than $m_R=23$, or $\\sim2\\mu$Jy), compared to 24 bursts with actual measurements (not upper limits) above the line, giving a fraction of $\\sim17$\\% in the dark category.  For {\\it Swift} there are 2 bursts with upper limits (GRB 050713B and  061222A) and 3 cases with measurements (GRB 060210, 070419B and 070508) below the line compared with 34 long  bursts above the line for a fraction of $\\sim17$\\% in the dark category,     the same as the pre-{\\it Swift} sample.  The conclusion is that {\\it Swift} is sampling the same source environments as previous   instruments.   The discovery of 3 cases of dark bursts with optical detections is particularly interesting.  One possible concern with this finding is that {\\it Swift} X-ray afterglows are contaminated in many bursts by emission components not from the external shocks, e.g. X-ray flares. In such instances, the Jakobsson et al. (2004) approach to define dark bursts is no longer relevant since it assumes that the X-ray and optical emission is from the same emission component, but separated by a cooling break.  However, the X-ray lightcurves for the {\\it Swift} dark bursts are smooth around 11 hr (and beyond the end of the X-ray plateau), with no significant contamination from other components. These are real ``dark'' bursts from both an observational and physics perspective.          Correlation analyses between optical and X-ray can help answer the question of whether these two afterglow components originate from the same physical processes. It is assumed in the Jakobsson et al (2004) study that both X-ray and optical emission arise from the external forward shock. Multiwavelength observations in the {\\it Swift} era reveal puzzling chromatic features of afterglow breaks (e.g., Panaitescu et al. 2006; Liang et al. 2007, 2008) that are not consistent with the simplest forward shock model. Models invoking non-forward-shock origin of X-ray afterglows have been discussed in the literature (e.g., Genet et al. 2007; Uhm \\& Beloborodov 2007; Ghisellini et al. 2007; Shao \\& Dai 2007; Panaitescu 2008). On the other hand, analyses suggest that the X-ray data are generally consistent with the temporal index and spectral index relations (e.g., Zhang \\& M\\'esz\\'aros 2004) predicted by the forward shock models, although not in every case. (Liang et al. 2007, Willingale, et al. 2007). The optical/X-ray data of some bursts (e.g., Grupe et al. 2007; Mangano et al. 2007) are      consistent with the same forward shock model. Regardless of the exact process, the analysis presented in this paper shows that generally optical/X-ray afterglow fluxes are correlated, which suggests that they are due to the same emission process. The few cases well below the correlation line are found to be dark due to extinction in the host galaxy.   For the first time we can search for dark short bursts. No short bursts are seen that fall below the dark-burst line. It is hard to find dark GRBs using this criterion since X-ray afterglow fluxes are also low for the short bursts.  However, there are some short bursts with bright X-ray afterglow, and, to date, none of those is seen to be highly deficient in optical afterglow.  Statistics are still small with only 5 optical detections, but if the observed trend continues we will be able to conclude that short bursts do not occur in regions with extremely high extinction as occurs for some long bursts.   We are beginning to have optical detections of bursts below the dark burst line. In one of the three dark bursts with detections (GRB 060210), the burst is found to have high extinction associated with its host galaxy, explaining the low optical flux (Curran et al. 2007b). By modeling the differences between $\\beta_{\\rm opt}$, $\\beta_X$, and $\\beta_{OX}$, and taking into account the Lyman$-\\alpha$ absorption ($z=3.91$), the authors find the $R-$band source extinction could be $3.9\\pm0.7$ mag ($\\nu_c > \\nu_O$) or $6.7\\pm0.6$ mag   ($\\nu_c < \\nu_O$). This is an important development in our understanding of dark GRBs. (For two of the three dark bursts with detections - 070419B and 070508 - there has not yet been sufficiently detailed follow-up work on the putative hosts for constraints to be placed on the host extinction.) Assuming that the dark bursts can be largely explained by extinction, then the optical - X-ray correlations, ignoring the dark bursts, would hold true. We note that new studies are being done to examine dark burst definitions (van der Horst et al. 2008).       \\subsection{Prompt Fluence and Flux Comparisons}   The comparison of fluences and peak fluxes in the prompt emission as shown in Figure 3 is a different kind of study than in the other two above.  In this case, the strong observed correlation and high degree of separation of short and long bursts is expected; brighter bursts with higher peak fluxes naturally have higher fluences and short bursts tend to have lower fluence for a given flux by the very fact of their short duration.  Within the short and long classes, the spread in fluence that is seen for a given peak flux is due to the diversity of durations and spectral indices.  Bursts with longer duration and hard spectra have higher fluences for the same peak flux.  It is interesting to note in Figure 3 that the short bursts tend to be fluence limited in the BAT, while long bursts tend to be peak flux limited.  This is due to the way BAT operates.  A valid GRB trigger requires a statistically significant excess in both the rate and image domains (Fenimore et al. 2004).  The ability to form an image depends on the number of photons collected on various trigger timescales, which is related to the burst fluence.  Even for relatively high peak-fluxes, short bursts can have low fluence  values and be limited in the number of photons available for the image trigger. On the other hand, long bursts tend to have higher fluences for a given peak-flux and become rate limited before the image limit is reached.  BAT also has a pure-image mode for triggering where very long duration GRBs and other transients are found by comparing sky images instead of having a rate trigger.  The lowest long-burst point in Figure 3 at a peak flux of $\\sim0.1$ was such an image-mode trigger for the very long ($T90 = 35$ min) and weak GRB 060218. A caveat on the above discussion is that the BAT trigger algorithm is complex with $\\sim500$ different trigger criteria evaluated. There are many different thresholds and limits coming into play for short and long burst triggering, with some mix of flux and fluence limits for both types.  This study was based on a 1 s binning for the gamma-ray fluxes. We have also investigated the effect of using a smaller bin size of 64 ms.  Smaller bins pick out larger peak flux values when there is short time structure or when the burst has a duration shorter than the bin size.  The effect of the smaller bin size is to shift the short bursts to the right (higher peak flux) relative to the long bursts by about a factor of 5. The larger bin size that we use allows for better statistics and is more reliable for long bursts. In either case, the short bursts tend to cluster toward lower fluences than long bursts."
    },
    "0808/0808.0671_arXiv.txt": {
        "abstract": "It has recently been pointed out that if the surface tension of quark matter is low enough, the surface of a strange star will be a crust consisting of a crystal of charged strangelets in a neutralizing background of electrons. This affects the behavior of the surface, and must be taken into account in efforts to observationally rule out strange stars. We calculate the thickness of this ``mixed phase'' crust, taking into account the effects of surface tension and Debye screening of electric charge. Our calculation uses a generic parametrization of the equation of state of quark matter. For a reasonable range of quark matter equations of state, and surface tension of order a few MeV/fm${}^2$, we find that the preferred crystal structure always involves spherical strangelets, not rods or slabs of quark matter. We find that for a star of radius 10 km and mass $1.5 M_\\odot$, the strangelet-crystal crust can be from zero to hundreds of meters thick, the thickness being greater when the strange quark is heavier, and the surface tension is smaller. For smaller quark stars the crust will be even thicker. ",
        "introduction": "\\label{sec:intro}  Quarks in their most familiar form are confined in protons and neutrons that make up standard nuclear matter. However, according to the ``strange matter hypothesis'' \\cite{Bodmer:1971we,Witten:1984rs} this form of matter might be metastable and the fully stable state would then be ``strange matter'', which contains roughly equal numbers of up, down, and strange quarks. Large (kilometer-sized) pieces of strange matter are ``strange stars'' (for a review see Ref.~\\cite{Weber:2004kj}); small nuggets of strange matter are ``strangelets'' \\cite{Farhi:1984qu}. The strange matter hypothesis remains a fascinating but unproven conjecture. In this paper we will assume that it is correct, and investigate the structure of the crust of a strange star.  The traditional picture of the surface of a strange star is a sharp interface, of thickness $\\sim$ 1 Fermi. Below the interface lies quark matter, the top layer of which is positively charged. Above the interface is a cloud of electrons, sustained by an electric field which could also support a thin nuclear matter crust in suspension above the quark matter~\\cite{Alcock:1986hz,Stejner:2005mw}, as long as the strange star is not too hot~\\cite{Usov:1997eg}.  However, if the surface tension $\\si$ of the interface between quark matter and the vacuum is small enough, the surface will take on a much more complicated structure. If $\\si$ is less than a critical value $\\sicrit$ then large lumps of strange matter are unstable against fission into smaller pieces \\cite{Jaikumar:2005ne,Alford:2006bx}. As a result, the simple surface described in the previous paragraph is unstable, and is replaced by a mixed phase involving nuggets of positively-charged strange matter in a neutralizing background of electrons. It is reasonable to guess that the ground state is a regular lattice, leading to a crust with a crystalline structure. (Note that this crystal is completely different from the Larkin-Ovchinnikov-Fulde-Ferrell phase of quark matter \\cite{Alford:2000ze,Casalbuoni:2003wh}, where the quark matter density is uniform, but the pairing gap varies in space.) Jaikumar, Reddy, and Steiner \\cite{Jaikumar:2005ne} conjecture that the strangelet crystal crust will actually be a multi-layer structure of mixed phases, analogous to the ``nuclear pasta'' phases that occur in models of the inner crust of a conventional neutron star \\cite{Ravenhall:1983uh}. At the outer edge of the crust, we expect a dilute low-pressure lattice of small strangelets in a degenerate gas of electrons. As we descend in to the star, the pressure rises, and the structure is modified (becoming denser, and perhaps changing to rods, slabs, cavities, etc). At a critical pressure $\\pcrit$ the mixed phase is no longer stable, and there is a transition to uniform neutral quark matter. As one burrows deeper into the star the pressure continues to rise, and there may be other phase transitions between different phases of quark matter \\cite{Alford:2007xm}, but those will not concern us here. Note that in this scenario the strange star depends on gravity for its existence. In the absence of gravity, it would undergo fission into strangelets.  Ref.~\\cite{Jaikumar:2005ne}, assuming zero surface tension and neglecting Debye screening, estimated that the mixed phase crust might be $40-100$~m thick, with $\\pcrit\\approx 1000~\\MeV^4$. This is an interesting result because if a strange star has a sufficiently thick crystalline crust, it might be hard to distinguish from the crust of a neutron star. Astrophysical properties that are sensitive to the crust include cooling behavior, neutrino and photon opacity during a supernova, the photon emission spectrum, glitches, and frequencies of seismic vibrations which are observed after giant flares in magnetars. For further discussion and references see Sec.~\\ref{sec:discussion}. This paper will make a more careful calculation of the properties of a strangelet crystal crust, including the effects of Debye screening \\cite{Heiselberg:1993dc} and surface tension.  We expect that the properties of the crust will emerge from a competition between various different contributions to the energy. Charge separation is often favored by the internal energy of the phases involved, because a neutral phase is always a maximum of the free energy with respect to the electrostatic potential (see \\cite{Ravenhall:1983uh,Glendenning:1992vb}; for a pedagogical discussion see \\cite{Alford:2004hz}). The domain structure is determined by surface tension (which favors large domains) and electric field energy (which favors small domains). Debye screening is important because it redistributes the electric charge, concentrating it in the outer part of the quark matter domains and the inner part of the surrounding vacuum, and thereby modifying the internal energy and electrostatic energy contributions.  To make an estimate of the thickness of the crust we need to calculate the equation of state of the mixed phase, i.e.~the energy density $\\ep_{\\rm mp}$ as a function of the pressure $p_{\\rm mp}$. The thickness of the crust for a star of mass $M$ is then \\beq \\Delta R = \\Rstar\\Bigl(\\frac{\\Rstar}{GM}-2\\Bigr)\\int_0^{\\pcrit} \\frac{1}{\\ep_{\\rm mp}} dp_{\\rm mp} \\ , \\label{dr} \\eeq in $\\hbar=c=1$ units. This expression follows from the Tolman Oppenheimer Volkoff equation \\cite{Tolman:1939jz,Oppenheimer:1939ne}, assuming that $\\De R\\ll \\Rstar$, and that everywhere in the crust the pressure is much smaller than both the local energy density and the average energy density of the whole star. These are very good approximations for the cases that we study.  We obtain $\\ep_{\\rm mp}$ as a function of $p_{\\rm mp}$ by dividing the strangelet lattice into unit cells (``Wigner-Seitz cells'') and calculating the pressure at the edge of a cell as a function of its energy density.  We study cells that are three-dimensional (a lattice of strangelets in a degenerate gas of electrons), two-dimensional (rods of strange matter in a degenerate gas of electrons) and one-dimensional (slabs of strange matter interleaved with regions of degenerate electron gas). Our approach is similar to that used in studying mixed phases of quark matter and nuclear matter in the interior of neutron stars \\cite{Maruyama:2007ey}.  We build on the formalism for a generic quark matter equation of state and infinitely-large Wigner-Seitz cells that was developed in \\cite{Jaikumar:2005ne,Alford:2006bx}. The main assumptions that we make are:\\\\ 1) Within each Wigner-Seitz cell we use a Thomas-Fermi approach, solving the Poisson equation to obtain the charge distribution, energy density, and pressure. This is incorrect for very small strangelets, where the energy level structure of the quarks becomes important \\cite{Madsen:1994vp,Amore:2001uf}; such corrections may be relevant for the very low pressure (outer crust) part of our results (see Sec.~\\ref{sec:discussion}). \\\\ 2) We assume our $D$-dimensional Wigner-Seitz cells to be $D$-spheres. In reality the cells will be unit cells of some regular lattice (cubic, hexagonal close packed, etc). However, as long as the cell is much bigger than the strangelet inside it, we expect this approximation to be reasonably accurate. We will only report results for cases where $\\Rcell>2R$ ($R$ being the radius of the quark matter in the center of the cell, which we expect will have a rotationally symmetric shape because of the surface tension). In some cases this assumption is violated, and we will then only be able to obtain a lower limit on the crust thickness (see Sec.~\\ref{sec:thickness-results}). \\\\ 3) We treat the interface between quark matter and the vacuum as a sharp interface, with no charge localized on it, which is characterized by a surface tension. We neglect any surface charge that might arise from the reduction of the density of states of strange quarks at the surface \\cite{Madsen:2000kb,Madsen:2001fu,Madsen:2008bx,Oertel:2008wr}. We also neglect the curvature energy of a quark matter surface \\cite{Christiansen:1997rc,Christiansen:1997vt}, so we do not allow for ``Swiss-cheese'' mixed phases, in which the outer part of the Wigner-Seitz cell is filled with quark matter, with a cylindrical or spherical cavity in the center, for which the curvature energy is crucial. Note that these phases would be expected to occur at higher pressure than the ones we study, so including them is likely to make the crust even thicker than we predict. \\\\ 4) We assume that the chemical potential for negative electric charge $\\mue$ is much less than the chemical potential for quark number $\\mu$. This allows us to expand the quark matter equation of state in powers of $\\mue$, and means that within the quark matter we can ignore the contribution of electrons to the charge and pressure. This is a very good approximation for small strange quark mass, which corresponds to small $\\nQ$ in our parameterization. For the largest value of $\\nQ$ that we study, $\\mue$ in neutral quark matter is close to 100 MeV, and the assumption is still reasonable.\\\\ 5) We assume that only electrons are present, with no muons. This is valid as long as $\\mue$ is less than the muon mass $m_\\mu$, which is true for all the cases that we study.\\\\ 6) We assume that $\\mue$ is always much greater than the electron mass. Thus in the degenerate electron gas, we can take the electrons to be massless, which simplifies the Thomas-Fermi calculation of their charge distribution. Since $\\mue$ drops monotonically from the center of the cell to its edge, this condition will only be violated for very large cells (very low pressures).\\\\ 8) We always work at zero temperature. The temperature of the surface of a compact star, even during a flare \\cite{Lyubarsky:2002cs}, is expected to be less than 100 keV, so we expect this to be a reasonable approximation. ",
        "conclusions": "\\label{sec:discussion}  The calculations described in this paper give us a more precise picture of the strangelet-crystal crust of a quark star. The results presented in Table~\\ref{tab:crusts} show that the thickness of the strangelet-crystal crust of a strange star is very sensitive to the surface tension $\\si$ of the interface between quark matter and the vacuum, and to the quark matter parameters  $\\nQ$ and $\\chiQ$ \\eqn{eqn:generic_EoS}, which determine the response of the quark matter to deviations of the electrostatic potential from its neutrality value. Our results are compatible with those of Ref.~\\cite{Jaikumar:2005ne} where an upper limit on the crust thickness was obtained by ignoring surface tension and Debye screening. The crust is thickest for large $\\nQ$ and small $\\chiQ$ (large $\\la_D$). As discussed in Sec.~\\ref{sec:thickness-results}, we find that values of surface tension in the physically expected range, around 1 to $10~\\MeV\\fm^{-2}$, reduce the thickness of the crust and may even eliminate it completely, but it remains possible that a quark star of radius 10~km could have a crust several hundred meters thick. From Table~\\ref{tab:factors} we see that for a smaller star the crust could be even thicker.  The geometry of the mixed phase in our crusts, on the other hand, shows no variation at all. It is always three-dimensional, containing spherical droplets of quark matter. We never find any case where a two-dimensional (rod) or one-dimensional (slab) geometry is favored.  Our calculations and results suggest two directions for future work: firstly, one could study phenomenological consequences of our understanding of the strangelet-crystal crust of a quark star. Secondly, one could improve on our treatment of the strangelet crystal, by relaxing some of the assumptions listed in Sec.~\\ref{sec:intro}.  The most obvious phenomenological task is to revisit computations of the frequencies of seismic vibrations which are observed after giant flares in magnetars \\cite{Watts:2006hk,Chugunov:2006kk}. Ref.~\\cite{Watts:2006hk} found that the strangelet crystal crust did not have the right spectrum of toroidal shear modes to account for current observations: it would be interesting to see whether taking in to account the surface tension and Debye screening affects that conclusion. Other aspects of the phenomenology of the crust could also be studied, for example (a) the thermal response of the crust to accretion \\cite{2000ApJ...531..988B}; (b) the role of the crust in the trapping of neutrinos and photons just after a type-II supernova \\cite{Burrows:2004vq}; (c) the spectrum of photons radiated from the surface of a quark star \\cite{Page:2002bj,Jaikumar:2004zy,Harko:2004ts}; (d) the contribution of the crust to the moment of inertia and glitches \\cite{1992ApJ...400..647G}; (e) the damping of $r$-modes in by shear viscosity in the crust \\cite{2000ApJ...529L..33B,Caballero:2008tx} (for quark stars, the contribution from the interior has been calculated \\cite{Madsen:1999ci,Jaikumar:2008kh}); (f) the thermal relaxation time of the crust and its response to the post-supernova ``cooling wave'' \\cite{Gnedin:2000me}. The thermal relaxation time of the crust depends on the thermal conductivity, for which we can make a very rough estimate using appendix A of Ref.~\\cite{Gnedin:2000me}. We find values of order a few hundred $\\MeV^2$ at $T\\sim 0.1~\\MeV$, which is comparable to the range $10^{18}~{\\rm erg}\\,{\\rm cm}^{-1}{\\rm s}^{-1}{\\rm K}^{-1}$ for low-density nuclear matter (Ref.~\\cite{Gnedin:2000me}, Fig.~4). We defer a more accurate calculation to future work.  To improve on our treatment, the most pressing issues are to use a realistic shape for the Wigner-Seitz cells (which should be unit cells of some regular lattice, rather than spheres), to include shell-model corrections for the smallest strangelets, and to allow for ``Swiss-cheese'' phases where most of the unit cell consists of quark matter, with a hole at the center containing electrons.  As discussed in Sec.~\\ref{sec:thickness-results}, the shape of the cell becomes important at very high pressures, and our use of the spherical approximation meant that in some cases we could only obtain lower limits on the crust thickness. Studying more realistic shapes is straightforward in principle, but would require a more demanding multidimensional numerical solution of the Poisson equation.  Shell-model corrections can be of order one MeV per quark for strangelets of size $R\\lesssim 5~\\fm$ \\cite{Madsen:1994vp,Amore:2001uf}, which is not negligible relative to our typical enthalpy per quark (Fig.~\\ref{fig:stability-search}), and may therefore affect our results for the outer part of the crust, where we predict strangelets as small as 3~fm (Fig.~\\ref{fig:crust}).  Treating Swiss-cheese phases would require us to include curvature energy as well as surface tension. This highlights the fact that we treated the interface between quark matter and vacuum in a very simplified way, as a zero-width interface with a surface tension. However, since the quark confinement distance is about 1~fm, the interface might well have structure on this distance scale. Like the shell-model effects described above, this could be relevant to the low-pressure regime, where the strangelets can be as small as a few fm. There are even indications that when such physics is taken into account, the CFL phase may undergo some degree of charge separation \\cite{Madsen:2001fu,Oertel:2008wr}, raising the possibility that there might be some sort of crystalline crust on quark stars made of CFL quark matter."
    },
    "0808/0808.2168_arXiv.txt": {
        "abstract": "We report the discovery of six infrared stellar-wind bowshocks in the Galactic massive star formation regions M17 and RCW49 from {\\it Spitzer} GLIMPSE (Galactic Legacy Infrared Mid-Plane Survey Extraordinaire) images. The InfraRed Array Camera (IRAC) on the {\\it Spitzer Space Telescope} clearly resolves the arc-shaped emission produced by the bowshocks. We combine {\\it Two Micron All-Sky Survey (2MASS),} {\\it Spitzer,} {\\it MSX,} and {\\it IRAS} observations to obtain the spectral energy distributions (SEDs) of the bowshocks and their individual driving stars. We use the stellar SEDs to estimate the spectral types of the three newly-identified O stars in RCW49 and one previously undiscovered O star in M17. One of the bowshocks in RCW49 reveals the presence of a large-scale flow of gas escaping the \\hii\\ region at a few $10^2$ \\kms. Radiation-transfer modeling of the steep rise in the SED of this bowshock toward longer mid-infrared wavelengths indicates that the emission is coming principally from dust heated by the star driving the shock. The other 5 bowshocks occur where the stellar winds of O stars sweep up dust in the expanding \\hii\\ regions. ",
        "introduction": "The Solar wind ends in a termination shock \\citep[e.g.][]{V105}, where the pressure of the heliosphere balances the ram pressure of the surrounding interstellar medium (ISM). Massive stars with more energetic winds generate much stronger shocks. In cases where the relative motion between the star driving the wind and the ambient ISM is large, the shock will be bent back around the star. If the relative velocity is supersonic, the ambient ISM gas is swept into a second shock, forming an arc-shaped ``bowshock'' that is separated from the termination shock by a contact discontinuity. Stellar-wind bowshocks have been reported for a variety of sources, including nearby runaway O stars \\citep{vBu88,vBu95,Nor97,BB05,CP07,FML07}, high-mass X-ray binaries \\citep{EC92,K97,HK02}, LL Ori-type stars \\citep{Bally00}, radio pulsars \\citep{GS06}, Galactic center O stars \\citep{geb04,geb06}, and mass-losing red giants \\citep{Martin07}. Recently, an infrared (IR) bowshock has been observed around the young A-type star $\\delta$ Vel \\citep{AG08}. Cometary \\hii\\ regions also resemble bowshocks, due either to density gradients in the ambient gas or to motion of the ionizing source with respect to the interstellar surroundings \\citep{vB90,AH06}. Both the direction of a bowshock and its ``standoff distance'' from the star generating the wind are determined by the velocity of the star with respect to the surrounding medium. In the case of runaway O stars, this is dominated by the high space motion of the star.  We report the detection of three mid-IR bowshocks in each of two massive star formation regions: M17 and RCW49. Two of the bowshocks in M17 are around known O stars. We will demonstrate that the other bowshocks are also around likely O stars. Since these stars are in or near expanding \\hii\\ regions, we find that the direction of the bowshock is determined principally by the flow of the ISM rather than the space motion of the star. ",
        "conclusions": "We have observed 6 prominent IR bowshocks in M17 and RCW49. These objects appear to be produced by the winds of individual O stars colliding with large-scale interstellar gas flows in the \\hii\\ regions. One bowshock,  M17-S3, may be the leading edge of an evaporating globule containing a newly-formed and previously undiscovered O star in the well-studied M17 region. All three bowshocks associated with RCW49 lead us to identify new candidate O stars. Our stellar classifications also suggest that the true distance to RCW49 is less than the kinematic distance of 6 kpc.  The bowshocks are bright at IR wavelengths due to emission from dust swept up from the ambient ISM and heated by radiation from the bowshock driving stars. As \\citet{AG08} note, IR excess emission from a bowshock could be attributed to the presence of a circumstellar disk, particularly when the bowshock morphology is not spatially well-resolved. This can be a pitfall for observational studies of accreting massive protostars.   The collective winds of the most luminous stars in young, massive clusters produce overlapping large-scale flows that hollow out thermally hot cavities in the parent molecular cloud \\citep{T03}. The largest bowshock presented here,  RCW49-S1, is evidence that the combined winds of the ionizing stars in Westerlund 2 have escaped the \\hii\\ region, creating a flow of hot gas moving at a few $10^2$ \\kms\\ that extends at least 16 pc away from RCW49.  The driving stars of the other 5 bowshocks are surrounded by ionized gas and dust of their natal \\hii\\ regions, where the density of the ambient medium ($n_0\\sim 10^3$ cm$^{-3}$) is sufficiently high to produce the observed bowshocks with a relative velocity of only 10--20 \\kms. The winds of the bowshock driving stars do not directly encounter the ${>}2000$ \\kms\\ winds from the most massive stars in the cluster. Instead, the bowshocks are shaped by the expansion of the ionized gas in the \\hii\\ regions relative to the orbital motions of the stars.  Eventually, supernova explosions will produce high velocity shock waves that heat and disperse the original gas cloud. In star forming regions like M17 and RCW49 that have not yet been disrupted by supernovae, IR bowshocks serve as interstellar ``weather vanes,'' indicating the speed and direction of large-scale gas flows at points within and around giant \\hii\\ regions."
    },
    "0808/0808.1111_arXiv.txt": {
        "abstract": "{The exploitation of clusters of galaxies as cosmological probes relies on accurate measurements of their total gravitating mass. X-ray observations provide a powerful means of probing the total mass distribution in galaxy clusters, but might be affected by observational biases and rely on simplistic assumptions originating from our limited understanding of the intracluster medium physics.} {This paper is aimed at elucidating the reliability of X-ray total mass estimates in clusters of galaxies by properly disentangling various biases of both observational and physical origin.} {We use N-body/SPH simulation of a large sample of $\\sim 100$ galaxy clusters and investigate total mass biases by comparing the mass reconstructed adopting an observational-like approach with the true mass in the simulations. X-ray surface brightness and temperature profiles extracted from the simulations are fitted with different models and adopting different radial fitting ranges in order to investigate modeling and extrapolation biases. Different theoretical definitions of gas temperature are used to investigate the effect of spectroscopic temperatures and a power ratio analysis of the surface brightness maps allows us to assess the dependence of the mass bias on cluster dynamical state. Moreover, we perform a study on the reliability of  hydrostatic and hydrodynamical equilibrium mass estimates using the full three-dimensional information in the simulation.} {A model with a low degree of sophistication such as the polytropic $\\beta$-model can introduce, in comparison with a more adequate model, an additional mass underestimate of the order of $\\sim 10 \\%$ at $r_{\\mathrm{500}}$ and $\\sim 15 \\%$ at $r_{\\mathrm{200}}$. Underestimates due to extrapolation alone are at most of the order of $\\sim 10 \\%$ on average, but can be as large as  $\\sim 50 \\%$ for individual objects. Masses are on average biased lower for disturbed clusters than for relaxed ones and the scatter of the bias rapidly increases with increasingly disturbed dynamical state. The bias originating from spectroscopic temperatures alone is of the order of $10 \\%$ at all radii for the whole numerical sample, but strongly depends on both dynamical state and cluster mass. From the full three dimensional information in the simulations we find that the hydrostatic equilibrium assumption yields masses underestimated by $\\sim 10-15 \\%$ and that masses computed by means of the hydrodynamical estimator are unbiased. Finally, we show that there is excellent agreement between our findings, results from similar analyses based on both Eulerian and Lagrangian simulations, and recent observational work based on the comparison between X-ray and gravitational lensing mass estimates.} {} ",
        "introduction": "\\label{intro.sec} Galaxy clusters are the largest virialized structure known in the universe. According to the hierarchical clustering model of structure formation they form by the gravitational collapse of the rare peaks of the primeval density field, on scales of the order of $\\sim 10$ Mpc. Within this scenario their formation and evolution is a sensitive function of the cosmological matter density parameter $\\Omega_{\\mathrm{m}}$ and the mass fluctuation amplitude $\\sigma_\\mathrm{8}$, where $\\sigma_\\mathrm{8}$ is the rms linear fluctuation on scales $8 \\, h^{\\mathrm{-1}}$ Mpc.  Measurements of their evolution rate can be used to asses the growth in mass of such structures, thereby providing a powerful method to constrain the geometry and matter content of the universe \\citep[see][and references therein]{vo05}. Moreover, because of the spatial extent of the collapse scale, the cluster baryonic fraction $f_{\\mathrm{b}}$ is expected to be close to the cosmic value $\\Omega_{\\mathrm{b}}/ \\Omega_{\\mathrm{m}}$. Measurements of $f_{\\mathrm{b}}$ at high redshifts  can be used to derive constraints on the equation of state of the dark energy \\citep{hai01,maj04,al08}.  The importance of cluster of galaxies as cosmological probes will be further strengthened with the upcoming high redshift surveys. This is of particular relevance in the new era of precision cosmology, in which studies of cluster evolution will provide independent tests with which to compare constraints on cosmological models extracted from observations of the cosmic background radiation \\citep[e.g.][]{spe07} and distance measurements of high redshift supernovae \\citep[e.g.][]{ton03,rie04,rie07}.  From the scenario here outlined it follows that in order to use cluster of galaxies as cosmological probes it is crucial to accurately measure their baryonic and total gravitating mass. The methods used to derive cluster masses are  mainly based on the velocity dispersion of the optical galaxy populations \\citep[e.g.][]{biv03,rin03}, observations of the X-ray emitting intracluster medium (ICM) \\citep[e.g.][]{fin01,rei02,et02,arn05,vikhlinin2006} and on gravitational lensing \\citep[e.g.][]{smi02,mahdavi08}.  Accurate mass estimates derived from X-ray data are based on the assumptions that both the total potential and the ICM distribution are spherically symmetric and that the ICM is in hydrostatic equilibrium in the cluster potential well. The latter assumption is usually justified by the fact that the estimated ICM sound crossing times are short when compared against cluster ages \\citep{sar86}.  Under these assumptions the ICM is a faithful tracer of the underlying matter distribution and the total mass profile can be determined from the gas radial density and temperature profiles. The density profile is recovered from deprojecting of the surface brightness profile, as measured from X-ray flux maps, whereas knowledge of the temperature profile requires the availability of spatially resolved spectroscopy. High quality data taken by the present generation of X-ray telescopes, {\\it Chandra} and {\\it XMM-Newton}, allows for nearby clusters accurate measurements of these quantities out to a large fraction of the cluster virial radius \\citep{ma98,vi05,pi05,san06,pra07,leccardi08}.  The reliability of cluster mass estimated through the X-ray method can be accurately studied using N-body/hydrodynamical simulations. The principal benefit over analytical methods is that the gas evolution can be treated self-consistently. The validity of the numerical approach is supported by X-ray observations, which show the existence of complex thermal structures and of merging activity \\citep[see][for a review]{ma07}.  Since early pioneering studies, hydrodynamical simulations have become a widely used tool to investigate cluster formation and evolution in different cosmological scenarios \\citep[cf.][]{vo05}. The numerical resolution of the simulations and the modelization of the cluster gas physics  has been  improved over the years. The latter now incorporates radiative cooling \\citep{yos00,lew20,per20,mua02,dav02}, metal enrichment of the ICM by supernovae and energy feedback \\citep{kay03,tor03,valda03,bor04,kay04,kra05}.   The accuracy of cluster X-ray mass estimators has been tested by means of N-body/hydrodynamical simulations in a variety of papers \\citep{evr90,evr96,kay04,rasia06,kay07,nagai07}. \\cite{evr90} first pointed out the existence of a significant bias in the binding mass estimates when using the isothermal $\\beta-$model. \\cite{evr96} confirmed that the source of this discrepancy is related to the isothermal and hydrostatic assumptions. The lack of validity of the hydrostatic assumption is observationally motivated by optical and X-ray maps, which show the existence of substructure with an ongoing merger activity, and is numerically supported by a number of authors \\citep[e.g.][]{kay04,ras04,nagai07}, who found that in the simulations the ICM is not perfectly in hydrostatic equilibrium. This implies the presence of residual gas bulk (laminar) and turbulent (random kinetic) motions, and leads to an underestimate of the masses because of additional non-thermal pressure support which is not accounted for by the  hydrostatic equilibrium equation.  With respect the isothermal $\\beta-$model the modelization of the ICM has been significantly improved \\citep[e.g.][]{vikhlinin2006} with observational progresses, which showed the existence of temperature profiles declining with radius \\citep{de02,vi05,pi05,san06,pra07}. These features are well reproduced out of the core radii in hydrodynamical simulations which incorporate cooling and feedback \\citep{mua02,kay03,tor03,valda03,bor04}.  In order to properly assess the reliability of cluster X-ray mass estimators it is however necessary to construct mock observations of simulated clusters which must reproduce spectroscopic measurements as expected from X-ray telescopes. This is motivated by the presence of complex thermal structures in the ICM, which bias the (measured) spectral fit temperatures towards lower values than the average emission weighted cluster temperatures defined theoretically \\citep{ma01,mazzotta04,valda06}. The dependence of X-ray mass estimators on spectral biases and other systematics has been investigated through hydrodynamical simulations by a number of authors \\citep{rasia06,kay07,nagai07,je07}. \\cite{nagai07} argued that mass estimates are biased low ($\\sim 5-20 \\%$) even for clusters identified as relaxed.  In order to properly disentangle observational biases which arise from spectroscopic measurements from those due to incomplete relaxation of the gas or from the ones caused by an inaccurate modeling of the radial profiles, it is however necessary to derive X-ray mass estimates from a large sample of simulated clusters. This is the main goal of this paper, in which we apply different X-ray mass estimators to a large set of high-resolution hydrodynamical simulations of galaxy clusters. The physics of the gas includes radiative cooling, star formation, chemical enrichment and energy feedback. The sample comprises $\\sim 100$ clusters, the size of the sample being a critical quantity in order to extract sub-samples large enough to give a meaningful statistics.  We discuss the dependence of the mass bias at different radii upon the adopted analytical models and the chosen radial range used to perform the fits of the profiles,  the cluster dynamical state as well as the impact on the mass bias which follows from the use of spectral temperatures. As a statistical indicator to quantify the cluster dynamical state through the analysis of X-ray maps we adopt the power ratio method \\citep[see][and references therein]{je05}. We also investigate the limit of applicability of the dynamical equilibrium equation when used to recover the cluster true masses in the presence of significant non-thermal gas pressure.  The paper is organized as follows. In Sect. \\ref{simulations.sec} we present the procedure for simulating the cluster sample. In Sect. \\ref{observables.sec} we describe how we generate and analyze the synthetic X-ray observations that are used in Sect. \\ref{mass.sec} to recover the total mass distribution. In Sect. \\ref{masssim.sec} we investigate the reliability of the hydrostatic and hydrodynamical mass estimators from the full three-dimensional information provided by the simulations. Finally, we discuss our main results and present our conclusions in Sect. \\ref{conclusions.sec} . ",
        "conclusions": "\\label{conclusions.sec} This work is aimed at elucidating the reliability of X-ray total mass estimates in clusters of galaxies using N-body/SPH simulation of a large sample of clusters. The physical modeling the gas includes radiative cooling, star formation, supernovae heating, and metal enrichment. The large number of simulated clusters enables us to derive very robust conclusions through a statistical analysis of the sample. The total mass is recovered adopting an observational-like approach and compared with the true mass in the simulations. Surface brightness and temperature profiles, that we generate from the simulations, are used to estimate the cluster mass at different overdensities ($r_{\\mathrm{2500}}, r_{\\mathrm{500}}$, and $r_{\\mathrm{200}}$) by means of the hydrostatic equilibrium equation. We explore various models and conditions under which the mass is reconstructed in order to entangle different mass biases. In addition, a power ratio analysis of the surface brightness maps allows us to assess the dependence of the mass bias on cluster dynamical state. Moreover, we perform a study on the reliability of  hydrostatic and hydrodynamical equilibrium mass estimates using the full three-dimensional gas distribution in the simulation.  In the following we list and discuss our main findings.  \\begin{enumerate} \\item Our analysis shows that it is very important to use analytical models with a large amount of parametric freedom when modeling the shape of the ICM temperature and surface brightness radial profiles. A model with a low degree of sophistication such as the polytropic $\\beta$-model can introduce a very large {\\it modeling bias}. Compared to the more sophisticated extended $\\beta$-model, which is found to be extremely accurate in following the slope changes of the gas profiles, it additionally leads to average mass underestimates of the order of $\\sim 5$, $10$, and $15\\%$ at $r_{\\mathrm{2500}}$, $r_{\\mathrm{500}}$, and $r_{\\mathrm{200}}$, respectively.  \\item The bias originating from extrapolating of the mass profiles beyond the radial range probed by observation can be extremely large. We find that the underestimate from extrapolation alone is of the order of $\\sim 10 \\%$ at $r_{\\mathrm{200}}$, and lower at smaller cluster-centric distances, when considering an average over the whole sample. However, for individual objects, the {\\it extrapolation bias} can be as large as  $\\sim 50 \\%$.  \\item The unrelaxed {\\it dynamical state} of a cluster can also lead to mass underestimates. The total mass is on average biased lower for disturbed clusters than for the relaxed ones. Furthermore, we find that the bias values are much more scattered around the mean for disturbed objects than for the relaxed ones. If mass-weighted temperature profiles are adopted, the mean mass bias is at most $\\sim -5 \\%$ at all cluster-centric distances for relaxed clusters, but can be as large as $\\sim -20 \\%$ for the most disturbed ones.  Our results are in excellent agreement with the findings of \\cite{je07}, as shown by the comparison of Fig. \\ref{fig:fig1} in this work and Fig. 9 (top panel) in their paper, where the correlation between the same quantities are shown. This agreement is of particular importance, considering that the analysis by \\cite{je07} is based on simulations performed with the adaptive mesh refinement code {\\it Enzo}. Our findings are also in good agreement with the results of \\cite{kay07}.  \\item Mass estimates derived using spectroscopic temperatures are lower that those derived from mass-weighted temperatures (i.e. the true gas temperatures). We find that the {\\it spectroscopic temperature bias}, i.e. the bias originating from spectroscopic temperatures alone, is of the order of $-10 \\%$ for the whole numerical sample. A similar value is found by \\cite{kay07}.  Mass underestimates derived adopting spectroscopic-like temperature profiles ($7$ and $17 \\%$ on average at $\\Delta=2500$ and 500 for the whole sample) are in good agreement with the values found by \\cite{nagai07} ($12$ and $16 \\%$ are the corresponding values), whose results are based on mock X-ray observations derived from simulations performed with an Eulerian code. We also find very good agreement by performing the same comparison for relaxed and unrelaxed clusters separately. We notice that total mass biases derived from the simulations by \\cite{nagai07} and adopting mass-weighted profiles are on average $-7 \\%$ at $r_{\\mathrm{500}}$ for relaxed clusters \\citep{lau07}. The corresponding value that we find from our Lagrangian simulations is $-5 \\%$. Considering the different nature of the numerical codes and the different implementation of the various physical processes, the agreement is extremely relevant.  Our results are in tension with the findings of \\cite{rasia06}, who find, on average, much stronger biases. We notice, however, that a fair comparison is not possible since the work by \\cite{rasia06} focussed only on 5 clusters.  Our analysis also shows that the spectroscopic bias depends on dynamical state. We find that the spectroscopic bias is rather small ($-2,-5$, and $-7 \\%$ at $\\Delta= 2500, 500$ and 200) for the most relaxed clusters and of the order of $-12 \\%$ for the most disturbed ones.  \\item While the mass bias derived from mass-weighted temperature profiles does not depend on true cluster mass, the one derived from spectroscopic-like temperature exhibits a strong {\\it mass dependence}. For clusters with $M(<r_\\mathrm{200}) < 1.8 \\times  10^{14} \\, h^{-1 }M_{\\odot}$ we find that, at all radii, the spectroscopic bias contributes on average to the total bias with less than $7 \\%$, and less than $4 \\%$ if one considers only the most relaxed clusters. On the other hand, for clusters with $M(<r_\\mathrm{200}) >1.8 \\times  10^{14} \\, h^{-1 }M_{\\odot}$ we find that the spectroscopic bias is larger: e.g. $-7 \\%$ and $-12 \\%$ for the relaxed clusters at $r_{\\mathrm{500}}$ and $r_{\\mathrm{200}}$, respectively.  \\item Even in the {\\it ideal case}, i.e. when the mass of relaxed clusters is estimated without extrapolation and adopting a very accurate model for the ICM density and temperature profiles, total masses are affected by the spectroscopic temperatures. In this case the mean total mass bias is $\\sim (-4,-3)$, (-12,-9), and $(-15,-3) \\%$ at $r_{\\mathrm{2500}}, r_{\\mathrm{500}}$, and $r_{\\mathrm{200}}$ for the (most, least) massive clusters in our simulated sample.  \\item Even if we assume that the true (mass-weighted) ICM temperature can be reconstructed, masses are biased low (by -2, -9, and $-12 \\%$ at $\\Delta=2500$, 500, and 200, respectively, when taking the average over the whole sample).  \\item The latter finding prompted us to investigate the possible {\\it violation of the hydrostatic equilibrium} by using the full, three-dimensional information provided by the simulations. In agreement with the results from our observational-like analysis, we find that the hydrostatic equilibrium assumption yields masses underestimated by $\\sim 10-15 \\%$. This implies that the origin of the bias is not originated by the X-ray reconstruction method. The same level of mass bias has been found by \\cite{ras04} and \\cite{burns07}.  \\item In order to elucidate the origin of this bias we compute masses using both {\\it hydrostatic and dynamical equilibrium} mass estimators. The dynamical equilibrium takes into account for both thermal and non-thermal pressure of the ICM. We find that the mass estimated through dynamical equilibrium are on average well estimated. The mean bias profiles of the \"feedback\" clusters simulated by \\cite{kay04} are in very good agreement with our average values at all overdensities.  \\item We explore the conditions of applicability of both estimators, i.e. spherical symmetry (torque parameter $\\tau_q\\ll1$) and small radial streaming velocities (radial Mach number $|<v_r>/c_s| \\ll1$). We find that the biases $b_{HE}$ and $b_{DE}$ do not depend on cluster true mass and that their scatter decreases with decreasing $\\tau_q$ and $|<v_r>/c_s|$.  \\item For clusters with small values of $\\tau_q < 0.15$ and $|<v_r>/c_s| < 0.1$ we find that the average mass underestimate found from the hydrostatic equilibrium estimator ($5 \\%$ at $r_{\\mathrm{2500}}$ and $10 \\%$ at $r_{\\mathrm{500}}$ and $r_{\\mathrm{200}}$) is extremely well corrected by adopting the dynamical equilibrium estimator.  \\end{enumerate}  Our results of course depend on the physics included in the simulation. Our modeling of the gas physical processes is incomplete and does not include, for example, energy feedback from active galactic nuclei (AGN). The inclusion of more realistic physics is particularly relevant in the cluster cores, since it is supposed to provide additional heating to the gas which could help to solve the cooling flow problem as well as the overcooling of baryons actually present in current simulations \\citep[e.g.][]{vo05}.  The effect on the thermal status of the ICM of energy feedback from the AGN are however not expected to modify in a significant way the ICM properties outside the cluster cores. The validity of this assumption is justified from the radial behavior of the measured temperature profiles, which are in good agreement at $ r \\gtrsim 0.1 \\times r_\\mathrm{200}$ with the present simulations \\citep{valda06} and the ones of other authors \\citep[e.g.][]{nagai07b}.   Importantly, we find that, in addition to cold gas clumps, also the diffuse cold gas component which is left after the removal of all resolved clumps substantially biases spectroscopic temperatures low. The good agreement of our temperature profiles with those found by other authors using different codes suggests that the amount of the cool gas component present in our simulations is not affected by numerical issues,   There is another issue, which is connected to the gas physical modeling in the simulations, that could be potentially relevant for our analysis. In the runs performed here the artificial viscosity is treated according to the standard SPH formulation \\citep{mo05}, which is a numerical scheme comparatively viscous \\citep{momo97} and inadequate to follow the development of fluid turbulence. It has been shown that the level of kinetic energy in random gas motion could be as high as $ \\sim 30 \\%$ of the gas thermal energy \\citep{dolag05,vazza06}, when the numerical viscosity scheme is generalized as in \\cite{momo97} to properly treat fluid turbulence.  It is worth noticing that very high levels of random motions are inconsistent with the upper limits recently derived from the comparison of X-ray and weak lensing mass estimates \\citep{mahdavi08,zhang08}. Direct measurements of the various sources of non-thermal pressure (gas motions, cosmic rays, magnetic fields, etc.) are extremely difficult, and the comparison between X-ray and lensing masses presently provides the most significant constraints on the level of non-thermal pressure in the ICM. \\cite{zhang08} obtain a ratio of $1.09 \\pm 0.08$ between the weak lensing and X-ray mass estimates extracted at $r_\\mathrm{500}$ from a sample of 19 clusters. At the same overdensity, \\citep{mahdavi08} find that the ratio between X-ray and lensing is $0.78 \\pm 0.09$ ($0.85 \\pm 0.10$ after correction for excess structure along the line of sight) for a sample of 18 clusters. The fairly good agreement between these recent observational results and the values found in the present analysis and other similar studies therefore implies a low level of turbulence present in the ICM, thereby suggesting that the gas physics outside the cluster cores is well approximated by the simulations presented here.  While we have shown that at present both theoretical models and observations seem to indicate a deviation from the hydrostatic equilibrium of the order of $ \\sim 10 \\%$, further investigations are needed in order to draw robust conclusions on this important issue. Future measurements of  X-ray and lensing masses for large samples and measurements of the ICM velocity structure will provide valuable information on the validity of the hydrostatic equilibrium and therefore on the reliability of X-ray clusters as cosmological probes.  Moreover, these measurements will also be of particular relevance for constraining gas physics in galaxy clusters. The amount of physical viscosity is in particular a key parameter which quenches the level of turbulence present in the ICM. Physical viscosity has been incorporated in hydrodynamic simulations of galaxy clusters by \\cite{si06} and it was shown that even a modest amount of  physical viscosity has significant consequences on ICM properties. Therefore X-ray mass estimates, extracted from a statistically meaningful sample of hydrodynamical SPH simulations which include  physical viscosity, could be profitably used to indirectly measure the level of ICM viscosity when contrasted against X-ray and weak lensing mass measurements.  These constraints will also likely have a significant impact on those scenarios in which the cooling flow problem is solved by providing additional heating to the gas through energy dissipation."
    },
    "0808/0808.3925_arXiv.txt": {
        "abstract": "{Nanoflares are small impulsive bursts of energy that blend with and possibly make up much of the solar background emission. Determining their frequency and energy input is central to understanding the heating of the solar corona. One method is to extrapolate the energy frequency distribution of larger individually observed flares to lower energies. Only if the power law exponent is greater than 2 is it considered possible that nanoflares contribute significantly to the energy input. } {Time sequences of ultraviolet line radiances observed in the corona of an active region are modelled with the aim of determining the power law exponent of the nanoflare energy distribution.} {A simple nanoflare model based on three key parameters (the flare rate, the flare duration, and the power law exponent of the flare energy frequency distribution) is used to simulate emission line radiances from the ions \\FeXIX, \\CaXIII, and \\SiIII, observed by SUMER in the corona of an active region as it rotates around the east limb of the Sun. Light curve pattern recognition by an Artificial Neural Network (ANN) scheme is used to determine the values. } {The power law exponents, $\\alpha\\approx2.8$, $2.8$, and 2.6 are obtained for \\FeXIX, \\CaXIII, and \\SiIII\\ respectively. } {The light curve simulations  imply a power law exponent greater than the critical value of 2 for all ion species. This implies that if the energy of flare-like events is extrapolated to low energies,  nanoflares could provide a significant contribution to the heating of active region coronae.} ",
        "introduction": "Heating the corona by the dissipation of current sheets was first suggested by \\citet{Gold64} and later developed to form the basis of the nanoflare heating model by \\citet{Levine74} and \\citet{Parker83, Parker88}. The idea is that current sheets arise spontaneously in coronal magnetic fields that are braided and twisted by random photospheric footpoint motions. These current sheets dissipate in many small-scale reconnection events, heating  and accelerating plasma in the coronal loops. In the corona, they would give rise to multiple unresolvable loop strands with specific observable signatures \\citep{Zirker94, WWH02, CK04, PK05}.  Recently \\citet{Aschwanden08} found evidence against such multi-temperature strands in TRACE coronal images. He concludes that nanoflare heating is only possible if it occurs in the chromosphere/transition region where heating across magnetic field lines can produce the isothermal loops seen in the corona. Irrespective of where the nanoflare energy input sites are, a key question is whether the energy of nanoflares is sufficient to heat the corona or not. Most of the individual nanoflares would be too small to detect and the majority would be small fluctuations on the overall background. That background could be produced by the blending of many small events.  The approach taken to estimate their contribution has been to extrapolate the energy frequency distribution of detectable flare-like events. The energy frequency distribution of larger flares tends to follow a power law distribution \\begin{equation} {dN\\over dE} \\sim E^{-\\alpha}, \\end{equation} where  $dN$ is the number of flares per energy interval $dE$.  The energy of small flares dominates if  $\\alpha>2$ \\citep{Hudson91}. This is therefore a critical parameter for the nanoflare heating model. The standard method to determine $\\alpha$ is to evaluate the energy of many flares in a series of observations and then plot their frequency in bins of energy $dE$. The majority of analyses based on this type of event counting deduce $\\alpha\\approx 1.7$ \\citep{Lin84, Shimizu95, AP02}, a value smaller than the critical 2. These results may, however, be misleading. For example, \\citet{Parnell04} demonstrated that one can obtain $\\alpha$ ranging from 1.5 to 2.6 for the same data set using different but still reasonable sets of assumptions for the analyses.  \\begin{figure} \\includegraphics[width=8.1cm,angle=0]{0911fig1.eps} \\caption{ EIT 195 \\AA ~images of the observed active region at two times, showing the position of the SUMER slit, indicated by the vertical line.} \\label{fig1} \\end{figure}  Here we take an alternative approach and model ultraviolet (UV) radiances observed by the Solar Ultraviolet Measurements of Emitted Radiation \\citep[SUMER;][]{Wetal95, Wetal97} in an active region corona, assuming that the radiance fluctuations and the nearly constant `background' emission are caused by small-scale stochastic flaring \\citep{PS04b, PS06}. The model has been applied successfully to UV radiance fluctuations in the quiet Sun \\citep{PS06}. The method compares light curves generated assuming random flaring with a power law frequency distribution to the light curves of an observed emission line. It has the advantage that it takes into account without bias weak, blended micro- and nanoflares that produce a nearly continuous background.  Here we apply this technique to off-limb time series recorded by SUMER. The three lines modelled, \\FeXIX~$\\lambda$\\,1118.07 (6.3~MK), \\CaXIII~$\\lambda$\\,1133.76 (2.2~MK) and \\SiIII~$\\lambda$\\,1113.23 (0.06~MK), cover two decades of formation temperature from the lower transition region to the hotter gas in the corona.  The analysis described here uses Artificial Neural Networks (ANNs) to find the optimum match to the three parameters of the model. The main advantage of this method over previous analyses based on the radiance distribution function   \\citep{PS06, SISP07} is that we are able to obtain quantitative values for all parameters, including $\\alpha$. Another advantage of the ANN method is that it concentrates on the number and shape of the emission peaks along the light curves with little weight on the low radiance pixels, which was a problem with the \\citet{SISP07} analysis. ",
        "conclusions": "\\label{results} \\subsection{Results}  In the present work,  PNN is used as a tool to extract the three flare model parameters required to reproduce the SUMER light curves.  All 35 \\FeXIX\\ and \\CaXIII, and 11 \\SiIII\\ SUMER light curves from the three days of observations were fed individually into the neural network and the parameters were obtained for each light curve separately. The final PNN outputs are shown in Table \\ref{tab1}. The bold numbers are the statistically maximum occurrence for each parameter. For example for \\FeXIX,  $\\alpha=2.8$ is found in more than 70\\% of the light curves. The minimum and maximum values, given on the left and right, indicate the scatter in the light curve parameters.  \\begin{table} \\centering \\caption{The SUMER spectral lines and the parameter values given by PNN.} \\label{tab1} \\smallskip \\begin{threeparttable} \\begin{tabular}{c c c c} \\hline\\hline\\\\ SUMER spectra  & \\multicolumn{3}{c}{PNN outputs\\tnote{1}}\\\\ lines & $\\alpha$ & $\\tau$ & $p_f$ \\\\ \\hline \\hline \\\\ \\FeXIX & 2.5 {\\bf 2.8} 3.1  & 9 {\\bf 14} 20 & 0.1 {\\bf 0.2} 0.7 \\\\ \\CaXIII & 2.6 {\\bf 2.8} 3.0 & 41 {\\bf 45} 45 & 0.8 {\\bf 0.9} 0.9 \\\\ \\SiIII & 2.4 {\\bf 2.6} 2.8 & 5 {\\bf 9} 12  & 0.2 {\\bf 0.3} 0.9 \\\\ \\hline \\end{tabular} \\begin{tablenotes} \\item[1]The most frequent values are given in bold, and the minimum and maximum on the left and right. $\\tau$ is per exposure time (90 s) and $p_f$ is per exposure time per 5\\arcsec$\\times$4\\arcsec\\ spatial element. \\end{tablenotes} \\end{threeparttable} \\end{table}  In each line there is  20\\% scatter in  $\\alpha$, and 50\\% scatter in $\\tau$. The range of $p_f$ values for \\FeXIX\\ and \\SiIII\\ is much broader, suggesting that events producing emission in these temperature ranges do not have the same rate everywhere but are seen in irregular bursts. We also note that the value of $\\tau p_f$ is roughly the same for both \\FeXIX\\ and \\SiIII, as suggested by their shape parameter (Fig.~\\ref{timeseries}). The \\CaXIII\\ light curves are all matched with a high value of $p_f$, consistent with the idea  that the 1~MK active region corona requires almost continuous flaring. The four times higher rate for \\CaXIII\\ than \\FeXIX\\ suggests that most of the \\CaXIII\\ emission is produced by heating events below the \\FeXIX\\ formation temperature (6.6~MK).  Example light curves obtained using these parameters are compared with the observed ones in Fig.~\\ref{lightcurves}. Both the \\SiIII\\ and \\CaXIII\\ simulations look remarkably similar to their observed light curves. The background radiance of the \\FeXIX\\ light curve is about a factor of 2 too low. The \\FeXIX\\ light curves had a $p_f$ ranging from 0.7 to 0.1, so we suspect that in this case the $p_f$ value is slightly too low. Also for \\FeXIX, the ratio $\\tau_r/\\tau_d$ deduced from the data is smaller than the fixed value 0.5 used here. This may influence the accuracy of the method.  The sensitivity of the PNN output depends on the training set. During the training session, the network must see all possible patterns that it is supposed to classify in the testing session. With 500 simulated light curves in the training set, PNN was not able to converge for several of the SUMER light curves. When we increased the number of simulated light curves to 6930, we were able to obtain unique parameters for all observed light curves.  \\begin{figure} \\includegraphics[width=7.5 cm]{0911fig6.eps} \\caption{Samples of the radiance time series: left panel: SUMER data, and right panel: simulation data obtained with the parameters given in Table 1. } \\label{lightcurves} \\end{figure}   \\subsection{Conclusions} The concept that the solar corona may be heated by numerous, randomly distributed, small flare-like events called nanoflares is considered by comparing simulated and observed emission line light curves. The difference between this and previous methods is the fully automated modelling of the light curve structure. There is no human decision required for background/event cut-off levels or best fit parameters.  The result is power law flare energy frequency exponents greater than 2.5 for all three emission lines considered, \\SiIII, \\CaXIII\\ and \\FeXIX. This is consistent with the corona being heated mainly by nanoflares, and demonstrates the importance of nanoflare 'background' emission in determining the power law exponents. The parameter with highest uncertainty or largest scatter is the flare rate, especially for the lines formed at transition region and hot flare temperatures. Coronal plasma at these temperatures is produced sporadically and is associated with more specific coronal and chromospheric loop structures than the general active region corona, so the scatter is to be expected.  The next step will be to determine the actual flare energies producing the nanoflare emission. This is a much more complicated exercise because the modelled light curves are observed in the corona which may be heated by events occurring lower in the atmosphere \\citep{Aschwanden08}, so that it requires a model for the energy transfer to the observation position."
    },
    "0808/0808.3578_arXiv.txt": {
        "abstract": "We introduce a new code for computing time-dependent continuum radiative transfer and non-equilibrium ionization states in static density fields with periodic boundaries.  Our code solves the moments of the radiative transfer equation, closed by an Eddingtion tensor computed using a long characteristics method.  We show that pure (i.e., not source-centered) short characteristics and the optically-thin approximation are inappropriate for computing Eddington factors for the problem of cosmological reionization.  We evolve the non-equilibrium ionization field via an efficient and accurate (errors $<1\\%$) technique that switches between fully implicit or explicit finite-differencing depending on whether the local timescales are long or short compared to the timestep.  We tailor our code for the problem of cosmological reionization.  In tests, the code conserves photons, accurately treats cosmological effects, and reproduces analytic Str\\\"omgren sphere solutions.  Its chief weakness is that the computation time for the long characteristics calculation scales relatively poorly compared to other techniques ($\\tlc \\propto N_{\\rm{cells}}^{\\sim1.5}$); however, we mitigate this by only recomputing the Eddington tensor when the radiation field changes substantially.  Our technique makes almost no physical approximations, so it provides a way to benchmark faster but more approximate techniques.  It can readily be extended to evolve multiple frequencies, though we do not do so here.  Finally, we note that our method is generally applicable to any problem involving the transfer of continuum radiation through a periodic volume. ",
        "introduction": "\\label{sec:intro}  The epoch of reionization is the current frontier in understanding how galaxies form and evolve over cosmic time.  After the Universe cooled sufficiently to recombine hydrogen atoms at redshift $z\\approx 1088$~\\citep{spe07}, the Universe was fully neutral.  Gravity grew ever-denser structures that, at $z\\sim 30-50$, were able to collapse into stars and/or black holes.  The radiation emitted from these first objects then began to re-ionize hydrogen.  By $z\\sim 6$, hydrogen reionization appears to be complete~\\citep{fan07}, and the diffuse intergalactic medium (IGM) has a neutral fraction of $\\sim 10^{-4}$.  Understanding this transition epoch is central to understanding the origin of galaxies and the evolution of the IGM.  It is a major science driver for a host of upcoming international telescope facilities, such as the {\\it James Webb Space Telescope} and the {\\it Atacama Large Millimeter Array}.  Reionization involves a complex interplay between nonlinear growth of structure, radiative cooling, star/black hole formation, chemical enrichment, and photon transport.  Numerical simulations are required to accurately model these highly nonlinear processes.  However, the large dynamic range and complex physics involved make this an extraordinarily challenging computational problem.  To obtain a full picture of reionization in the context of currently-favored hierarchical structure formation models, it is imperative that simulations include processes of star formation, galaxy formation, and IGM evolution, along with feedback processes that connect all three.  Cosmological hydrodynamic simulations accounting for these processes are now achieving maturity, thanks to improving algorithms and computing power.  However, the inclusion of radiation transport complicates matters immensely.  A cosmological radiative hydrodynamics code that can accurately evolve a representative volume with sufficient dynamic range to study how galaxies reionize the Universe would be a major development towards understanding reionization. In this paper we provide a step towards that end by introducing a new accurate moment-based method for calculating radiative transfer (RT) in a cosmological context.  Time-dependent radiative transfer is one of the most difficult components to treat in any theoretical study of the reionization epoch owing to the problem's well-known high dimensionality.  Consequently, over the past decade, a number of approximate treatments have emerged that seek to render it more tractable through well-motivated physical approximations. The most flexible methods are the fully analytic treatments~\\citep[for example,][]{mad99,wyi03,fu04,ili05,kra06}.  These generally involve assuming values for quantities such as the gas clumping factor and the recombination rate that are averaged over all space or, in the case of the excursion-set formalism~\\citep{fu04}, over the volume of an ionized region.  In exchange, they readily allow for broad surveys of parameter space to be performed.  The next step in the direction of a full solution is taken by the semi-numerical methods~\\citep{cia00,mes07,gei08,cho08}, which combine numerically-generated density fields with analytic treatments for radiative transfer using techniques such as the excursion-set formalism in order to account more realistically for source bias and the effects of inhomogeneous density fields.  These treatments offer a dramatic increase in realism over purely analytic calculations at modest additional computational cost.  However, they have some difficulty accounting fully for the consequences of inhomogeneous density fields such as shadowing and the tendency for low-density regions to have a lower neutral fraction during the later stages of reionization~\\citep{cho08}.  These problems arise from the need of semi-numerical models to make assumptions regarding the shape of the ionized regions surrounding individual sources and the nontrivial relationship between dark matter and gas densities in the nonlinear regime.  Some of these difficulties are avoided in models that actually solve the radiative transfer equation on numerically-generated density fields but without fully accounting for radiative feedback on the sources (\\citealt{cia01,sok01,mel06,mcq07,ili07a}; see also \\citealt{ili06} for a very useful comparison of a number of techniques).  Nonetheless, obtaining realistic baryonic density and emissivity fields in such contexts still presents considerable challenges~\\citep{mcq07}.  Additionally, while parametrized treatments for radiative feedback have been introduced in such models in order to study, for example, whether the photoevaporation of minihaloes extends the epoch of reionization~\\citep{ili07a,mcq07}, the simplified nature of these studies leaves their results open to question~\\citep{mes07}.  Hence, while each of these methods has yielded an abundance of insight into reionization and warrant continued development, the need is emerging for a complete solution to the radiative transfer equation that is merged self-consistently with hydrodynamical calculations.  A few fully radiative-hydrodynamic codes have been used to study the reionization epoch~\\citep{gne01,cen02,rij06}.  Unfortunately, the techniques that these codes have introduced do not yet enjoy widespread use owing to their high computational expense, even though such studies are crucial for tuning the assumptions employed in more simplified treatments and verifying their conclusions.  Moreover, despite the enormous gain in realism that these codes have afforded the reionization community, they too involve some physical approximations such as the use of the ``optically thin variable Eddington tensor\" approximation~\\citep{gne01} and a reduced~\\citep{gne01} or an infinite speed of light~\\citep{cen02,rij06}.  In this work, we present a moment method solution to the radiative transfer equation~\\citep{aue70} and test it on static density fields. Our technique is highly flexible, involves a minimum of physical approximations, and can readily be combined with existing hydrodynamical calculations.  It is similar to the method presented by~\\citet{sto92}, but with several differences.  First, we optimize our code only for cubical simulation volumes with periodic boundaries, as this is typical of cosmological simulations.  Second, we derive our Eddington tensors from a long characteristics (LC) calculation in order to minimize artifacts owing to poor angular and spatial resolution.  Finally, we include a treatment for nonequilibrium ionizations and account for the cosmological terms in the radiative transfer equation.  In a follow-up paper, we will present its implementation within a cosmological galaxy formation code.  We begin in Section~\\ref{sec:rt} by casting the radiative transfer equation into the form in which we solve it and summarizing our numerical method.  In Section~\\ref{sec:fedd}, we compare the performance of long characteristics versus two other time-independent radiative transfer techniques in order to select a method for deriving the Eddington tensor, which we need in order to close our moment hierarchy.  After demonstrating that long characteristics introduces the fewest unphysical artifacts, we optimize it for computing reionization using a suite of realistic albeit low-resolution integrations.  In Section~\\ref{sec:nlte}, we discuss our technique for evolving the nonequilibrium ionization field. In Section~\\ref{sec:full_alg}, we summarize our iterative scheme for weaving these ingredients into a self-consistent calculation. In Section~\\ref{sec:tests}, we subject our code to a number of standard tests.  Finally, we summarize our method and results in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this paper, we introduced a method that accurately and efficiently computes continuum radiative transfer in static density fields.  The code uses a moment-based approach to solve the equation of comoving radiative transfer, with the Eddington tensors obtained using a long characteristics method.  We compared several techniques for computing the Eddington tensors that are needed to close the moment hierarchy and demonstrated that, of the three methods that we investigated, only the method of long characteristics has the ability to compute highly inhomogeneous radiation fields without introducing numerical artifacts.  We found through direct measurement that the computation times for our long characteristics and moments modules scale with the number of computational cells $N_{\\rm{cells}}$ as $N_{\\rm{cells}}^{1.5}$ and $N_{\\rm{cells}}^{1.0}$, respectively.  Next, we introduced a hybrid method for computing the evolution of nonequilibrium ionization fields and demonstrated that it is accurate to 1\\% throughout our computational domain.  We combined this with our radiative transfer code via an efficient iterative algorithm.  The final code is regulated by a number of parameters, and we characterized how these parameters impact computation time and accuracy using a suite of low-resolution convergence tests.  We subjected our method to a number of standard problems in continuum radiative transfer.  First, we verified that our code accurately computes the growth of an \\hii\\ region about a source in the classical (static) case as well as in the case of an expanding medium.  We found that, in both cases, the radius of the resulting ionized region agrees with analytic expectations to within the computation's spatial resolution at all times. Next, we tested whether our code is able to produce shadows behind opaque regions.  In agreement with previous work, we found that the moment method introduces a small diffusion into the shadowed region~\\citep{hay03}. Nevertheless, we found that the strength of the radiation field in the shadowed region was up to only 1\\% of the value in the unshadowed region. Finally, we computed the reionization of an expanding cosmological volume and found qualitative agreement with other work in the literature.  Our code currently accounts for radiative and collisional ionization of hydrogen and helium as well as radiative recombination.  It does not account for recombination radiation.  Additionally, it does not follow the evolution of the temperature field, hence it does not account for photoionization suppression of star formation in low-mass halos, photoionization heating, recombination cooling, or shock formation.  In the future we plan to expand on our code in several ways.  First, we have found that the computation time increases dramatically as the universe becomes optically thin because the LC line integrals traverse more cells before terminating either because they arrive at the source or because they reach $\\taumax$.  However, in this regime, the optically-thin approximation, which is roughly ten times faster than LC, becomes increasingly valid.  For this reason, we plan to study how to transition smoothly from LC to the optically thin approximation without introducing accuracy errors.  Second, we will generalize our method to multifrequency radiative transfer, which is necessary for studying, for example, ionization front hardening or \\heii\\ reionization. Finally, we plan to merge our radiative transfer scheme with our version of the cosmological galaxy formation code {\\sc GADGET-2}."
    },
    "0808/0808.1261_arXiv.txt": {
        "abstract": "Due to quantum fluctuations, spacetime is foamy on small scales.  The degree of foaminess is found to be consistent with the holographic principle. One way to detect spacetime foam is to look for halos in the images of distant quasars. Applying the holographic foam model to cosmology we \"predict\" that the cosmic energy density takes on the critical value; and basing only on existing archived data on active galactic nuclei from the Hubble Space Telescope, we also \"predict\" the existence of dark energy which, we argue, is composed of an enormous number of inert ``particles\" of extremely long wavelength.  We speculate that these ``particles\" obey infinite statistics. ",
        "introduction": "Like everything else, spacetime is conceivably subject to quantum fluctuations.  So we expect that spacetime, probed at a small enough scale, will appear complicated --- something akin in complexity to a turbulent froth that John Wheeler has dubbed \"quantum foam,\" also known as \"spacetime foam.\"\\footnote{In the gravitational context, the phenomenon of turbulence is indeed intimately related to the properties of spacetime foam.  See Ref. \\cite{Jejjala2008}.}  But how large are the fluctuations in the fabric of spacetime? To quantify the problem, let us recall that, if spacetime indeed undergoes quantum fluctuations, there will be an intrinsic limitation to the accuracy with which one can measure a distance, for that distance fluctuates.  Denoting the fluctuation of a distance $l$ by $\\delta l$, on general grounds, we expect $\\delta l \\gtrsim l^{1 - \\alpha} l_P^{\\alpha}$, where $l_P = \\sqrt{\\hbar G/c^3}$ is the Planck length, the characteristic length scale in quantum gravity, and we have denoted the Planck constant, gravitational constant and the speed of light by $ \\hbar$, $G$ and $c$ respectively. The parameter $\\alpha \\sim 1$ specifies the different spacetime foam models.  In this talk we will concentrate on the model corresponding to $\\alpha = 2/3$, which has come to be known as the holographic model \\cite{wigner,Karol}, so called because it is found to be consistent with the holographic principle \\cite{tHooft,susskind}, according to which, the information content inside any three dimensional region of space can be encoded on the two dimensional surface around the region, like a hologram. For comparison, we will also consider the random-walk model \\cite{Amelino1999} corresponding to $\\alpha = 1/2$.  Contents of this talk: Applying nothing more than quantum mechanics and some rudimentary black hole physics, we derive the holographic model of spacetime foam. Applying the holographic model to cosmology, we \"predict\" that the cosmic energy density takes on the critical value (i.e., the fractional density parameter of the universe $\\Omega \\cong 1$), consistent with observation.  Then aided by some archived data on quasar or AGN from the Hubble Space Telescope, we are led to conclude that dark energy exists. Furthermore we are naturally led to speculate that the constituents of dark energy, unlike ordinary matter, obey an exotic statistics known as infinite statistics in which all representations of the particle permutation group can occur. ",
        "conclusions": ""
    },
    "0808/0808.0267_arXiv.txt": {
        "abstract": "The primary scientific goal of the GRIPS mission is to revolutionize our understanding of the early universe using $\\gamma$-ray bursts. We propose a new generation gamma-ray observatory capable of unprecedented spectroscopy over a wide range of $\\gamma$-ray energies (200 keV--50 MeV) and of polarimetry (200--1000 keV). Secondary goals achievable by this mission include direct measurements of supernova interiors through $\\gamma$-rays from radioactive decays, nuclear astrophysics with massive stars and novae, and studies of particle acceleration near compact stars, interstellar shocks, and clusters of galaxies. ",
        "introduction": "Gamma-ray bursts (GRB) are the most luminous sources in the sky, and thus act as signposts throughout the Universe. The long-duration sub-group is produced by the explosion of massive stars, while short-duration GRBs likely originate during the merging of compact objects. Both types are intense neutrino sources, and being stellar sized objects at cosmological scales, they connect different branches of research and thus have a broad impact on present-day astrophysics.  Identifying objects at redshift \\gax 6.5 has become one of the main goals of modern observational cosmology, but turned out to be difficult. GRBs offer a promising opportunity to  identify high-$z$ objects, and moreover even allow us to investigate the host galaxies at these redshifts. GRBs are a factor 10$^{5-7}$ brighter than quasars during the first hour after explosion, and a favourable relativistic k-correction implies that they do not get fainter beyond $z$$\\sim$3. Yet, present and near-future ground- and space-based sensitivity limits the measurement of redshifts at $z$$\\sim$13 (as $H$-band drop-outs), because GRB afterglows above 2.5 $\\mu$m are too faint by many magnitudes % for 8--10\\,m telescopes. % Thus, a completely different strategy is needed to step beyond redshift 13  to measure when the first stars formed.  Fortunately, nuclear physics offers such a new strategy. Similar to X-ray and optical absorption lines due to transitions between electronic levels, resonant absorption processes in the nuclei exist which leave narrow absorption lines in the $\\gamma$-ray range. The most prominent and astrophysically relevant are the nuclear excitation and Pygmy resonances (element-specific narrow lines between 5--9 MeV), the Giant Dipole resonance  (GDR; proton versus neutron fluid oscillations; $\\sim$ 25 MeV; two nucleons and more) and the Delta-resonance (individual-nucleon excitations, 325 MeV; all nucleons, including H!). Such resonant absorption only depends on the presence of the nucleonic species, and not on ionization state and isotope ratio. They imprint well-defined spectral features in the otherwise featureless continuum spectra of GRBs (and other sources). This is completely new territory (Iyudin \\etal\\ 2005), but with the great promise to measure redshifts directly from the gamma-ray spectrum, i.e. without the need for optical/NIR identification!  Technically, this new strategy requires sensitive spectroscopy in the 0.2--50 MeV band. The detection of GRBs requires a large field of view. Therefore, the logical detection principle is a Compton telescope. In addition, such detectors  can be tailored to have a high polarisation sensitivity. Polarimetry is the last property of high-energy electromagnetic radiation which has not been utilized in its full extent, and promises to uniquely determine the emission processes in GRBs, as well as many other astrophysical sources. With its large field of view, such a detector will not only scan 80\\% of the sky within one satellite orbital period of 96 min., but also provide enormous grasp for measuring the diffuse emission of nucleosynthesis products and cosmic-ray acceleration. ",
        "conclusions": ""
    },
    "0808/0808.2180_arXiv.txt": {
        "abstract": "We have determined the distance to NGC~4258 using observations made with the Hubble Space Telescope (HST) and the Wide Field, Advanced Camera for Surveys (ACS/WFC). We apply a modified technique that fully accounts for metallicity effects on the use of the luminosity of the tip of the red giant branch (TRGB) to determine one of the most precise TRGB distance moduli to date: $\\mu(TRGB) = 29.28 \\pm 0.04$ (random) $\\pm 0.12$ (systematic) mag ($7.18 \\pm 0.13 \\pm 0.40$ Mpc). We discuss this distance modulus with respect to other recent applications of the TRGB method to NGC~4258, and with several other techniques (Cepheids and masers) that are equally competitive in their precision, but different in their systematics. ",
        "introduction": "This is the first in a short series of papers using a refined methodology for determining distances using the discontinuity in the I-band magnitude of the red giant branch luminosity function as a standard candle (Lee et al. 1993), the so-called TRGB (tip of the red giant branch) method. Here we apply a new methodology in correcting for the now well understood and precisely calibrated metallicity effects on the TRGB magnitude (see Section 4.1 and Madore et al. 2008).  Our first target is the spiral galaxy NGC~4258. It is nearby, and therefore very highly resolved, not only into its bright, high-mass Population I disk stars, but also into its fainter, but still accessible, low-mass Population II halo stars. NGC~4258 contains many known Cepheids that have been discovered and used as distance indicators in multiple observing campaigns using HST. Its halo has been resolved and studied on equally as many occasions, revealing a broad, richly populated giant branch for TRGB distance determination. The uniqueness of NGC~4258 lies at its center, where a Keplerian-rotating disk of water masers has proper motions and radial velocities that can be cross-compared and modeled with essentially one additional free parameter: the distance. As such, the independently calibrated Population~I (Cepheid) and Population~II (TRGB) distance scales both converge on and cross at NGC~4258, where they can be compared to that from simple geometry (maser method). No other galaxy provides such an environment for testing the distance scale. That said, it must also be emphasized that NGC~4258 is still only one object, and its uniqueness means that there is no independent check on the maser distance methodology itself, its random errors, or its systematics.  Without prejudice as to which (if any) of the three distance determination methods discussed here is better (understood or calibrated) at this point, we now proceed to present a new and improved determination of the TRGB distance using HST ACS/WFC data from one of our approved and scheduled programs, and archival WFPC2 data as a consistency check. We compare these results with previous TRGB results, and with the other past and published methods. ",
        "conclusions": ""
    },
    "0808/0808.0185_arXiv.txt": {
        "abstract": "We have acquired near-infrared spectra of Kuiper belt objects 2003 UZ117, 2005 CB79 and 2004 SB60 with NIRC on the Keck I Telescope. These objects are dynamically close to the core of the 2003 EL61 collisional family and were suggested to be potential fragments of this collision by \\citet{darin}. We find that the spectra of 2003 UZ117 and 2005 CB79 both show the characteristic strong water ice absorption features seen exclusively on 2003 EL61, its largest satellite, and the six other known collisional fragments. In contrast, we find that the near infrared spectrum of 2004 SB60 is essentially featureless with a fraction of water ice of less than 5\\%. We discuss the implications of the discovery of these additional family members for understanding the formation and evolution of this collisional family in the outer solar system. ",
        "introduction": "The only known collisional family in the Kuiper belt was detected because of the unique spectral properties the family members (Brown et al. 2007). The near-infrared spectra of all other non-volatile rich Kuiper belt objects (KBOs) lie on a continuum between those whose spectra contain moderate amounts of water ice absorptions to those whose spectra  are  essentially featureless \\citep{2007Natur.446..294B, kris08}. In contrast, the near-infrared spectra of 2003 EL61, its largest satellite, and five other small KBOs resemble laboratory spectra of pure crystalline water ice \\citep{1999ApJ...519L.101B,kris06,2007ApJ...655.1172T,2007A&A...466.1185M,2007Natur.446..294B, kris08}. Remarkably, these objects are also relatively clustered in orbital element space. The largest object, 2003 EL61, had been previously suggested to have experienced a massive collision that imparted its fast (4-hour) rotation, stripped off most of its icy mantle leaving it with a density close to rock (2.7 g/cc), and formed its two satellites \\citep{2006ApJ...639.1238R, 2006ApJ...639L..43B, 2008AJ....135.1749L}. \\citet{2007Natur.446..294B} concluded that the four extremely water ice rich KBOs and the large satellite of 2003 EL61 were in fact fragments of the icy mantle of the proto-2003 EL61 that had been ejected during a massive collision.   While collisional families in the asteroid belt can be identified by their dynamical clustering alone, families in the Kuiper belt are much harder to identify because the collisional ejection velocities can be a significant fraction of an object's orbital velocity.  Collisional fragments can therefore be spread out over a wide range of orbital element space with the fastest ejected fragments having significantly different orbits from the family core. The 2003 EL61 family members were all identified on the basis of their unique surface properties: they are the only known objects in the Kuiper belt with extremely pure water ice spectra and neutral visible colors. It is only because of this unique surface signature that detection of the 2003 EL61 family members was possible.  In order to determine if other known KBOs could be fragments of the 2003 EL61 collision, \\citet{darin} integrated the orbits of 131 high inclination KBOs to  determine their proper orbital elements and minimum ejection velocities away from the 2003 EL61 family core.  They determined the minimum ejection velocity a KBO must have had in order to reach its present orbit from the modeled location of the family forming impact. In Figure 1 we show object H-magnitude vs. minimum ejection velocity (from \\citet{darin}) for the known KBOs closest to the family core.  For KBOs in known resonances, we plot the ejection velocity accounting for resonance diffusion \\citep{darin}. A large fraction of these objects have never been observed to determine if they have IR spectral signatures consistent with the 2003 EL61 family members.   It is important to note that though all of the known 2003 EL61 family members have neutral visible colors, visible color or visible spectroscopy alone \\citep{PinAlons2008} is not sufficient for family member identification. There are many KBOs with neutral visible colors that do not also have strong water ice absorptions and are not members of the 2003 EL61 collisional family. Therefore, while visible colors can be used to rule out objects as potential family members, near infrared spectroscopy to determine water ice absorption depths is necessary for definitive family member identification.  Identifying additional family members and characterizing the extent of the spread of fragments throughout the Kuiper belt may provide insight into the physics of this giant impact event in the outer solar system. In this letter we present near-infrared spectra obtained with the Keck I telescope of KBOs 2003 UZ117, 2005 CB79 and 2004 SB60, objects that are located 67, 97, and 221 m/s away from the modeled 2003 EL61 family core respectively (Fig 1). 2003 UZ117 and 2005 CB79 were suggested to be the most likely additional family candidates by \\citet{darin}. ",
        "conclusions": "2003 UZ117 and 2005 CB79 contain high fractions of pure water ice on their surfaces and appear dynamically related to the other members of the 2003 EL61 collisional family.  We therefore conclude that they are members of this collisional family. Thus far, all objects within 130 m/s of the modeled collision center \\citep{darin} have been shown to have the same unique strong water ice spectral signatures. 2004 SB60, at 221 m/s from the family core, has an essentially  featureless spectrum inconsistent with 2003 EL61 and its family members but comparable to many other small KBOs. The flat spectrum we observed is consistent with infrared photometric observations of 2004 SB60 with HST by  \\citet{2007DPS....39.3906S}. Infrared spectral or spectrophotometric observations of more potential fragments with minimum ejection velocities from $\\sim$150 to $\\sim$300 m/s could help constrain the extent to which fragments were scattered throughout the Kuiper belt and could tell us about the physics of the giant impact itself. In addition to the distribution of fragments, any collisional model would also need to explain the the presence of the two satellites of 2003 EL61.  Though it does not have the spectral signature of the fragments of the 2003 EL61 collision, 2004 SB60 is an interesting object in its own right. \\citet{2008Icar..194..758N} have found this object to be one of the few high inclination small binary objects in the Kuiper belt. It is worth noting the possibility that not all fragments of the 2003 EL61 collision necessarily have the unique strong water ice signature. Fragments from different locations in the initial parent body may have had different initial compositions. However, without this unique spectral signature, identification of an object as a family member is currently impossible.  Another object near the center of the 2003 EL61 family is the third largest KBO, 2005 FY9.  2005 FY9 has a velocity closer to the center of the collision (150 m/s) than one of the known fragments but has a spectrum dominated by methane, not water ice absorptions. We find it intriguing that two of the largest KBOs are so close to each other in orbital element space but can think of no reason why they should be physically related.  The spectral signature of extremely pure water ice found on 2003 EL61 and its fragments and nowhere else in the Kuiper belt is the only reason definitive identification of the collisonal family was possible. A collision with a non differentiated parent body would likely not produce such a distinct spectral signature in its fragments. Future large scale surveys of the Kuiper belt such as PanSTARRS and LSST are expected to increase the numbers of known KBOs by over an order of magnitude.  With higher numbers of objects, overdensities in certain regions may be revealed and families without unique spectral signatures may be able to be identified by their dynamics alone."
    },
    "0808/0808.3904_arXiv.txt": {
        "abstract": "{} {We present an accurate characterisation of the high-resolution X-ray spectrum of the Narrow Line Seyfert 1 galaxy Arakelian 564 and put it in to context with other objects of its type by making a detailed comparison of their spectra.} {The data are taken from 5 observations with the \\textit{XMM-Newton} Reflection Grating Spectrometer and fitted with various spectral models.} {The best fit to the data identifies five significant emission lines at 18.9, 22.1, 24.7, 29.0 and 33.5\\AA~due O\\,VIII Ly$\\alpha$, O\\,VII(f), N\\,VII Ly$\\alpha$, N\\,VI(i) and C\\,VI Ly$\\alpha$ respectively. These have an RMS velocity of $\\sim1100$\\,km\\,s$^{-1}$ and a flow velocity of $\\sim-600$\\,km\\,s$^{-1}$, except for the O\\,VII(f) emission line, which has a flow velocity consistent with zero. Two separate emitting regions are identified. Three separate phases of photoionized, X-ray absorbing gas are included in the fit with ionization parameters log\\,$\\xi=-0.86$, 0.87, 2.56 and column densities $N_H=0.89, 2.41, 6.03\\times10^{20}$\\,cm$^{-2}$ respectively. All three phases show this to be an unusually low velocity outflow ($-10\\pm100$\\,km\\,s$^{-1}$) for a narrow line Seyfert 1. We present the hypothesis that the BLR is the source of the NLR and warm absorber, and examine optical and UV images from the \\textit{XMM-Newton} Optical Monitor to relate our findings to the characteristics of the host galaxy.} {} ",
        "introduction": "\\label{564intro}  Narrow Line Seyfert 1 galaxies (NLS1s) were initially classified by \\cite{osterbrock85} and are defined as having H$\\beta$ FWHM\\,$\\leq$\\,2000\\,km\\,s$^{-1}$. Their X-ray spectra often display more rapid variability (e.g.~\\nocite{boller96}Boller et al. 1996;~\\nocite{laor94}Laor et al. 1994) and are steeper in the soft X-ray band than those of `normal' broad line Seyfert 1s (e.g.~\\nocite{leighly99}Leighly 1999). From the \\textit{ROSAT} All Sky Survey it was found that approximately half of the sources in  soft X-ray selected samples of AGN are NLS1s (\\nocite{grupe96}Grupe 1996;~\\nocite{hasinger97}Hasinger 1997). They are believed to be `high-state' active galaxies, with low black hole masses for their luminosity, and so with high accretion rates relative to Eddington (\\nocite{pounds95}Pounds et al. 1995;~\\nocite{boroson02}Boroson 2002). \\newline  Observations indicate that the majority of Seyfert 1s display evidence for `warm absorbers' (\\nocite{blustin05}Blustin et al. 2005), i.e. clouds of photoionized X-ray absorbing gas in our line of sight (e.g.~\\nocite{komossa00}Komossa 2000), yet still very little is known of the absorbers origin, location or structure. The warm absorber is often characterised by its ionization parameter, $\\xi$. This is a measure of the ionization state of the material and is defined as  \\begin{equation} \\xi=\\frac{L}{nr^2} \\label{564_xi_eq} \\end{equation} where $L$ is the source luminosity (in erg\\,s$^{-1}$), $n$ is the gas number density (in cm$^{-3} $) and $r$ is the distance of the absorber from the central engine (in cm) (\\nocite{tarter69}Tarter et al. 1969). \\newline  Arakelian 564 is the brightest known NLS1 in the 2--10\\,keV range (L$_{2-10\\,keV}=2.4\\times10^{43}\\,$erg\\,s$^{-1}$, \\nocite{turner01} Turner et al. 2001). This, coupled with its close proximity, z=0.02468 (\\nocite{huchra99}Huchra et al. 1999), and relatively low Galactic column density, N$_H=5.34\\times10^{20}$cm$^{-2}$ (\\nocite{kalberla05}Kalberla et al. 2005), make it a very interesting target to investigate. This object has been studied across all wavebands (e.g.~\\nocite{shemmer01}Shemmer et al. 2001;~\\nocite{romano04}Romano et al. 2004). Here we have merged data from 5 observations with the \\textit{XMM-Newton} Reflection Grating Spectrometer (RGS), allowing us to explore, in unprecedented detail, the high-resolution X-ray spectrum of this object. \\newline  \\citet{matsumoto04} report on a 50\\,ks \\textit{Chandra} HETGS observation. They fit the hard X-ray spectrum with a power law of photon index of $2.56\\pm0.06$ and add a blackbody, of temperature $0.124\\pm0.003$\\,keV, to fit the soft excess. They detect an edge-like feature at 0.712\\,keV, and two phases of absorption by photoionized gas (phase 1: log\\,$\\xi\\sim1$, N$_{H}\\sim10^{21}$\\,cm$^{-2}$; phase 2: log\\,$\\xi\\sim2$, N$_{H}\\sim10^{21}$\\,cm$^{-2}$).~\\citet{vignali04} have analysed two \\textit{XMM-Newton} observations from June 2000 and June 2001. They fit the spectrum with a steep power law ($\\Gamma=2.50-2.55$) plus a soft blackbody component (kT$\\sim$140\\,--\\,150\\,eV). They also identify an edge-like feature in the EPIC data at $\\sim0.73$\\,keV, which they interpret as an O\\,VII K absorption edge, and find evidence for a broad iron emission line, at $\\sim1$\\,keV, in one observation. ~\\citet{crenshaw02} have analysed the UV spectrum using data from the STIS onboard the \\textit{Hubble Space Telescope}. They identify UV absorption lines due to Ly$\\alpha$, N\\,V, C\\,IV, Si\\,IV and Si\\,III centred at a radial velocity of $-190$\\,km\\,s$^{-1}$ relative to the systemic velocity. They identify this with a dusty lukewarm absorber with column density N$_{H}=1.62\\times10^{21}$\\,cm$^{-2}$ and ionization U=0.033 (this ionization is equivalent to log\\,$\\xi\\simeq1.2$, converted using the method of \\nocite{george98}George et al. 1998). \\newline  \\citet{papadakis06} analysed the most recent (2005), and longest ($\\sim100$\\,ks), observation with \\textit{XMM-Newton}, focusing on the EPIC data. They find that the soft excess cannot be parametrized either by a multiple power law or by a power law plus a single blackbody, in contrast to previous findings. They fit the MOS and PN spectra with either a power law of slope $2.43\\pm0.03$ and two blackbodies, with kT$\\sim$0.15 and $\\sim$0.07\\,keV, or with a relativistically blurred photoionized disk reflection model. They also find evidence for two phases of photoionized X-ray absorbing gas with ionization parameters log\\,$\\xi\\sim1$ and $\\sim2$ and column densities N$_{H}\\sim2$ and $5\\times 10^{20}$\\,cm$^{-2}$. \\newline  The same 2005 data were also analysed by~\\citet{dewangan07}. They concentrate their analysis on the EPIC data, but also study the RGS spectrum and find two warm absorber phases with ionization parameters log\\,$\\xi\\sim2$ and $<0.3$, column densities $N_H\\sim4\\times10^{20}$ and $2\\times10^{20}$cm$^{-2}$ and flow velocities $v\\sim-300$ and $-1000$\\,km\\,s$^{-1}$ respectively. \\newline  Here we investigate the combined RGS spectrum of Arakelian 564. Section~\\ref{obs_data_red} gives information on the individual observations and describes our data reduction process. The spectral fitting and best fit results are presented in Section~\\ref{analysis}. A full discussion of these results and comparison with other observations is given in Section~\\ref{discussion}. ",
        "conclusions": "We present a detailed analysis of the photoionized, X-ray absorbing and emitting gas present in the NLS1 galaxy Arakelian 564. The high-resolution X-ray spectra are from five observations with \\textit{XMM-Newton}'s Reflection Grating Spectrometer. \\newline  The absorption profile is fitted with three phases of X-ray absorbing gas with average RMS velocity $\\sim60$\\,km\\,s$^{-1}$ and flow velocity $\\sim-10$\\,km\\,s$^{-1}$. Each phase has a different ionization (log $\\xi=-0.86, 0.87, 2.56$) and column density ($N_H=0.89\\times10^{20}, 2.41\\times10^{20}, 6.03\\times10^{20}$\\,cm$^{-2}$ respectively). Note that one continuous distribution of gas is unlikely, but not ruled out. The low flow velocity of the warm absorber phases can be partially explained by assuming the low ionization phases to be interstellar gas. The lower ionization phases are expected to contribute significantly to the UV spectrum, implying that the UV and X-ray absorbers are connected. \\newline  The best fit to the RGS data requires five significant emission lines due to O\\,VIII Ly$\\alpha$ (18.9\\AA), O\\,VII(f) (22.1\\AA), N\\,VI(i) (29.0\\AA), N\\,VII Ly$\\alpha$ (24.7\\AA) and C\\,VI Ly$\\alpha$ (33.5\\AA). By examining the outflow velocities and density restrictions on these emission lines we have shown that they originate in two separate regions, the BLR and the NLR. It is possible that these regions and the absorber are connected. The outflowing broad line region expands as it travels out from the core, maintaining a similar ionization level. This gas then picks up dust from the torus, and slows down. As the gas continues to expand it then displays the characteristics of the NLR and the diffuse, dusty warm absorber. UV observations suggest a disturbed host galaxy. A tidal stream or dust and gas accumulating in the inner regions of the host galaxy offer an alternative origin for the low ionization absorber."
    },
    "0808/0808.1289_arXiv.txt": {
        "abstract": "Observations using the {\\it Spitzer Space Telescope} provided the first detections of photons from extrasolar planets. {\\it Spitzer} observations are allowing us to infer the temperature structure, composition, and dynamics of exoplanet atmospheres.  The {Spitzer} studies extend from many hot Jupiters, to the hot Neptune orbiting GJ\\,436. Here I review the current status of {\\it Spitzer} secondary eclipse observations, and summarize the results from the viewpoint of what is robust, what needs more work, and what the observations are telling us about the physical nature of exoplanet atmospheres. ",
        "introduction": "The powerful astrophysical leverage provided by transits enables us to study extrasolar planets directly, i.e., by detection of their emergent radiation.  The {\\it Spitzer Space Telescope} has provided the bulk of these detections.  The first {\\it Spitzer} measurements of exoplanet secondary eclipses were announced in 2005.  Two independent groups (Charbonneau et al. 2005; Deming et al. 2005) measured eclipses for two different planets, using two different {\\it Spitzer} instruments, and obtained very similar results (Figure~1).  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=2.8in]{f1.eps} \\caption{First detections of exoplanet thermal emission using the {\\it Spitzer Space Telescope}. Plotted are the secondary eclipses of TrES-1 at 8\\,$\\mu$m (top, Charbonneau et al. 2005), and HD\\,209458b at 24\\,$\\mu$m (bottom, Deming et al. 2005)}. \\label{fig1} \\end{center} \\end{figure}  Since each eclipse was independently measured to $\\sim 6\\sigma$ significance, exoplanet thermal emission was securely detected.  The discovery of transits in HD\\,189733b (Bouchy et al. 2005) provided an opportunity to measure exoplanet thermal emission at higher signal-to-noise ratio.  Initial observations of HD\\,189733b at 16\\,$\\mu$m (Deming et al. 2006) showed an eclipse of the planet at 32$\\sigma$ significance (Figure~2), and subsequent work using the IRAC instrument has detetcted the planet's flux to 60$\\sigma$ precision at 8\\,$\\mu$m (Knutson et al. 2007).  This extraordinary level of precision in measuring exoplanet thermal emission allows many intersting studies that could hardly have been imagined when the first extrasolar planets were detected by radial velocity studies.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=2.2in]{f2.eps} \\caption{Eclipse of HD\\,189733b at 16\\,$\\mu$m (Deming et al. 2006)}. \\label{fig2} \\end{center} \\end{figure}  In this review I summarize highlights from {\\it Spitzer} secondary eclipse measurements, with some discussion of transmission spectroscopy during transit. The quality of work in this field has been uniformly high, but I will summarize the results from the viewpoint of what is robust, and what I believe needs more work and clarification. ",
        "conclusions": ""
    },
    "0808/0808.0070_arXiv.txt": {
        "abstract": "{ The growth of supermassive black holes (SMBH) through accretion is accompanied by the release of enormous amounts of energy which can either be radiated away, as happens in quasars, advected into the black hole, or disposed of in kinetic form through powerful jets, as is observed, for example, in radio galaxies. Here, I will present new constraints on the evolution of the SMBH mass function and Eddington ratio distribution, obtained from a study of AGN luminosity functions aimed at accounting for both radiative and kinetic energy output of AGN in a systematic way. First, I discuss how a refined Soltan argument leads to joint constraints on the mass-weighted average spin of SMBH and of the total mass density of high redshift ($z\\sim 5$) and ``wandering'' black holes. Then, I will show how to describe the ``downsizing'' trend observed in the  AGN population in terms of cosmological evolution of physical quantities (black hole mass, accretion rate, radiative and kinetic energy output). Finally, the redshift evolution of the AGN kinetic feedback will be briefly discussed and compared with the radiative output of the evolving SMBH population, thus providing a robust physical framework for phenomenological models of AGN feedback within structure formation. ",
        "introduction": "Black holes in the local universe come into two main families according to their size, as recognized by the strongly bi-modal distribution of the local black hole mass function (see Fig.~\\ref{fig:mf}). While the height, width and exact mass scale of the stellar mass peak should be understood as a by-product of stellar (and binary) evolution, and of the physical processes that make supernovae and gamma-ray bursts explode, the supermassive black holes one is the outcome of the cosmological growth of structures and of the evolution of accretion in the nuclei of galaxies, likely modulated by the mergers these nuclear black holes will experience as a result of the hierarchical galaxy-galaxy coalescences.  \\begin{figure*}[t!] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{shankar_mf.ps}} \\caption{ \\footnotesize The local black hole mass function, plotted as $M \\times \\phi_M$, in order to highlight the location and height of the two main peaks. The stellar mass black holes peak has been drawn assuming a log-normal distribution with mean mass equal to 5 solar masses, width of 0.1 dex and a normalization yielding a density of about $1.1 \\times 10^{7}$ $M_{\\odot}$ Mpc$^{-3}$ \\citep{fukugita:04}. The supermassive black hole peak, instead, contribute to an overall density of about $4.3 \\times 10^{5}$ $M_{\\odot}$ Mpc$^{-3}$ \\citep{merloni:08}} \\label{fig:mf} \\end{figure*}  In the recent  literature, it has become customary  to introduce the works on cosmological aspects of AGN astrophysics by referring to the strong role they most likely play in the galaxy formation process throughout cosmic history. Indeed, a new paradigm has emerged, according to which the feedback energy released by growing supermassive black holes (i.e. AGN) limits the stellar mass growth of their host galaxies in a fundamental, generic, but yet not fully understood fashion.  The strongest observational evidence for such a schematic picture emerged in the last decade. The search for the local QSO relics via the study of their dynamical influence on the surrounding stars and gas carried out since the launch of the {\\it Hubble Space Telescope} \\citep[see e.g.][and references therein]{richstone:98,ferrarese:08} led ultimately to the discovery of tight scaling relations between SMBH masses and properties of the host galaxies' bulges \\citep{gebhardt:00,ferrarese:00,tremaine:02,marconi:03}, clearly pointing to an early co-eval stage of SMBH and galaxy growth. A second piece of evidence comes from X-ray observations of galaxy clusters, showing that black holes are able to deposit large amounts of energy into their environment in response to radiative losses of the cluster gas. From studies of the cavities, bubbles and weak shocks generated by the radio emitting jets in the intra-cluster medium (ICM) it appears that AGN are energetically able to balance radiative losses from the ICM in the majority of cases \\citep[see][and references therein]{birzan:08}.   Nevertheless, the physics of AGN heating in galaxy cluster is still not well established, neither have the local scaling relations proved themselves capable to uniquely determine the physical nature of the SMBH-galaxy coupling. As a consequence, a large number of feedback models have so far been proposed which can reasonably well reproduce these relations. From the observational point of view, the crucial test for most models will be a direct comparison with the high-redshift evolution of the scaling relations.  There is, however,  another benchmark, based on existing data, all models have to be tested upon: the evolution of the SMBH mass function and of the predicted energy output (either in radiative or kinetic form) needed to offset gas cooling and star formation in galaxies.   Here I present our recent attempt to reconstruct the history of SMBH accretion in order to follow closely the evolution of the black hole mass function, as needed in order to test various models for SMBH cosmological growth as well as those for the black hole-galaxy co-evolution. Similar to the case of X-ray background synthesis models, where accurate determinations of the XRB intensity and spectral shape, coupled with the resolution of this radiation into individual sources, allow very sensitive tests of how the AGN luminosity and obscuration evolve with redshift, we have argued that accurate determinations of the local SMBH mass density and of the AGN (bolometric) luminosity functions, coupled with accretion models that specify how the observed AGN radiation translates into a black hole growth rate, allow sensitive tests of how the SMBH population (its mass function) evolves with redshift. By analogy, we have named this exercises `AGN synthesis modelling' \\citep{merloni:08}. In performing it, we have taken advantage of the fact that the cosmological evolution of SMBH is markedly simpler than that of their host galaxies, as individual black hole masses can only grow with time, and SMBH do not transform into something else as they grow. Moreover, by identifying active AGN phases with phases of black holes growth, we can follow the evolution of the population by solving a simple continuity equation, where the mass function of SMBH at any given time can be used to predict that at any other time, provided the distribution of accretion rates as a function of black hole mass is known (see \\S~\\ref{sec:integral}).  \\begin{figure*}[t!] \\begin{tabular}{cc} \\resizebox{0.48\\hsize}{!}{\\includegraphics[clip=true]{rhobh_mass_083.ps}}& \\resizebox{0.48\\hsize}{!}{\\includegraphics[clip=true]{zev_lambda_083_masses.ps}}\\\\ \\end{tabular} \\caption{\\footnotesize Redshift evolution of the SMBH mass density (left) and average Eddington ratio (right), calculated for different BH mass bins (in solar mass units). In the right hand plot, black lines and grey shaded area represent the overall (mass-wighted) average Eddington ratio.} \\label{fig:zev_mdot} \\end{figure*}  In order to carry out our calculation, we assumed that black holes accrete in just three distinct physical states, or ``modes'': a radiatively inefficient, kinetically dominated mode at low Eddington ratios (LK, the so-called ``radio mode'' of the recent literature), and two modes at high Eddington ratios: a purely radiative one (radio quiet, HR), and a kinetic (radio loud, HK) mode, with the former outnumbering the latter by about a factor of 10. Such a classification is based on our current knowledge of state transitions in stellar mass black hole X-ray binaries as well as on a substantial body of works on scaling relations in nearby SMBH. It allows a relatively simple mapping of the observed luminosities (radio cores, X-ray and/or bolometric) into the physical quantities related to any growing black hole: its accretion rate and the released kinetic power.  In this work I will focus on just a few specific aspects of the derived SMBH evolution, in particular on the redshift evolution of the mass function and Eddington ratio distribution, on the constrains we put on the mass-weighted average spin of the SMBH population, and on the kinetic energy output of growing black holes. A more detailed discussion of the methodology, as well as a wider exploration of our results can be found in \\citet{merloni:08}.  All results will be shown accounting for the intrinsic uncertainties of the adopted luminosity functions. We estimated that these uncertainties can be evaluated by comparing different analytic parametrization of the same data sets; specifically, we adopted the LDDE and MPLE parametrization for the hard X-ray luminosity function of \\citet{silverman:08}, and two alternative parametrizations for the flat-spectrum radio luminosity function of \\citet{dunlop:90} and \\citet{dezotti:05}. ",
        "conclusions": "I have outlined some recent results of our work aimed at pinning down as accurately as possible the cosmological evolution of active galactic nuclei and of the associated growth of the supermassive black holes population.  In particular, I have focused here on the global (integrated) constraints on the mass-weighted average spin of SMBH, and I have discussed in some details the specific ways in which these are tighten to the very interesting open issues regarding the population of high-redshift SMBH and that of black holes ejected from galactic nuclei due to gravitational wave recoil in merger events.  I have also discussed a few generic properties of the kinetic energy output of growing black holes, emphasizing the importance of a late phase of low-luminosity, jet-dominated accretion onto the most massive objects.  The richness of details we have been able to unveil demonstrates that times are ripe for comprehensive unified approaches to the multi-wavelength AGN phenomenology. At the same time, our results should serve as a stimulus for semi-analytic and numerical modelers of structure formation in the Universe to consider more detailed physical models for the evolution of the black hole population."
    },
    "0808/0808.0300_arXiv.txt": {
        "abstract": "In this paper we introduce a new linear filtering technique, the so-called matrix filters, that maximizes the signal-to-interference ratio of compact sources of unknown intensity embedded in a set of images by taking into account the cross-correlations between the different channels. By construction, the new filtering technique outperforms (or at least equals) the standard matched filter applied on individual images. An immediate application is the detection of extragalactic point sources in Cosmic Microwave Background images obtained at different wavelengths. We test the new technique in two simulated cases: a simple two-channel case with ideal correlated color noise and more realistic simulations of the sky as it will be observed by the LFI instrument of the upcoming ESA's Planck mission. In both cases we observe an improvement with respect to the standard matched filter in terms of signal-to-noise interference, number of detections and number of false alarms. ",
        "introduction": "\\IEEEPARstart{T}{he} detection of faint pointlike sources is a task that is common to many branches of Astronomy, from the search for protostars in gas-rich nebulae to the study of active galactic nuclei in the confines of the observable universe. Since the angular size of these objects is smaller than the angular resolution of the telescopes that are used to observe them, they appear as \\emph{point sources} with the shape of the telescope point spread function.  A case of particular interest is the detection of extragalactic point sources (EPS) in microwave wavelengths. In the microwave range of the electromagnetic spectrum, the sky is dominated by the so-called Cosmic Microwave Background (CMB) radiation, a relic of the hot and dense first moments of the universe. The study of the CMB is one of the hottest research topics in modern Cosmology. For a short review on the CMB, see~\\cite{cmbrev06}. The CMB signal is mixed with other signals of astrophysical origin, mainly the emission from our own Galaxy and from a large number of extragalactic objects including radio galaxies, dusty galaxies and galaxy clusters. For the typical angular resolution of current CMB experiments, ranging from a few arcmin to one degree, most of these extragalactic objects appear as point sources. From the standpoint of CMB, these point sources are contaminants that must be removed; from the standpoint of extragalactic Astronomy, however, they provide a valuable source of information, particularly at microwave wavelengths where the properties of sources are poorly understood~\\cite{bluebook}. In both cases, techniques for the detection of faint point sources are needed.  For single images taken at a given wavelength the problem is equivalent to the general problem of detecting a number of objects, all of them with a known waveform but unknown positions and intensities, embedded in additive noise (not necessarily white). In the field of microwave Astronomy, wavelet techniques~\\cite{vielva01,vielva03,MHW206,wsphere,NEWPS07}, Bayesian approaches~\\cite{hobson03,powell08}, matched filters~\\cite{tegmark98,barr03,can06} and other related linear filtering techniques~\\cite{sanz01,naselsky02,herr02b,herr02c,can04a,can05a,can05b} have proved to be useful. The common feature of all these techniques is that they rely on the prior knowledge that the sources have a distinctive spatial behaviour (i.e.~a known spatial profile, plus the fact that they appear as compact objects as opposed to `diffuse' random fields) that helps to distinguish them from the noise.  Most of the current CMB experiments are able to observe the sky at several different wavelength bands. In particular, the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{wmap0,hinshaw07} is observing the sky at 23, 33, 41, 61 and 94 GHz and the upcoming ESA's Planck satellite~\\cite{planck_tauber05} will observe the sky in nine frequency channels ranging from 30 to 857 GHz. Multi-wavelength data can be used to improve the detection/separation of astrophysical components. For example, CMB can be separated from Galactic dust and synchrotron emission using well-established methods; a non-exhaustive list of them include Independent Component Analysis~\\cite{fastica02,fastica07}, Maximum Entropy Method~\\cite{MEM97,vlad02,barreiro_MEM04}, Internal Linear Combination~\\cite{wmap0,eriksen04} and Wiener filtering~\\cite{bouchet99}, among others. Moreover, specific techniques for the detection of compact sources whose spectral behaviour is well known have been proposed in the literature; a typical example is the detection of the so-called Sunyaev-Zel'dovich effect due to galaxy clusters~\\cite{herr02a,chema02,schaf06,melin06}.  All the previous component separation techniques, however, are in trouble when dealing with extragalactic point sources. EPS form a very heterogeneous population constituted by a large number of objects with very different physical properties, from radio-emitting active galactic nuclei to dusty star-forming galaxies. Since each source is a unique object, with an spectral emission law that is different from the spectral emission law of any other, the component separation problem is underdetermined: the number $N$ of components to be separated is much larger than the number $M$ of channels available. Attempts to simplify the problem by grouping the extragalactic sources into classes of objects with similar spectral behaviour are in most cases unsatisfactory.  So far we have considered two approaches: on the one hand, it is still possible to work channel by channel, separately, by using the filtering techniques above mentioned. But in that case a valuable fraction of the information that multi-wavelength experiments can offer is wasted. On the other hand, standard multi-wavelength component separation techniques are impracticable because the number of physical components involved is too high. An intermediate approach is to design filters that are able to find compact sources thanks to their distinctive spatial behaviour while at the same time do incorporate some multi-wavelength information, without pretending to achieve a full component separation. In this paper we propose a new filtering technique that makes no assumptions about the spectral behaviour of the sources, but that makes use of some multi-wavelength considerations, namely: \\begin{itemize} \\item When a source is found in one channel, it will be also present in the same position in all the other channels. \\item The spatial profile of the sources may differ from channel to channel, but it is a priori known. For example, for a microwave experiment the source profile is equal to the antenna response of the experiment's radiometers\\footnote{In general, of its detectors. Some microwave experiments use bolometers instead of radiometers, but the distinction is irrelevant for the purposes our discussion.}, that is well known. \\item The second order statistics of the background in which the sources are embedded is well known or it can be directly estimated from the data by assuming that point sources are sparse. \\end{itemize} \\noindent We find that the new technique we propose takes the form of a matrix of filters that can be applied to any number $M$ of channels. In the case of a single channel, the matrix of filters defaults to the standard matched filter. In section~\\ref{sec:math} we will introduce the new formalism. In section~\\ref{sec:tests} we will illustrate the new method with two different tests: a simple toy model and a more realistic simulation that corresponds to a small region of the sky as it will be observed by the upcoming Planck mission. Finally, in section~\\ref{sec:conclusions} we will draw some conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In this work we have introduced a new filtering technique, the matrix filters, that can be of utility for the detection of extragalactic point sources in multi-wavelength experiments of the Cosmic Microwave Background. Matrix filters are designed to maximize the signal-to-interference ratio of compact sources embedded in a set of images (``channels'') by taking into account the cross-correlations between the different channels. For the case of a single-channel experiment and circularly symmetric sources, the matrix of filters defaults to the standard matched filter. For the case of $M$ totally uncorrelated channels and circularly symmetric sources, the matrix of filters becomes a diagonal matrix whose non-zero elements are the matched filters corresponding to each channel. In a general case, the non diagonal elements of the matrix are non zero.  Since the matrix filters contain as a particular case the matched filter and they are designed to maximize the signal-to-interference ratio, then the individual amplifications for each channel (defining by amplification the quotient between the signal-to-interference ratios after and before filtering) are greater or (in the worst case) equal to the amplifications obtained by the standard matched filters.  We have tested the matrix filters in two simulated cases: a simplistic two-channel case with ideal noise and a more realistic simulation of the sky as it will be observed by the LFI instrument of the upcoming ESA's Planck mission. In the first test we observe that matrix filters clearly outperform the standard matched filter for one of the channels, while for the other channel their performance is very similar to that of the matched filter. Other simulation tests with different parameters we have carried out show a similar behaviour: there is a significant gain for at least one of the channels with respect to standard matched filters.  For the second test case we have considered one single patch of the sky, observed at 30, 44 and 70 GHz. The purpose of this test was just to give an example of how the matrix filters work in a more realistic case. In this example, matrix filters outperform the matched filter both in the flux limit they can reach and in the ratio between false alarms and true detections. This result seems to indicate that matrix filters could help to obtain better catalogues of extragalactic point sources in future CMB experiments. However, this result has been obtained for a single simulation of a small portion of the sky, and therefore it may not be extrapolable to the whole sky or to any other experiment. The study of the application of matrix filters to the whole microwave sky for the Planck mission is the subject of a future work."
    },
    "0808/0808.2135.txt": {
        "abstract": "We present a time-dependent  multi-zone code for simulating the variability of Synchrotron-Self Compton (SSC) sources. The code adopts a  multi-zone pipe geometry for the emission region, appropriate for simulating emission from a standing or propagating shock in a collimated jet. Variations in the injection of relativistic electrons in the inlet propagate along the length of  the pipe cooling radiatively. Our  code for the first time takes into account the non-local, time-retarded nature of synchrotron self-Compton (SSC) losses that are thought to be dominant in  TeV blazars. The observed synchrotron and SSC emission is followed self-consistently  taking into account light travel time delays. At any given time, the emitting  portion of the pipe depends on the frequency and the nature of the  variation followed. Our simulation employs only one additional physical parameter relative to one-zone models, that of the pipe length and is computationally very efficient, using  simplified expressions for the SSC processes. The code will be useful for observers modeling GLAST,  TeV, and X-ray observations of SSC blazars. ",
        "introduction": "}   In blazars, radio loud active galaxies with their relativistic jets pointing close to our  line of sight  \\citep{blandford78}, the observed radiation is dominated by  relativistically beamed emission from the sub-pc base  of the jet. The blazar spectral energy distribution (SED)  consists of two components. The first one, peaking at sub-mm to X-ray energies is almost certainly due to synchrotron radiation, while the second one peaking at MeV to TeV energies is believed to be of inverse Compton (IC) nature, with both components produced by  the same population of relativistic electrons. The nature of the IC-scattered seed photons   is still not clear, with both external optical-UV photons from the broad line region \\citep{sikora94} and IR photons from the putative molecular torus \\citep{blazejowski00}, as well as  synchrotron photons (SSC, e.g. Maraschi, Ghisellini \\& Celotti 1992) contributing. It is believed that  in the case of powerful blazars peaking at MeV to GeV energies, external seed photons  from the broad line region dominate  the IC scattering, while for weaker lineless  blazars, peaking at $\\sim$ TeV energies, SSC is the dominant emission mechanism. Recent observational results (e.g. D'Arcangelo et al. 2007; Marscher et al. 2008), however, place the blazar emission site beyond the broad line region, lending support  to the possibility that even in powerful blazars the GeV emission process  may be pure SSC. For a review of  leptonic models, as well as hadronic models for blazar emission (e.g.  Aharonian 2000)   see \\citet{boettcher06}.    Due to the small angular size of the blazar emission region,  it is not possible to spatially resolve the emitting region. Because of this, information about the structure of the emitting source can be obtained only through  multiwavelength variability studies. Particularly telling is the variability of the emission produced by  the highest energy electrons, because these electrons lose energy very quickly and  exist only close to the sites where they have been produced. The goal of multiwavelength variability  campaigns, involving in many cases observations from radio up to TeV energies,  is to study the  characteristics of blazar variability, such as correlations and/or  time delays between  different energies, spectral characteristics of the observed variability,  and the amplitude of  variability as a function of energy.   Most notable amongst blazars are the so called TeV blazars  for which the synchrotron emission peaks at X-ray energies and the SSC emission peaks at TeV energies, as they present the active galaxies  producing the highest confirmed electron energies. Variations of TeV blazars in these two bands can be extremely  rapid (TeV doubling times as short as a few min; Aharonian et al. 2007), suggesting highly relativistic sub-pc scale  flows (Doppler factors $\\delta\\sim 50$;  e.g. Begelman, Fabian \\& Rees 2008) that decelerate substantially  \\citep{georganopoulos03,ghisellini05} to match the much slower speeds required by VLBI observations \\citep{piner04, piner08}.  The TeV and X-ray variations are usually well correlated (e.g. Fossati et al. 2008; Maraschi et al. 1999; Sambruna et al. 2000),  as expected, because they present variations by the same electron population. Usually, the lower energy emission  within each of these bands peaks with small time delays relative to the higher energy emission (e.g. Fossati et al. 2000), while  the X-ray and TeV spectra become harder with increasing flux (e.g.  Takahashi et al. 1996). In certain  cases, however, the X-ray and TeV variability do not seem to be correlated in a simple way (e.g. Aharonian et al. 2005). An intriguing  variability pattern is that of the so-called `orphan' flares,  rare TeV flares  that are not accompanied by X-ray  flares (e.g. Krawczynski et al. 2004,  B\\l a\\.zejowski et. al. 2005). While the correlated X-ray - TeV flares can be understood through an increase of the high energy emitting electrons, orphan TeV flares defy such a straightforward explanation.   Models of  blazar emission to date have, for the most part,  been in some form of homogeneous one zone models (e.g. Mastichiadis \\& Kirk 1997; Krawczynski,  Coppi, \\&  Aharonian  2002).  Such models, although  appropriate for modeling the steady-state emission of a source, cannot simulate variability faster than the zone light crossing time. The basic limitation of one zone models stems from the fact that  the high energy variability of both the synchrotron and SSC components is produced by high energy electrons with cooling times shorter than the light crossing time. Even if  we assume that a disturbance in the radiating plasma (e.g. a higher density)  instantaneously propagates across the zone,  the received radiation would be smeared out for timescales shorter than the light crossing time, due to light travel time delays from different parts of the source (\\S \\ref{section:testing}; also Chiaberge \\& Ghisellini 1999), and no variability faster than the light crossing time would be observed. One, therefore, cannot use one zone models to infer the source structure from high energy variability.  Inhomogeneous variability models of increasing degree of sophistication have attempted to overcome the problems of one-zone models. The basic idea is to overcome the unphysical instantaneous injection throughout the source by adopting a specific geometry for the plasma flow that includes an inlet for injecting the radiating plasma. Variations in the injected plasma  propagate and produce variations in the emissivity. Calculations of  the received emission that take into account the light-travel times that radiation from different parts of the source  takes  to reach the observer, produce light curves that, at least,  do not violate causality. How physically realistic these light curves are depends on the approximations used and on the  characteristics of the source to be modeled. For example,  synchrotron and IC losses from photons external to the source are local processes in the sense that, at a given point in the flow, the energy loss rate only depends on the local magnetic field and external photon field energy density, and not on the photon production throughout the  source.  This is not the case with SSC losses, because synchrotron photons produced throughout the source at earlier  times - to take into account the light travel time from one point of the source to another - contribute to the photon energy density responsible for the SSC losses and to the emissivity at a given point and time in the source. To properly model sources like TeV blazars, in which SSC losses are important or even dominant, these considerations have to be taken into account.   There have been a few  attempts during the last fifteen years to take these spatial considerations into account.  \\cite{gomez94} considered a conical jet  with a  constant bulk Lorentz factor flow in which the electron plasma and the magnetic field undergo adiabatic evolution only  and calculated the  radio variability induced by a shock wave propagating along the jet. \\cite{georganopoulos98a,georganopoulos98b} studied a parabolic jet that hydrodynamically accelerates and focuses  to a conical geometry, and  by following the synchrotron energy losses of the emitting electrons  reproduced the  radio to X-ray  light curves  of the X-ray bright blazar PKS 2155-304. This resulted in a frequency-dependent source size, in agreement with the fact that the variability timescale of synchrotron radiation increases with decreasing frequency. It also reproduced the usually observed soft lags (variations at soft X-rays being preceded by variations at hard X-rays) and the counterclockwise X-ray flux - X-ray spectral index loops (e.g. Takahashi 1996, Maraschi 1999, Kataoka 2000, Ravasio 2004), both manifestations of radiative cooling dominating the energetics of the high energy  electrons.    \\cite{kirk98} developed a semi-analytical model  in which low energy electrons are injected in a zone where they undergo acceleration and eventually escape. The acceleration zone is assumed to move with a certain velocity, leaving behind the freshly accelerated electrons that cool through synchrotron radiation.  Variations in the injection rate of low energy electrons in the acceleration zone result in variations of the emissivity, which are integrated over the volume of the source, taking into account time delays, to produce the observed multifrequency synchrotron light curves. This model includes a treatment of particle acceleration and it is able to reproduce the uncommon hard lags (variations at hard X-rays  preceded by variations at soft X-rays) and clockwise  X-ray flux - X-ray spectral index loops \\citep{zhang02,ravasio04}, both manifestations of electrons still accelerating, just before  reaching the maximum electron Lorentz factor, where the acceleration and radiative loss timescales are comparable. It does not include, however, SSC considerations.  \\cite{chiaberge99} studied  the synchrotron and SSC emission,  from a homogeneous one zone model in which they assumed an instantaneous plasma injection, but taking into account the time delays with which the external observer would observe the variability (a similar approach was also taken by Kataoka et al. 2000). They also studied a case similar to that of Kirk et al. (1998), but without treating particle acceleration, by splitting the source into smaller one zone models that evolved autonomously, in the sense that ({\\sl i}) the SSC emission inside any one of their single zones uses as seed photons  only the synchrotron photons produced in that zone and ({\\sl ii}) the SSC energy losses in every zone are caused only by the synchrotron photons produced in that zone.  This simplified approach is a good approximation for following the energetics of the electrons if the source is synchrotron dominated (because the SSC losses, although inappropriately calculated, are negligible), but does not produce realistic SSC light curves, because it does not calculate the emission due to upscattering synchrotron photons produced in other parts of the source in retarded times.    A significant improvement was introduced by \\cite{sokolov04} who incorporated in the calculation of the SSC emission from a given location in an inhomogeneous source the synchrotron photons produced throughout the source in retarded times. This produces accurate SSC light curves, provided that the SSC losses  that were still treated as a local process are negligible. In a follow up paper, \\cite{sokolov05},  considered also external Compton photons from the broad line region and the molecular torus. The challenge for inhomogeneous multi zone models  for sources such as  the TeV emitting blazars is  the calculation of the non-local,  time-retarded SSC losses induced by photons produced in other parts of the source.  Here,  we present such an inhomogeneous model that, for the first time, takes into account the non-local, time-delayed  source emission on  the SSC losses. We assume that a power law of relativistic electrons is injected at the inlet  of a pipe, and that the electrons flow downstream and cool radiatively. Variations in the injected electron distribution propagate downstream and manifest themselves as frequency dependent variability. This allows us to model high energy multiwavelength variability in a self-consistent manner. In \\S \\ref{section:onezone} we describe the one zone model, which we use  as a building block for the multizone model, and we show that, by construction, one zone models  cannot simulate variability produced by high energy electrons with radiative cooling time shorter than the electron light crossing time from the single zone. In \\S \\ref{section:multizone} we describe our multizone, pipe-geometry model with emphasis on the coupling between subsequent zones and on the  calculation of the local photon field due to non-local, time-delayed emission throughout the source. This is followed in \\S \\ref{section:results} by a comparison of the code with analytical results and  a series of case studies. We conclude  in \\S \\ref{section:discussion} with a discussion of additional considerations that can be used as starting points for future work. ",
        "conclusions": "} We presented  a multizone code that for the first time takes into account the  non-local, time-retarded nature of SSC losses. This  code is currently the only multizone model that incorporates the non-local, time-delayed  SSC  losses, and as such is uniquely suitable for modeling the results of  multiwavelength campaigns  at radio, optical, X-ray and $\\gamma$-ray energies, with the additional constraints in the critical and unexplored  for TeV blazars GeV  GLAST regime.    As we argued, the results of one zone codes for the critical high energy regime of both the synchrotron and SSC components are problematic, and should not be used to infer the physical conditions in the source through variability modeling.  We described our multizone code,  tested it successfully against known analytical results, and presented a small number of variability case studies. The case studies we presented, although based on the same underlying steady-state configuration, exhibited very different variability patterns. This means that detailed modeling of broadband SEDs and simultaneous multiwavelength variability can be used to infer what is actually the cause of a given observed variability pattern, providing reliable constraints on the particle acceleration taking place.  Orphan flares can be reproduced assuming an increase of the injection of the low energy electrons, but not assuming  the injection of a very high energy electron population, as we also showed analytically.  The fact that this plausible variation  cannot produce orphan flares significantly   narrows the   parameter space for events that can produce such events, possibly in agreement with their observed scarcity.  The  code we described  can  run  with a typical workstation in a reasonable time of at most a few minutes at a resolution of $\\sim 10$ bins per decade of observing frequency, $\\sim 10$ bins per decade of electron energy, and $\\sim 50$ zones. To achieve this we  employed a pipe geometry, and  adopted an energy conserving $\\delta$-function approximation for the  SSC  emissivity, as well as a step function approach to take into account the change from the Thomson to Klein-Nishina IC scattering cross-sections. Adopting  these approximations is problematic  for situations where IC scattering of narrow photon distributions (e.g. line emission from the broad line region or even a blackbody spectrum characterized by a typical photon energy $\\epsilon_0$) is important. In this case the adoption of the step function cross section description would create a strong artificial feature on the EED localized at the transition from the Thomson to the Klein-Nishina regimes at $\\gamma\\propto 1/\\epsilon_0$, which would then propagate to the emitted spectra through the $\\delta$-function IC emissivity. For SSC systems, however, where  the seed photons are spread over many decades in energy, the resulting spectra are good approximations of those   produced using the full expressions for the synchrotron and SSC emissivities as well as the full Klein-Nishina cross section.  Including the above considerations, as well as the processes of synchrotron opacity and pair production through $\\gamma$-ray absorption  within the source, would  increase the execution time up to levels marginally comfortable for the typical workstation. Most probably, such an extension of the code would require parallelization.  A more desirable upgrade of the code would drop the assumption of no lateral gradients in  the plasma characteristics by  switching to a two-dimensional geometry, in which the electron distribution and the SSC photon energy density are allowed to change laterally to the flow direction. Such considerations may be relevant to the recently observed $\\sim   0.75 $ days  delay between the IR and the X-ray variability in 3C 273 \\citep{mchardy07}.    These authors argued that the delay may be attributed to the time it takes for the SSC photon energy density to built up as the SSC photons are transversing   the cross section of the flow. This upgrade will scale the computation time roughly by $N^2$, increasing it  from $\\sim$  few minutes to  $\\sim$  several hours. We note here that our formalism can be extended to  treat velocity profiles in term of the decelerating flow \\citep{georganopoulos03} or the spine sheath model \\citep{ghisellini05} that have developed to address the lack of superluminal motions in TeV blazars (e.g. Piner \\& Edwards 2004).  Another upgrade that can be incorporated in the existing code, this time with a minimal computational overhead, is that of a zone for particle acceleration, following the formalism of \\cite{kirk98}. In this case, in the first zone of the model, low energy  electrons will be injected and allowed to accelerate while suffering radiative losses due to synchrotron and non-local SSC. These particles will subsequently escape into the pipe and flow downstream. This configuration will require a different numerical scheme for the acceleration zone, since there most particles are advected upward in energy space, but there is a possibility, in a time-dependent scenario, of the highest energy particles being advected downward, while the rest of the electrons are still advected upwards. The benefit of including particle acceleration in the code is that it will allow us to study cases of hard lags/ counterclockwise loops in the X-ray hardness - X-ray flux diagrams thought to result when acceleration and loss timescales are comparable. Such a  code could  be used to model the observed curved X-ray spectra of high peak frequency blazars in the framework of episodic particle acceleration \\citep{perlman05}."
    },
    "0808/0808.2582_arXiv.txt": {
        "abstract": "The muon charge ratio of ultrahigh energy cosmic rays may provide information to detect the composition of the primary cosmic rays. We propose to extract the charge information of high energy muons in very inclined extensive air showers by analyzing their relative lateral positions in the shower transverse plane. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.2061_arXiv.txt": {
        "abstract": "Firmani et al. proposed a new Gamma Ray Burst (GRB) luminosity relation that showed a significant improvement over the $L_{iso} - E_{peak}$ relation.  ($L_{iso}$ is the isotropic peak luminosity and $E_{peak}$ is the photon energy of the spectral peak for the burst.) The new proposed relation simply modifies the $E_{peak}$ value by multiplying it by a power of $T_{0.45}$, where $T_{0.45}$ is a particular measure of the GRB duration.  We begin by reproducing the results of Firmani for his 19 bursts.  We then test the Firmani relation for the same 19 bursts except that we use independently measured values for $L_{iso}$, $T_{0.45}$, and $E_{peak}$, and we find that the relation deteriorates substantially.  We further test the relation by using 60 GRBs with measured spectroscopic redshifts, and find a relation that has a comparable scatter as the original $L_{iso} - E_{peak}$ relation.  That is, a much larger sample of bursts does not reproduce the small scatter as reported by Firmani et al.  Finally, we investigate whether the Firmani relation is improved by the use of any of 32 measures of duration (e.g., $T_{90}$, $T_{50}$, $T_{90}/N_{peak}$, the fluence divided by the peak flux, $T_{0.30}$, and $T_{0.60}$) in place of $T_{0.45}$.  The quality of each alternative duration measure is evaluated with the root mean square of the scatter between the observed and fitted logarithmic $L_{iso}$ values.  Although we find some durations yield slightly better results than $T_{0.45}$, the differences between the duration measures are minimal. We find that the addition of a duration does not add any significant improvement to the $L_{iso} - E_{peak}$ relation.  We also present a simple and direct derivation of the Firmani relation from {\\it both} the $L_{iso} - E_{peak}$ and Amati relations.  In all we conclude that the Firmani relation neither has an independent existence nor does it provide any significant improvement on previously known relations that are simpler. ",
        "introduction": "To date, eight separate luminosity relations have been identified for long duration Gamma-Ray Bursts (GRBs) (Schaefer and Collazzi, 2007). These relations correlate a burst's peak bolometric luminosity with various light curve and spectral parameters. The possible utility of these relations is that we can then use GRBs as tracers of the high-redshift universe.  One of these eight relations is the $L_{iso} - E_{peak}T_{0.45}$ relation proposed by Firmani et al. (2006, hereafter named the Firmani relation). Here, $E_{peak}$ describes the peak of the $E*F(E)$ curve (proportional to $\\nu F_{\\nu}$ ), which is the photon energy of the peak spectral flux; $L_{iso}$ is the isotropic luminosity of the burst measured bolometrically (1-10,000 kev in the burst rest frame). $T_{0.45}$ is the Reichart definition of a GRB time duration (Reichart et al. 2001) where the duration is the total time interval of the brightest bins in the light curve that contains 45$\\%$ of the burst fluence. The Firmani relation was presented as an improvement over the $L_{iso} - E_{peak}$ relation. Nineteen GRBs were used to demonstrate a tight correlation with a reduced chi-square of 0.7 over 16 degrees of freedom; the resulting luminosity relation being $L_{iso} \\propto {E_{peak}}^{1.62} {T_{0.45}}^{-0.49}$.  The reported scatter in the Firmani relation is substantially smaller than those of most other GRB luminosity relations.  This result offers the hope of substantial improvement in the accuracy of GRBs for cosmological distance measures.  We have no physical reason to expect that the addition of a duration should make for a tighter relation. Nonetheless, we should look to see if we can get tighter luminosity relations from duration definitions other than $T_{0.45}$, as $T_{0.45}$ may not be the optimal duration to use.  We have no physical reason to expect that any one definition of duration is best, while the particular choice of the Reichart definition was used only for historical reasons no longer of any relevance.  For example, we could still use the Reichart definition, but measure the duration over a different percentage of the burst fluence.  So perhaps the use of a duration based on 30\\% or 60\\% ($T_{0.30}$ or $T_{0.60}$) might be better.  Alternative definitions of duration can be considered instead of the Reichart formulation. For example, we can try to define the duration as equalling the total fluence divided by the peak flux to get a sort of equivalent width; alternatively, we could use the familiar $T_{90}$ or $T_{50}$ durations.  A wide variety of alternative durations can be defined, and we do not know which one will be optimal.  In this Section 2, we first reproduce the Firmani relation for the original data of Firmani et al. (2006) as a test that we are using identical fitting procedures. Then we test the Firmani relation with a set of independent data for the same 19 bursts.  A further test of the Firmani relation is made with a much larger sample of 60 bursts. In Section 3, we test many duration definitions in a Firmani-like relation to see which produces the `tightest' correlation. Section 4 contains a simple and forced derivation of the Firmani relation from two other luminosity relations. ",
        "conclusions": "In a recent independent preprint, Rossi et al. (2008) also examined the Firmani relation, in particular with a comparison to the Amati relation.  They use an extended sample of 40 {\\it BeppoSAX} and {\\it Swift} bursts, with little overlap with our sample of 60 GRBs.  Their best fit is somewhat different from those in Eqs 1, 4, or 7; with their fitted Firmani relation scaling as $L_{iso} \\propto E_{peak}^2 T_{0.45}^{-1}$.  They realized that this Firmani relation is essentially identical to the Amati relation (Amati et al. 2002), which gives the isotropic energy emitted in gamma radiation over the whole burst duration as $E_{\\gamma ,iso} \\propto E_{peak}^2$.  With the reasonable approximation that the total energy in the light curve equals the peak luminosity times the duration ($E_{\\gamma ,iso} \\approx L_{iso} T_{0.45}$), the Amati relation ($E_{\\gamma ,iso} \\propto E_{peak}^2$) is transformed into their Firmani relation ($L_{iso} \\propto E_{peak}^2 T_{0.45}^{-1}$).  While our exponents in the Firmani relation are somewhat different, the Rossi derivation demonstrates that  the Fermani relation has a physical basis that is close to that of the Amati relation.  Rossi et al. (2008) further go on to show that the scatter in their Firmani relation is comparable to that in the Amati relation, which is another way of saying that the two relations are not independent.  Here, we will directly derive the Firmani relation from {\\it both} the $L_{iso} - E_{peak}$ and Amati relations.  Let us start with the relation $L_{iso} \\propto E_{peak}^{1.68}$ as given in Schaefer (2007).  This can be rearranged as $L_{iso} \\propto E_{peak}^{1.9} (E_{peak}^2/L_{iso})^{-0.69}$.  The Amati relation ($E_{\\gamma ,iso} \\propto E_{peak}^2$) can be inserted to get $L_{iso} \\propto E_{peak}^{1.9} (E_{\\gamma ,iso}/L_{iso})^{-0.69}$.  Now we can select one of our duration definitions with $\\tau = S_{bolo}/P_{bolo}$.  The ratio of fluence to peak flux will equal to the ratio of the burst energy and the peak luminosity, so we have $\\tau = E_{\\gamma ,iso}/L_{iso}$.  This can now be substituted to obtain $L_{iso} \\propto E_{peak}^{1.9} \\tau^{-0.69}$.  We see that we have just derived Eq. 8 with $\\xi = 1.9$ and $\\eta = -0.69$, values which are characteristic of the fitted Firmani relation (cf. Eq. 7).  With this, we see that the Firmani relation has no independent existence because it is only a combination of two simpler relations.  Thus, given any two of these three relations, we can derive the third. A question is which of these is more fundamental. The inherent problem with making this assessment is that it comes down to how one identifies the more fundamental of relations that really address different physics. By Occam's Razor, the more fundamental relation would be one that has the fewest parameters, while accurately and efficiently yielding a luminosity. Thus, the Firmani relation is less fundamental as it uses more parameters for the fit. Of the two relations that remain, the one that has the most `utility' will be the one with the least amount of scatter in its calibration curve.  We were able to reproduce Firmani's results over his small sample of 19 bursts. However, when we substitute independent values for $L_{iso}$, $E_{peak}$ and $T_{0.45}$, we see a substantial broadening around the model. That is, the Firmani relation is not robust on the use of alternative input data. In addition, when we extend the test to a larger sample of 60 bursts, the scatter becomes substantially larger again. Indeed, this scatter is comparable to the scatter in the original $E_{peak} - L_{iso}$ relation.  That is, the Firmani relation is not robust for the use of additional bursts.  These failures of the Firmani relation  have dashed our hopes raised by the tight calibration curves displayed in Firmani et al. (2006).  It also suggests that the addition of a duration does not significantly improve the $L_{iso} - E_{peak}$ relation.  The larger point of interest is that no duration shows a significant advantage over the $E_{peak} - L_{iso}$.  While it might be possible that that a relation involving $\\mathcal{T}_{30}/N_{peak}$ might really have a smaller scatter than the Firmani relation, the improvements are small and not significant. This leads us to conclude that the addition of a duration is not doing enough to improve the $L_{iso} - E_{peak}$ relation to be considered to be a separate luminosity relation.  We conclude that the Firmani relation is not useful for several reasons:  First, the Firmani relation is simply derived by putting together two well-known, simpler, and independent luminosity relations, and thus it has no separate existence.   Second, it is not robust for the inclusion of independent input data or for the extension to many more GRBs.  Third, the real scatter for the Firmani relation does not live up to the hope generated by the original report, with the scatter being comparable to those of the luminosity relations from which it is derived.  In all, we can see no utility or advantage to using the Firmani relation.    We would like to thank the Louisiana State University Board of Regents and the Louisiana Space Consortium for their support."
    },
    "0808/0808.2899_arXiv.txt": {
        "abstract": "We use $16\\;\\rm \\mu m$, Spitzer-IRS, blue peakup photometry of 50 early-type galaxies in the Coma cluster to define the mid-infrared colour-magnitude relation. We compare with recent simple stellar population models that include the mid-infrared emission from the extended, dusty envelopes of evolved stars. The K$_{\\rm s}$-[16] colour in these models is very sensitive to the relative population of dusty Asymptotic Giant Branch (AGB) stars. We find that the \\emph{passively evolving} early-type galaxies define a sequence of approximately constant age ($\\sim 10$~Gyr) with varying metallicity. Several galaxies that lie on the optical/near-infrared colour-magnitude relation do not lie on the mid-infrared relation. This illustrates the sensitivity of the K${\\rm s}$-[16] colour to age. The fact that a colour-magnitude relation is seen in the mid-infrared underlines the extremely passive nature of the majority (68 \\%) of early-type galaxies in the Coma cluster. The corollary of this is that 32\\% of the early-type galaxies in our sample are \\emph{not} `passive', insofar as they are either significantly younger than 10 Gyr or they have had some rejuvenation episode within the last few Gyr. ",
        "introduction": "Early-type galaxies (ellipticals and S0s, ETGs hereafter) are ancient stellar populations whose epoch of formation is related to the large-scale structure formation in the Universe. This picture is supported by various studies that have used optical line-strength indices to determine evolutionary parameters of ETGs in the cluster and field environments (see Renzini 2006 for a review). Among these studies, some authors (Bernardi et al. 2005, Clemens et al. 2006, Thomas et al. 2005, S{\\'a}nchez-Bl{\\'a}zquez et al. 2006a, Annibali et al. 2007) suggest that cluster ETGs have a luminosity weighted, mean stellar age 1-2 Gyr older than those in the field. Recently, Trager, Faber \\& Dressler (2008) challenged this view. Analyzing the Coma cluster, they found that the 12 ETGs in their sample have mean single stellar population equivalent ages of 5-8 Gyr, with the oldest systems being $\\leq$ 10 Gyr old. This average age is remarkably similar to the mean age of ETGs in low density environments (see e.g. Annibali et al. 2007 and references therein).  \\begin{figure*} \\centerline{ \\includegraphics[scale=0.45]{31.ps} \\hskip-3mm \\includegraphics[scale=0.45]{49.ps} \\hskip-3mm \\includegraphics[scale=0.45]{78.ps} \\hskip-3mm \\includegraphics[scale=0.45]{95.ps} } \\vskip-5mm \\centerline{ \\includegraphics[scale=0.45]{129.ps} \\hskip-3mm \\includegraphics[scale=0.45]{143.ps} \\hskip-3mm \\includegraphics[scale=0.45]{144.ps} \\hskip-3mm \\includegraphics[scale=0.45]{148.ps} } \\vskip-5mm \\centerline{ \\includegraphics[scale=0.45]{167.ps} \\hskip-3mm \\includegraphics[scale=0.45]{172.ps} \\hskip-3mm \\includegraphics[scale=0.45]{240.ps} \\hskip-3mm \\includegraphics[scale=0.45]{psf.ps} } \\caption{Example background subtracted, PBCD, IRS peakup images of the Coma ETGs. Contour levels are $2^{n/2}/100$ where n=(-1,0,1,2\\dots) $\\rm mJy/pixel$ with a dashed contour for negative values at the lowest level. The grey-scale in each plot is scaled to the peak source flux. The last panel is an image of the peak-up point spread function. The cross in each panel gives the optical source position. The grey-scale has a square-root stretch and contour levels are $2^{n/2}/100$ where n=(1,3,5\\dots) $\\rm mJy/pixel$. Pixels are {1}\\farcs{8} on a side.} \\label{fig:maps1} \\end{figure*}  The emission from elliptical galaxies longward of $\\sim 10\\;\\rm \\mu m$ has a large contribution from the circumstellar dust around evolved stars, such as those on the AGB. This dust emission has been detected by Spitzer in early-type galaxies where it is seen as a wide emission feature near $10\\;\\rm \\mu m$ with another broad feature near  $18\\;\\rm \\mu m$ (Bressan et al., 2006). The spatial profile in the mid-infrared is similar to the optical profile (Temi et al., 2008, Athey et al. 2002), re-enforcing the notion that the origin is stellar rather than ISM. This cicumstellar dust has also been detected directly as an extended envelope around nearby AGB stars (Gledhill \\& Yates, 2003). Longer wavelength infrared emission from elliptical galaxies is sometimes detected (Leeuw et al., 2004; Marleau et al., 2006; Temi, Brighenti \\& Mathews, 2007.), but this cooler component is probably associated with a diffuse interstellar medium, rather than with evolved stars directly.  For objects older than $0.1\\;\\rm Gyr$ the mid-infrared emission traces stars during their mass-losing AGB phase, and this evolutionary phase is still detected in objects with ages typical of globular clusters such as 47 Tucanae (Lebzelter et al., 2006). Because the luminosity of an AGB star depends on its mass and the main sequence turn-off mass decreases with time, the mid-infrared emission depends on the age of the stellar population. If all the stars in an early-type galaxy were created in an instantaneous burst $10\\;\\rm Gyr$ ago, the mid-infrared emission would be dominated by stars on the AGB with an initial mass of slightly less than $1\\;\\rm M_{\\odot}$ (for solar metallicity).  The strength of the silicate emission from the dusty envelopes of evolved stars is a function of both age and metallicity. As pointed out by Bressan, Granato \\& Silva (1998) however, the dependence of this emission on age and metallicity is different to that in the optical.  A colour-magnitude relation which includes a band containing the silicate emission from AGB stars therefore traces the AGB population as a function of the total luminosity, and can, in principle, be used with the optical relation to disentangle the effects of age and metallicity.  Here, we use $16\\;\\rm \\mu m$ images of 50 ETGs in the Coma cluster, made with the blue peakup detector of the IRS on Spitzer. We combine these with archival Spitzer IRAC images at $4.5\\;\\rm\\mu m$ and K-band images to construct the mid-infrared colour-magnitude relation. As objects at the distance of the Coma cluster were too faint for IRS spectroscopy we include 4 galaxies in the Virgo cluster for which we have IRS spectra. These 4 objects were selected to have spectra that show no evidence of activity, such as recent star formation or an AGN, and are intended as a `template' for passive objects against which to compare the more distant objects of Coma.  The Coma cluster is both very rich and dynamically relaxed, and contains some of the the most massive galaxies in the local Universe. As such, it is an ideal place to study the star formation history of ETGs.    \\begin{figure} \\centerline{ \\includegraphics[scale=0.4]{coma_Re_log.ps} } \\caption{Relation between the effective H-band radii ($1.65\\;\\rm \\mu m$, Gavazzi et al. 2000) and the $16\\;\\rm \\mu m$ effective radii. The latter have been estimated by convolving a model $r^{1/4}$ de Vaucouleur's law with the IRS blue peakup PSF (see text). Crosses are ellipticals and diamonds are lenticulars (classified either as S0/E, S0, SB0, or S0/a). The solid line is a least squares fit to the data whereas the dotted line corresponds to equal Re at both wavelengths. The horizontal dot-dashed line corresponds to the radius containing half of the encircled energy for the IRS blue peakup PSF.} \\label{fig:radii} \\end{figure} ",
        "conclusions": "We present $16\\;\\rm \\mu m$, Spitzer-IRS, blue peakup images of a sample of 50 ETGs in the Coma cluster. We compare these with archival IRAC images at  $4.5\\;\\rm \\mu m$ and 2MASS, $K_s$ band images at $2.2\\;\\rm \\mu m$.  We make the following conclusions.  \\begin{itemize}  \\item{Our IRS blue peakup images show no evidence that the $16\\;\\rm \\mu m$ emission is anything other than stellar in origin}  \\item{The region within a dusty AGB star envelope where dust is hot enough to emit strongly at $16\\;\\rm \\mu m$ is not vulnerable to environmental effects such as the ram-pressure of an ISM wind or dust sputtering by hot gas. Dust grains only spend $\\sim 100\\;\\rm yr$ near $\\sim 300\\;\\rm K$ and this region is sufficiently dense to survive against ram-pressure disruption.}  \\item{We identify the mid-infrared colour-magnitude relation of passively evolving ETGs as the lower envelope of the galaxy distribution in the K$_s$ - [16] vs K$_s$ plane; that is, the minimum value of K$_s$ - [16] at a given K$_s$ magnitude. $\\sim$68\\% of the galaxies in our sample lie on this colour-magnitude relation. These galaxies cannot have had any episode of star formation accounting for more than $\\sim 1\\%$ of the total stellar mass within the last few Gyr. These are genuinely `passively evolving' objects. The remaining objects are either significantly younger than 10 Gyr or have undergone a rejuvenation event in the recent past. This result is at odds with the most recent estimate of the fraction of old objects based on optical spectroscopy (Trager et al. 2008).}   \\item{We construct the mixed optical-NIR-MIR two colour diagram and, by means of updated simple stellar population models, we show that the addition of mid-infrared data allows a much better separation of the effects of age and metallicity, which are rather degenerate in either the optical or mid-infrared when taken in isolation. In this plane galaxies populating the colour-magnitude relation trace a sequence of \\emph{varying metallicity at approximately constant age}. Although this conclusion was already consistent with the optical-NIR colour-magnitude relation, a correlation between age and metallicity could equally well explain the relation. Indeed, comparison with the optical-NIR colours shows that \\emph{a number of galaxies that lie on the optical-NIR relation are significantly displaced from the mid-infrared relation}, with redder K$_s$ - [16] colours. The mid-infrared colour-magnitude diagram therefore shows a sequence of metallicity for old, passive galaxies, with younger objects displaced towards redder K$_s$ - [16] colours.}  \\item{The oldest elliptical galaxies in our sample have luminosity weighted, mean stellar ages of 10.5 Gyr and metallicities within a factor of two of the solar value. No galaxy in our sample is as old or metal poor as the globular cluster 47 Tuc.}  \\item{Although the addition of the K$_s$ - [16] colour allows us to identify objects with significantly younger, luminosity weighted, mean stellar ages, we cannot distinguish between genuinely `young' objects and those that have undergone a minor rejuvenation event. However, given that even a period of recent (last few Gyr) star formation that accounts for less than 1\\% of the total stellar mass will shift a galaxy off the mid-infrared colour-magnitude relation, the latter option seems far more likely. Unambiguous resolution of this issue will require the infrared spectroscopic capabilities of future space observatories.}  \\item{There is evidence that those galaxies that have an excess in the K$_s$ - [16] colour are found preferentially at smaller cluster-centric radii. As the interaction between the dusty AGB star envelopes and the ISM does not directly effect the $16\\;\\rm \\mu m$ emission, the excess may be caused by ``rejuvenation'' episodes induced by the cluster environment.}  \\end{itemize}"
    },
    "0808/0808.3584_arXiv.txt": {
        "abstract": "We estimate cluster ages from lithium depletion in five pre-main-sequence groups found within 100\\,pc of the Sun: TW~Hydrae Association, $\\eta$~Chamaeleontis Cluster, $\\beta$~Pictoris Moving Group, Tucanae-Horologium Association and AB~Doradus Moving Group.  We determine surface gravities, effective temperatures and lithium abundances for over 900 spectra through least squares fitting to model-atmosphere spectra.  For each group, we compare the dependence of lithium abundance on temperature with isochrones from pre-main-sequence evolutionary tracks to obtain model dependent ages. We find that the $\\eta$\\,Chamaelontis Cluster and the TW~Hydrae Association are the youngest, with ages of $12{\\pm6}$\\,Myr and $12{\\pm8}$\\,Myr, respectively, followed by the $\\beta$~Pictoris Moving Group at $21{\\pm9}$\\,Myr, the Tucanae-Horologium Association at $27{\\pm11}$\\,Myr, and the AB Doradus Moving Group at an age of at least $45$\\,Myr (where we can only set a lower limit since the models -- unlike real stars -- do not show much lithium depletion beyond this age).  Here, the ordering is robust, but the precise ages depend on our choice of both atmospheric and evolutionary models.  As a result, while our ages are consistent with estimates based on Hertzsprung-Russell isochrone fitting and dynamical expansion, they are not yet more precise.  Our observations do show that with improved models, much stronger constraints should be feasible: the intrinsic uncertainties, as measured from the scatter between measurements from different spectra of the same star, are very low: around 10\\,K in effective temperature, 0.05\\,dex in surface gravity, and 0.03\\,dex in lithium abundance. ",
        "introduction": "\\label{s:intro}  Triggered largely by the discovery of young stars in the {\\em ROSAT} X-ray Satellite All-Sky Survey, over the last decade several nearby pre-main-sequence (PMS) star groups have been identified (for a review, see \\citealt{zuc04}).  Ranging in age from roughly 6\\,Myr to $\\sim$100\\,Myr, five separate groups can be distinguished: TW Hydrae Association (TWA), $\\eta$~Chamaeleontis cluster ($\\eta$\\,Cha), $\\beta$~Pictoris Moving Group (BPMG), Tucanae-Horologium association (TUCHOR) and AB~Doradus Moving Group (ABD).  The common space motions and localized sky positions suggest that these groups are likely connected to the Sco-Cen star forming region located $\\sim$100\\,pc away in the southern hemisphere.  \\citet{mam99} and \\citet{son03} have traced back the space motion of members in BPMG, TWA and $\\eta$\\,Cha and argue that the groups are related to a star formation burst in the Sco-Cen region as a result of the passing of the Carina arm $\\sim$60\\,Myr ago.  These groups, because of their close vicinity, are excellent laboratories for studying star and planet formation.  Well constrained ages are necessary to make conclusions about timescales of, e.g., disk dissipation and planet formation.  Already, observations from the same sample presented in this paper have revealed that accretion disks can last up to $\\sim$10\\,Myr, but beyond this is rare \\citep{jay06}.  \\defcitealias{bar98}{BCAH98} Previously, ages have been derived from Hertzsprung-Russell (HR) diagram fitting, group dynamics and lithium abundance measurements. \\citet{luh04} provide an HR diagram isochrone age for $\\eta$\\,Cha of $6^{+2}_{-1}$\\,Myr derived from the evolutionary models of \\citeauthor{bar98} (\\citeyear{bar98}; hereafter \\citetalias{bar98})  and \\citet{pal99}, which agrees well with the dynamical expansion age of 6.7\\,Myr determined by \\citet{jil05}. The dynamical age of TWA has been harder to determine because of inconsistent space motions among its more than 30 members. An inferred age of $8.3\\pm0.8$\\,Myr is given to TWA based on the dynamical motion of four members \\citep{del06}. However, the likely complex dynamical evolution of TWA has led to several plausible evolutionary scenarios. \\citet{mak05} attribute this complex evolution to a chance encounter with Vega, while \\citet{law05} suggest TWA is composed of two separate groups, based on bimodal rotation period distributions with distinctly separate ages of $\\sim$\\,10\\,Myr and $\\sim$\\,17\\,Myr. More recently, \\citet{bar06} finds a conservative age of $10^{+10}_{-7}$\\,Myr by comparing ages from HR diagram isochrone comparisons (from BCAH98) and lithium abundances.  The slightly older group BPMG has an estimated age of $12^{+8}_{-4}$\\,Myr based on HR diagram isochrone comparisons (from \\citetalias{bar98}) and lithium abundances \\citep{zuc01}, with three dimensional motions that are consistent with a dynamical expansion age of 11.5\\,Myr \\citep{ort02}. \\citet{fei06} independently derived an age of $13^{+4}_{-3}$\\,Myr for the recently confirmed wide binary system of 51~Eri and GJ\\,3305, part of BPMG.  Known to be older than BPMG, but younger than the Pleiades, TUCHOR has an age of 20--40\\,Myr based on H$\\alpha$ measurements, X-ray luminosity, rotation and lithium abundances in comparison to other young clusters like TWA and the Pleiades \\citep{zuc00,ste00}.  Perhaps, the most debated age is that of the ABD group. \\citet{zuc04b} derived an age of 50\\,Myr by comparing the H$\\alpha$ emission strength of ABD members to members of the younger TUCHOR association, in addition to fitting its three M-type members to HR diagram isochrones. In contrast, \\citet{luh05} compared HR isochrones of ABD members to those of two well-observed clusters with ages of 50 and 125\\,Myr and suggested that the ABD group is coeval with the Pleiades at an age of 100--125\\,Myr. The latter age is strongly supported by \\citet{ort07}, who compute full 3D galactic orbits of ABD and the Pleiades cluster and show the dynamics of the two groups can be traced back to a common origin of $119\\pm20$\\,Myr ago.  A relatively new approach to age estimates is to use the evolution of the lithium abundance for low-mass, partially and fully convective PMS stars \\citep{bil97,jef05}. The initiation and duration of lithium depletion in PMS stars is dependent on mass and is very sensitive to the central temperature. Lithium is converted into helium in $p,\\alpha$ reactions in cores of low-mass stars when the temperature reaches $2.5\\times10^{6}$\\,K. The lower the stellar mass, the longer the time it takes to reach this critical temperature. For example, a 0.6\\,\\Msun\\ star begins to burn lithium at an age of 3\\,Myr, while a lower mass star at 0.1\\,\\Msun\\ begins to burn lithium at an age of 40\\,Myr. Stars with $M<0.06$\\,\\Msun\\ never reach this typical temperature, while stars with 0.6\\,\\Msun\\ $<M<$ 1.2\\,\\Msun\\ burn lithium for a short period (1--2\\,Myr) until a radiative core develops, and more massive pre-main sequence stars do not destroy lithium in their envelope at all. The result of these processes is a dip in the lithium abundance as a function of luminosity (and consequently, a function of effective temperature), only affecting stars with spectral types late than F5. As a group ages, this dip becomes deeper and widens on the cool end as progressively cooler stars reach the critical core temperature.  The cool end of this dip has been used to date coeval groups containing late M-dwarf stars by identifying the lithium depletion boundary (LDB). The LDB marks the luminosity above which all stars will have depleted their lithium. The lithium is very quickly depleted in these low-mass stars, so the LDB marks a sharp jump from initial to near depleted lithium abundances. As the temperature in the cores of PMS stars increases in time, the LDB will shift to cooler temperatures as a cluster ages. LDB ages have been determined for the Pleiades (125 $\\pm$ 8 Myr), $\\alpha$\\,Per (90 $\\pm$ 10 Myr), IC\\,2391 (53 $\\pm$ 5 Myr), and NGC\\,2547 (35 $\\pm$ 4 Myr) \\citep{sta98,bar99,sta99,bar04,jef05}.  In this paper, we fit more than 900 multi-epoch, high-resolution spectra, introduced in \\S\\ref{s:obs}, of 121 low-mass PMS stars to synthetic spectra created from PHOENIX model atmospheres (see \\S\\ref{s:models}). We begin by identifying the model with the best-fitting surface gravity, $\\log g$, and effective temperature, $T_\\mathrm{eff}$ for each observed spectrum, as described in \\S\\ref{s:gravtemp}. In \\S\\ref{s:Li}, we use these best-fit $T_\\mathrm{eff}$ and $\\log g$ to find the lithium abundance by fitting model spectra to both the 6\\,104\\,\\AA\\ and 6\\,708\\,\\AA\\ lithium features in the observations, but now using lithium abundance as a free parameter.  For comparison, we also measure the equivalent width (EW) of the 6\\,708\\,\\AA\\ lithium doublet which allows us to chronologically order the groups based solely on empirical measurements. We proceed by comparing the measured lithium abundance distribution to that predicted by PMS evolutionary models of \\citetalias{bar98} and \\citet{sie00} and estimate model dependent ages for each of the PMS groups. ",
        "conclusions": "\\label{s:conclusions}  We have measured effective temperatures, surface gravities, lithium equivalent widths and lithium abundances for 121 low-mass PMS stars from five nearby, PMS groups ranging in age from 8--125\\,Myr by performing least-squares fits of high resolution spectra to synthetic spectra created from PHOENIX model atmospheres \\citep{hau98}. To investigate the reliability of our measurements we compare the derived $T_\\mathrm{eff}$ and $\\log g$ with isochrones from PMS evolutionary models \\citepalias{bar98} as well as temperatures derived from spectral types.  Isochrones from PMS models for $\\log N_{\\mathrm{Li}}$ as a function of $T_\\mathrm{eff}$ are visually compared to the observed distribution. We find agreement between ages derived from PMS isochrones of \\citetalias{bar98} and \\citet{sie00} to ages calculated from other methods such as dynamical expansion ages. We find that $\\eta$\\,Cha and TWA have ages of $12{\\pm6}$\\,Myr and $12{\\pm8}$\\,Myr, respectively. BPMG has an age of $21{\\pm9}$\\,Myr, and TUCHOR has an age of $27{\\pm11}$\\,Myr. We can only constrain a tight lower limit for ABD, with an age greater than $45$\\,Myr, since, according to the PMS models, there is no more lithium depletion after $\\sim\\!45$\\,Myr for stars with $4\\,000<T_{\\mathrm{eff}}<6\\,000\\,$K. However, the halting of lithium depletion at this age and temperature is inconsistent with observations of radial velocity standards which demonstrate more depletion than the ABD group and the model predictions. Finally, we find that some of the ultrafast rotators in our sample have significantly less lithium depletion than other stars in the same group at the same temperature.  The consistent determination of $T_\\mathrm{eff}$ and $\\log g$ between multiple epochs ($\\sigma_T \\simeq 10$\\,K, $\\sigma_{\\log g} \\simeq 0.05\\,$dex) means that, in principle, we should be able to constrain those parameters with this precision.  As revealed by Fig.~\\ref{fig:logg_vs_t}, however, there is an apparent systematic offset between the $\\log g$ inferred from model spectra and $\\log g$ expected from models of stellar evolution (\\S\\ref{s:LiRes}). To account for these offsets, we introduce rough conservative external errors by examining how the measured parameters depend on each other. We find that the systematic errors in $\\log g$ of 0.5\\,dex lead to systematic errors of 100\\,K in our ability to constrain the $T_{\\mathrm{eff}}$. These external errors, similarly lead to offsets in the measured lithium abundances of 0.15\\,dex.  The small internal errors that we have measured imply that currently our accuracy is limited by the models.  With further improvements in the atmospheric models, there is a potential of comparing $T_\\mathrm{eff}$ and $\\log g$ directly to evolutionary models, thereby finding an age constraint independent of other estimators, such as color-magnitude diagram or lithium depletion boundary fitting.  Our data set would be well-suited for use with such future improved models.  Another use of our data set would be to derive both overall metallicity and abundances for individual elements.  While we do not believe this would affect our derived temperatures, etc., to a significant degree, it may be interesting to see how uniform the abundances are within (and between) groups, and whether there is any dependence on binarity, etc."
    },
    "0808/0808.1254_arXiv.txt": {
        "abstract": "% The observational link between Supermassive Black Holes (SMBH) and galaxies at low redshift seems to be very tight, and statistically the global evolution of star formation activity and BH accretion activity also seem to trace each other closely. However, pinning down the co-evolution of galaxies and BH on an object-by-object basis remains elusive. I present results from new models for the joint evolution of galaxies, SMBH, and AGN, which may be able to help resolve some of the observational puzzles. A unique aspect of these models is our treatment of self-regulated BH growth based on hydrodynamic simulations of galaxy-galaxy mergers. Although these models do quite well at reproducing the observed evolution of galaxies, they do not reproduce the observed history of BH accretion, predicting too much early accretion and not enough at late times. I suggest two possible resolutions to this problem. ",
        "introduction": "There is strong evidence that the present-day properties of Supermassive Black Holes (SMBH) and their host galaxies are tightly linked --- for example we observe a tight relationship between BH mass and the mass of the host spheroid in nearby galaxies \\citep[e.g.][]{haering:04}. However, we do not yet know whether this BH-galaxy mass relationship evolved over cosmic time, or how galaxies and SMBH arrived onto it. Perhaps most importantly, we do not yet understand the physical origin of this relationship.  Mergers have often been proposed as a mechanism that can drive gas into the nuclei of galaxies and fuel both starbursts and black hole accretion. \\citet{hopkins:05} used a large suite of hydrodynamic simulations of galaxy mergers to characterize the episode of black hole growth and the AGN accretion ``lightcurves'' in these events. \\citet{hopkins:06} then used this model to test whether observed galaxy merger rates, the build-up of red/spheroidal galaxies, and AGN luminosity functions are consistent with the picture in which mergers both trigger AGN activity and cause morphological and color transformation. They concluded that there is excellent statistical agreement between all of these quantities.  However, establishing a direct link between mergers and AGN activity has proven controversial. Low luminosity AGN, which are more common and can be studied in more detail at relatively low redshift ($z<1$), do not show strong evidence for being predominantly in morphologically disturbed hosts or close pairs \\citep[e.g.][]{li:06,pierce:06}. Higher luminosity quasars seem to be more often (though not always) associated with disturbed hosts or to have close companions \\citep{bahcall:97,bennert:08,letawe:08}, but the samples with high resolution imaging are small. It may be the case that AGN are most easily identified in the very late stages of mergers, when the two nuclei have completely coalesced and most signs of morphological disturbance have become invisible. Especially at high redshift, at the typically available image sensitivity, these very late stage merger remnants might well appear to be ``normal'' spheroidal-type galaxies.  Of course it is also possible, even likely, that there are multiple modes of BH growth. There are ample suggestions in the literature that AGN activity could also be fed by bars or by stochastic accretion of cold gas in galactic nuclei. \\citet{hopkins_hernquist:06} presented a model for fueling of BH in isolated disks by stochastic accretion, and argued that while this fueling mode may be important in lower luminosity objects and at late times, most of the growth of today's massive BH occurred in the merger-driven mode at high redshift.  Another sign of co-evolution is that the global histories of star formation and BH accretion seem to trace each other remarkably closely over the whole range of lookback times over which these quantities have robust observational estimates. The global star formation rate density is almost exactly a factor of 2000 times the BH accretion rate density, from redshift $z\\sim5$ to the present. Moreover, even when divided by mass or accretion rate, ``matched'' populations of galaxies and black holes at $z<1$ also trace one another's activity (Zheng et al. in prep). However, we know that at least at $z<1$, the bulk of the star formation activity is associated with \\emph{isolated} disks, not mergers or spheroids \\citep{bell:05}. Therefore, it is also strange that these two kinds of activity trace each other, but do not seem to be occuring in the same types of objects (Zheng et al. in prep).  At higher redshift, there may be a larger mismatch between between the two kinds of activity. \\citet{faucher-giguere:08} obtained constraints on the global star formation density at $2<z<4.2$ from the Lyman-$\\alpha$ forest opacity in QSO spectra. They found that the hydrogen photoionization rate, and hence the star formation rate density, was remarkably flat over this redshift range, in contrast to the sharply peaked QSO luminosity density, which falls off sharply at $z>2$. This may indicate that there is a time delay between SF activity and QSO activity, which is naturally explained in the merger picture (any galaxy with cold gas can form stars, but galaxies have to ``wait'' for a major merger before significant accretion onto the BH can occur).  One way to test this picture, in which mergers are responsible for triggering both AGN activity and morphological and spectrophotometric transformation of galaxies, is to build cosmological models that treat the growth of galaxies, black holes, and AGN self-consistently. In addition, many astronomers now believe that the energy released by accreting black holes may play a crucial role in regulating galaxy formation.  Here we describe one such semi-analytic model for the joint formation of galaxies, black holes, and AGN, and present some predictions from these models. ",
        "conclusions": ""
    },
    "0808/0808.2055.txt": {
        "abstract": "In this paper, we present a systematical study of brane worlds of string theory on $S^{1}/Z_{2}$. In particular,  starting with the toroidal compactification of the Neveu-Schwarz/Neveu-Schwarz sector  in (D+d) dimensions, we first obtain an effective $D$-dimensional action, and then compactify one of the $(D-1)$ spatial dimensions by introducing two orbifold branes as its boundaries. We divide the whole set of the gravitational and matter field equations into two groups, one holds outside the two branes, and the other holds on them. By combining  the Gauss-Codacci and Lanczos equations, we write down explicitly the general gravitational field equations on each of the two branes, while using distribution theory we express the matter field equations on the branes in terms of the discontinuities of the first derivatives of the matter fields.  Afterwards, we address three important issues: (i) the hierarchy problem; (ii) the radion mass; and (iii) the localization of gravity, the 4-dimensional Newtonian effective potential and the Yukawa corrections due to the gravitational high-order Kaluza-Klein (KK) modes.  The mechanism of solving the hierarchy problem is essentially the combination of the large extra dimension  and  warped factor mechanisms together with the tension coupling scenario. With very conservative arguments, we find that the radion mass is of the order of $10^{-2}\\; GeV$. The gravity is localized on the visible brane,  and the spectrum of the gravitational KK modes is discrete and can be of the order of TeV. The corrections to the 4-dimensional Newtonian potential from the higher order of gravitational KK modes are exponentially  suppressed and can be safely neglected in current experiments. In an appendix, we  also present  a systematical and pedagogical study of the Gauss-Codacci equations and Israel's junction conditions across a (D-1)-dimensional hypersurface, which can be either spacelike or timelike. ",
        "introduction": "\\renewcommand{\\theequation}{1.\\arabic{equation}} \\setcounter{equation}{0}  Superstring and M-theory all suggest that we may live in a world that has more than three spatial dimensions.\u00a0 Because only three of these are presently observable, one has to explain why the others are hidden from detection.\u00a0 One such explanation is the so-called Kaluza-Klein (KK) compactification, according to which the size of the extra dimensions is very small (often taken to be on the order of the Planck length).\u00a0 As a consequence, modes that have momentum in the directions of the extra dimensions are excited at currently inaccessible energies.  Recently, the braneworld scenarios \\cite{ADD98,RS1} has dramatically changed this point of view and, in the process, received a great deal of attention. At present, there are a number of proposed models (See, for example, \\cite{branes} and references therein.). In particular, Arkani-Hamed {\\em et al} (ADD) \\cite{ADD98} pointed out that the extra dimensions need not necessarily be small and may even be on the scale of millimeters.\u00a0 This model assumes that Standard Model fields are confined to a three (spatial) dimensional hypersurface (a 3-brane) living in a larger dimensional bulk  while the gravitational field propagates in the  bulk.\u00a0 Additional fields may live only on the brane or in the  bulk, provided that their current undetectability is consistent with experimental bounds. One of the most attractive features of this model is that it may potentially resolve the long standing hierarchy problem, namely the large difference in magnitudes between the Planck and electroweak scales, %\\bq %\\lb{1.0} %\\frac ${M_{pl}}/{M_{EW}} \\simeq 10^{16}$, %\\eq where $M_{pl}$ denotes the four-dimensional Planck mass with $M_{pl} \\sim 10^{16} \\; TeV$, and $M_{EW}$  the electroweak scale with $M_{EW} \\sim  TeV$.  Considering a N-dimensional spacetime and assuming that the extra dimensions are homogeneous and finite, we find \\bqn \\label{1.1} S^{(N)}_{g} &\\sim& -{M^{N-2}}  \\int{dx^{4} dz^{n}\\sqrt{- g^{(N)}} R^{(N)}} \\nb\\\\ &=&  -{M^{N-2}V_{n}}\\int{d^{4}x  \\sqrt{- g^{(4)}} R^{(4)}}\\nb\\\\ &\\simeq& -{M^{2}_{pl} }\\int{d^{4}x  \\sqrt{- g^{(4)}} R^{(4)}}, \\eqn where  $V_{n}$ denotes the volume of extra dimensions, $n \\equiv N-4$, and $M$  the N-dimensional fundamental Planck mass, which is related to $M_{pl}$ by \\bq \\lb{1.2} M = \\left({M^{2}_{pl}}/{V_{n}}\\right)^{1/(2+n)}. \\eq Clearly, for any given   extra dimensions, if $V_{n}$  is large enough, $M$ can be as low as the electroweak scale, $M \\simeq M_{EW} \\simeq TeV$. Therefore, if we consider  $M$ as the fundamental scale and $M_{pl}$ the deduced one, we can see that the hierarchy between the two scales is exactly due to the dilution of the spacetime in high dimensions, whereby the hierarchy problem is resolved. Table top experiments show that Newtonian gravity is valid at least down to the size $R \\sim 44 \\; \\mu{m}$ \\cite{Hoyle04}. From the above we can see that for $n \\ge 2$ the N-dimensional Planck mass $M$ can be lowered down to the electroweak scale  from the four-dimensional Planck scale.   In a different model, Randall and Sundrum (RS1) \\cite{RS1} showed that if the self-gravity of the brane is included, gravitational effects can be localized near the Planck brane at low energy and the 4D Newtonian gravity is reproduced on this brane. In this model, the extra dimensions are not homogeneous, but warped. One of the most attractive features of the model is that it will soon be fully explored by LHC \\cite{DHR00}. In the RS1 scenario, the mechanism to solve the hierarchy problem is completely different \\cite{RS1}. Instead of using large dimensions, RS used the warped factor, for which the mass $m_{0}$ measured on the invisible (Planck) brane is related to the mass $m$ measured on the visible (TeV) brane by $m = e^{-ky_{c}}m_{0}$, where $e^{-ky_{c}}$ is the warped factor. Clearly, by properly choosing the distance $y_{c}$ between the two branes, one can lower $m$ to the order of $TeV$, even $m_{0}$ is of the order of $M_{pl}$. It should be noted that the five-dimensional Planck mass $M_{5}$ in the RS1 scenario is still in the same order of $M_{pl}$. In fact, the 5-dimensional action $S^{5)}_{g}$ can be written as \\bqn \\lb{1.3} S^{(5)}_{g} &\\sim& -{M^{3}}\\int{dx^{4} d\\phi\\sqrt{- g^{(5)}} R^{(5)}}\\nb\\\\ &=& -{M^{3}}\\int^{\\pi}_{-\\pi}{r_{c}e^{-2kr_{c}|\\phi|}d\\phi} \\int{d^{4}x \\sqrt{- g^{(4)}} R^{(4)}} \\nb\\\\ &\\simeq& -{M^{2}_{pl} }\\int{d^{4}x  \\sqrt{- g^{(4)}} R^{(4)}}, \\eqn where now we have \\bq \\lb{1.4} M^{2}_{pl} = {M^{3}}{k^{-1}}\\left(1 - e^{-2ky_{c}}\\right) \\simeq M^{2}_{5}, \\eq for $k \\simeq M_{5}$ and $ky_{c} \\simeq 35$.  Another long-standing problem  is the cosmological constant problem: Its theoretical expectation values from quantum field theory exceed its observational limits by $120$ orders of magnitude  \\cite{wen}. Even if such high energies are suppressed by supersymmetry, the electroweak corrections  are still $56$ orders higher. This problem was further sharpened by recent observations of supernova (SN) Ia, which  reveal the revolutionary  discovery that our universe has lately been in its accelerated expansion phase  \\cite{agr98}. Cross checks from the cosmic microwave background radiation  and large scale structure all confirm it \\cite{Obs}. In Einstein's theory of gravity, such an expansion can be achieved by a tiny positive cosmological constant, which is well consistent with all observations carried out so far \\cite{SCH07}. Because of this remarkable result, a large number of ambitious projects have been aimed to distinguish the cosmological constant from dynamical dark energy models \\cite{DETF}. Since the problem is intimately related to quantum gravity, its solution is expected to come from quantum gravity, too. At the present, string/M-Theory is our best bet for a consistent  quantum theory of gravity, so it is  reasonable to ask what string/M-Theory has to say about the cosmological constant.  In the string landscape \\cite{Susk}, it is expected there are many different vacua with different local cosmological constants \\cite{BP00}. Using the anthropic principle, one may select the low energy vacuum in which we can exist. However, many theorists still hope to explain the problem without invoking the existence of ourselves. In addition, to have a late time accelerating universe from string/M-Theory, Townsend and Wohlfarth \\cite{townsend} invoked a time-dependent compactification of pure gravity in higher dimensions with hyperbolic internal space to circumvent Gibbons' non-go theorem \\cite{gibbons}. Their exact solution   exhibits a short period of acceleration. The solution is the zero-flux limit of spacelike branes \\cite{ohta}. If non-zero flux or forms are turned on,  a transient acceleration exists for both compact internal hyperbolic and flat spaces \\cite{wohlfarth}. Other accelerating solutions  by compactifying more complicated time-dependent internal spaces can be found  in \\cite{string}.   Recently,  we studied brane cosmology in the framework of  both string theory \\cite{WS07,WSVW08} and the Horava-Witten (HW) heterotic M Theory  \\cite{GWW07,WGW08}  on $S^{1}/Z_{2}$. From a pure numerology point of view, we found that the   4D  effective cosmological constant  can be cast in the form, \\bq \\lb{1.5} \\rho_{\\Lambda} = \\frac{\\Lambda_{4}}{8\\pi G_{4}} = 3\\left(\\frac{R}{l_{pl}}\\right)^{\\alpha_{R}}\\left(\\frac{M}{M_{pl}}\\right)^{\\alpha_{M}} M_{pl}^{4}, \\eq where $R$ denotes the typical size of the extra dimensions, $M$ is the energy scale of string or M theory, and $(\\alpha_{R}, \\alpha_{M}) = (10, 16)$ for string theory and $(\\alpha_{R}, \\alpha_{M}) = (12, 18)$ for the HW heterotic M Theory. In both cases, it can be shown that for $R \\simeq 10^{-22} \\; m$ and $M_{10} \\simeq 1\\; TeV$, we obtain $\\rho_{\\Lambda} \\sim \\rho_{\\Lambda, ob} \\simeq 10^{-47}\\; GeV^{4}$.  When orbifold branes are concerned, a critical ingredient is the radion stability. Using the mechanism of Goldberger and Wise \\cite{GW99}, we showed that the radion is stable.  Such studies were also generalized to the HW heterotic M Theory \\cite{HW96}, and found that, among other things, the radion is stable and has a mass of order of $10^{-2} \\; GeV$ \\cite{WGW08}.   In this paper, we shall give a systematical study of brane worlds of string theory on $S^{1}/Z_{2}$. Similar studies in 5-dimensional spacetimes have been carried out in the framework of both string theory \\cite{WSVW08} and M Theory \\cite{WGW08}. However, to have this paper as much independent as possible, it is difficult to  avoid repeating some of our previous materials, although we would  try our best to keep it to its minimum. The rest of the paper is organized as follows: In Sec. II, we consider the toroidal compactification of the Neveu-Schwarz/Neveu-Schwarz (NS-NS) sector  in (D+d) dimensions, and obtain an effective $D$-dimensional action. Then, we compactify one of the $(D-1)$ spatial dimensions by introducing two orbifold branes as the boundaries along this compactified dimension. In Sec. III, we divide the whole set of the gravitational and matter field equations into two groups, one holds outside the two branes, and the other holds on each of them. Combining  the Gauss-Codacci and Lanczos equations, we write down explicitly the general gravitational field equations on the branes, while using distribution theory we are able to express the matter field equations on the branes in terms of the discontinuities of the first derivatives of the matter fields. In Sec. IV, we study the hierarchy problem, while in Sec. V, we consider the radion mass by using  the Goldberger-Wise mechanism \\cite{GW99}. In Sec. VI we study the localization of gravity, the 4-dimensional effective potential and high order Yukawa corrections. In Sec. VII,  we present our main conclusions with some discussing remarks. We also include an appendix, in which we present a systematical and pedagogical study of the Gauss-Codacci equations and Israel's junction conditions across a surface, where the metric coefficients are only continuous, i.e., $C^{0}$, in higher dimensional spacetimes. To keep such a treatment as general as possible, the surface can be either spacelike or timelike.   Before turning to the next section, we would like to note that  in 4-dimensional spacetimes there exists Weinberg's no-go theorem for the adjustment of the cosmological constant \\cite{wen}. However, in higher dimensional spacetimes, the 4-dimensional vacuum energy on the brane does not necessarily give rise to an effective 4-dimensional cosmological constant. Instead, it may only curve the bulk, while leaving the brane still flat \\cite{CEG01}, whereby Weinberg's no-go theorem is evaded. It was exactly in this vein, the cosmological constant problem was studied in the framework of brane worlds in 5-dimensional spacetimes \\cite{5CC} and 6-dimensional supergravity \\cite{6CC}. However, it was soon found that in the 5-dimensional case hidden fine-tunings are required \\cite{For00}. In the 6-dimensional case such fine-tunings may not be needed, but it is still not clear whether loop corrections can be as small as  expected \\cite{Burg07}.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\renewcommand{\\theequation}{7.\\arabic{equation}} \\setcounter{equation}{0}  In this paper, we have  systematically studied the brane worlds of string theory on $S^{1}/Z_{2}$. Starting with the toroidal compactification of the Neveu-Schwarz/Neveu-Schwarz sector  in (D+d) dimensions, in Sec. II.A we have first obtained an effective $D$-dimensional action given by Eq.(\\ref{2.16}) for  non-vanishing dilaton field and flux with an effective potential given by Eq.(\\ref{2.17}). Then, in Sec. II.B we have compactified one of the $(D-1)$ spatial dimensions by adding two orbifold branes as the boundaries of the spacetime along the compactified dimension.  Variations of the total action with the metric and matter fields yield, respectively, the gravitational and matter field equations. This has been done in Sec. III and given by Eqs.(\\ref{3.3})-(\\ref{3.ee}). Dividing the whole set of the  field equations into two groups, one holds outside the two branes, and the other holds on  them, in Sec. III.A we have first written down the field equations outside the two branes, Eqs.(\\ref{3.ef})-(\\ref{3.ej}), while in Sec. III.B, we have written down explicitly the general gravitational field equations on each of the two branes,  Eqs. (\\ref{3.15})-(\\ref{3.17}), by combining  the Gauss-Codacci and Lanczos equations. On the other hand, by using the distribution theory,  we  have also been able to write down the matter field equations on the branes in terms of the discontinuities of the first derivatives of the matter fields, Eqs. (\\ref{3.27a})-(\\ref{3.28}).   In the study of orbifold branes, one of the most attractive features is that it may resolve the long standing hierarchy problem. In Sec. IV, we have shown explicitly how it can be solved in the current setup. The mechanism is essentially the combination of the ADD large extra dimension \\cite{ADD98} and RS warped factor \\cite{RS1} mechanisms together with the tension coupling scenario \\cite{Cline99}. In order to solve the hierarchy problem in the current setup,  the tensions of the branes are required to be in the order of the cosmological constant. %, which might represent   fine-tuning. %It should be noted that  this result %is quite general, and applicable to a large class of brane-world scenarios %\\cite{branes}.  Another important issue in brane worlds is the radion stability and radion mass \\cite{branes}. Previously, we showed that the radion is stable \\cite{WS07}. In this paper, we have devoted Sec. V to study the radion mass. With some very conservative arguments, we have found that the radion mass is of the order of $10^{-2}\\; GeV$, which is by far beyond its current observational constraint, $m_{\\varphi} > 10^{-3}\\; eV$.  In Sec. VI we have also shown that the gravity is localized on the visible (TeV) brane, in contrast to the RS1 model in which the gravity is localized on the Planck (hidden) brane \\cite{RS1}. In addition, the spectrum of the gravitational KK modes is discrete, and given explicitly by Eq.(\\ref{7.23}), which can be of the order of TeV. The corrections to the 4D Newtonian potential from the higher order gravitational KK modes are exponentially  suppressed and can be safely neglected [cf. Eq.(\\ref{7.27})].   In Appendix, we have also presented a systematical and pedagogical study of the Gauss-Codacci equations and Israel's junction conditions across a surface, which can be either spacelike or timelike,  in higher dimensional spacetimes.   It should be noted that, when studied the radion stability, we have ignored the backreaction of the perturbations. Although it is expected that the main results obtained here will be continuously valid even after taking such backreaction into acocunt, as what exactly happened in the Randall-Sundrum model \\cite{radion}, it would be very interesting to show explicitly that this is indeed the case.  Other important issues that have not been addressed in this paper include  the constraints from the solar system tests \\cite{solar}, and linear perturbations  in the current setup."
    },
    "0808/0808.0702_arXiv.txt": {
        "abstract": "\\noindent A number of positive and null results on the time variation of fundamental constants have been reported. It is difficult to judge whether or not these claims are mutually consistent, since the observable quantities depend on several parameters, namely the coupling strengths and masses of particles. The evolution of these coupling-parameters over cosmological history is also a priori unknown. A direct comparison requires a relation between the couplings. We explore several distinct scenarios based on unification of gauge couplings, providing a representative (though not exhaustive) sample of such relations. For each scenario we obtain a characteristic time dependence and discuss whether a monotonic time evolution is allowed. For all scenarios, some contradictions between different observations appear. We show how a clear observational determination of non-zero variations would test the dominant mechanism of varying couplings within unified theories. ",
        "introduction": "Any observation of a time variation of ``fundamental constants'' would be a far-reaching discovery. There are various claims for a detection, and many more observations indicating a null variation: criteria for judging their mutual consistency would be useful. We investigate whether a simple scenario exists which can account for several observational claims simultaneously and is consistent with unification of Standard Model (SM) gauge couplings. % Several observations motivate such a study. First, the claimed deviations from the present value of the fine structure constant $\\alpha$, or the proton-to-electron mass ratio $\\mu$, observed in quasar absorption systems. Second, the discrepancy between the primordial \\lise\\ abundance expected from standard nucleosynthesis (BBN) and seen in old halo stars, which may be explained by a variation of ``constants'' within a unified framework \\cite{DSW07,CocNunes06}. Third, the theoretical insight that scalar fields cannot be exactly constant over the entire cosmological evolution, and that a possible ``late'' time evolution can play an important role in the dynamics of the expanding Universe. A time variation of couplings arising from the evolution of a ``cosmon'' field in so-called ``quintessence scenarios'' \\cite{Wetterich:1987fk,DvaliZ} would link these variations to observables in cosmology \\cite{Wetterich:2002ic,Nunes:2003ff}.  Due to the many unknowns of the underlying particle physics models and the partly contradictory present observational situation, a systematic treatment is not easy, and may even seem premature. Nevertheless we consider a first attempt to be useful, in order to discuss strategies that can be used to compare variations of different observables. The power of the proposed method will only become clear if and when future observations present a less ambiguous picture.  The basic approach in the present paper relates the fractional variations of different fundamental couplings $G_k$, such as the fine structure constant $\\alpha$, the proton-electron mass ratio $\\mu$ or the ratio of the nucleon mass to the Planck mass, by an assumption of {\\it proportionality}, with fixed ``unification coefficients'' $d_k$. The choice of the values of $d_k$ is in turn determined within different scenarios of varying parameters in unified theories (GUTs) where the gauge couplings of the Standard Model converge at a unification scale $M_X$. The assumption of time-independent coefficients $d_k$ covers a large class of possible models for varying couplings. This assumption is, however, not a necessity, and we will describe specific quintessence models where it may not be realized in a forthcoming paper \\cite{Part2}.  In Section~\\ref{sec:Data} of this paper we review observational determinations of the variation or constancy of couplings, considering five types of methods: early-universe cosmology, astrophysical spectroscopy, nuclear physics in the Earth and the Solar System, gravitational physics, and atomic clock comparisons. In Section~\\ref{sec:Unification} we introduce the unification of couplings (Grand Unified Theory, GUT) and determine the implications of unification and supersymmetry (SUSY) for the Standard Model parameters and for the observables we consider. We further define six unified scenarios by considering different possibilities for the variation of the Fermi scale and the superpartner masses. Within each scenario we reduce the various observational results to constraints on the time evolution of a single fundamental coupling, and discuss the mutual consistency of observations.  In Section~\\ref{sec:Epochs_and_evolutionFactors} six cosmological epochs are introduced, and the constraints on variation deduced in Section~\\ref{sec:Unification} are collected into a set of ``evolution factors'' for each unified scenario. These evolution factors are a measure of the overall size of coupling evolution between a given epoch and the present. We then determine for each scenario to what extent a monotonic variation over time can be consistent with the data. Section~\\ref{sec:Conclude} draws some general conclusions.  In a subsequent paper \\cite{Part2} we will investigate the scalar field dynamics that could give rise to a small but nonzero variation of couplings. The presence of a cosmologically varying degree of freedom gives rise to important additional effects. It affects gravity on large scales, altering the expansion of the Universe and potentially giving rise to the observed late-time acceleration. Also on local scales, a light field weakly coupled to matter produces long-range forces which are tightly constrained \\cite{Will} by Solar System precision tests of gravity and the null results of experiments testing the Weak Equivalence Principle. Combining all these considerations leads to tighter constraints on models but also offers more possibilities to test them. ",
        "conclusions": "\\label{sec:Conclude} Within Grand Unified Theories, different measurements of the variation of fundamental constants can be consistently reduced to a variation of a few ``unification parameters'', namely the unification scale $M_X/M_{\\rm P}$, gauge coupling $\\alpha_X$, the Fermi scale $\\vev{\\phi}/M_X$ and SUSY-breaking masses $\\tilde{m}/M_X$. We define various GUT-scenarios for varying couplings by the assumption of proportionality of fractional variations of the unification parameters.  Assuming that couplings really vary, this is a way of excluding such GUT scenarios by demanding consistency of the implied variations. The assumption of proportionality permits us to project all observations into constraints on a common evolution factor $l(z)$ for each scenario. We show that different GUT scenarios yield different time evolutions of $l(z)$ assuming that certain claimed measurements of varying constants are correct. We confirm that ``simple'' models which have only one fundamental parameter varying ($\\alpha_X$ or $\\vev{\\phi}/M_X$) result in inconsistent variations. However, combined variations of these two parameters, as described in scenarios 5 and 6, lead to results more consistent with the possible quintessence-induced time variations of fundamental couplings which we investigate in \\cite{Part2}. %   Specifically, one may ask whether the claimed observations of variations in $\\alpha$ \\cite{Murphy:2003mi} and $\\mu$ \\cite{Reinhold:2006zn} are mutually consistent, and whether they are consistent with an explanation of the apparent primordial \\lise-depletion by varying couplings. Within a hypothesis of constant Yukawa couplings, which results in identical fractional variations of all quark and lepton masses, we investigated arbitrary variations of $\\alpha_X$, $\\vev{\\phi}/M_X$ and $M_X/M_{\\rm P}$. For scenarios with supersymmetry we also assumed that the SUSY-breaking masses vary proportional to $\\vev{\\phi}$, but the effect of such a variation only appears at higher order and is probably not crucial.  We have not found a scenario with a monotonic time evolution $l(z)$ that makes all three signals or hints of variation mutually consistent. A monotonic evolution requires either to discount one of the ``signals'' by substantially increasing its uncertainty, or to alter our assumptions by including additional time variation of some Yukawa couplings.  Our investigation shows how the variations of different couplings in the Standard Model may be compared. If the observational situation becomes clearer and at least one nonzero time variation is established, such methods may be used for new tests of the idea of grand unification.  \\subsection*{Note added} Shortly before the completion of this paper a new determination of the variation of $\\mu$ appeared \\cite{King:2008ud} reporting a reanalysis of spectra from the same two H$_2$ absorption systems as \\cite{Reinhold:2006zn}, and adding one additional system at $z\\simeq 2.8$. The results of the new analysis are not consistent with the previous claim indicating a nonzero variation, either considering all three systems or the two previously considered. The stringent null bound of the new analysis, $\\Delta \\mu/\\mu = (2.6 \\pm 3.0) \\times 10^{-6}$, would disfavour all scenarios except those where the fractional variation of $\\mu$ was of the same order as or smaller than that of $\\alpha$. This would require us to approach the ``special'', apparently fine-tuned values of $\\tilde{\\gamma}$ discussed in Section~\\ref{sec:Specials}, for which $\\mu$ variation (and any deviation from the standard \\lise\\ abundance at BBN) are suppressed.  \\subsection*{Acknowledgements}  We acknowledge useful discussions with M.~Pospelov, R.~Trotta, P.~Molaro, P.~Avelino, J.~Berengut and V.~Flambaum, and invaluable correspondence and discussions with M.~Murphy. T.\\,D. is supported by the {\\em Impuls- and Vernetzungsfond der Helmholtz-Gesellschaft}.   \\appendix \\newcommand{\\appsection}[1]{\\let\\oldthesection\\thesection \\renewcommand{\\thesection}{Appendix \\oldthesection}"
    },
    "0808/0808.2641_arXiv.txt": {
        "abstract": "We study the potential of GLAST to unveil particle dark matter properties with gamma-ray observations of the Galactic center region. We present full GLAST simulations including all gamma-ray sources known to date in a region of 4 degrees around the Galactic center, in addition to the diffuse gamma-ray background and to the dark matter signal. We introduce DMFIT, a tool that allows one to fit gamma-ray emission from pair-annihilation of generic particle dark matter models and to extract information on the mass, normalization and annihilation branching ratios into Standard Model final states. We assess the impact and systematic effects of background modeling and theoretical priors on the reconstruction of dark matter particle properties. Our detailed simulations demonstrate that for some well motivated supersymmetric dark matter setups with one year of GLAST data it will be possible not only to significantly detect a dark matter signal over background, but also to estimate the dark matter mass and its dominant pair-annihilation mode. ",
        "introduction": "The Gamma-ray Large Area Space Telescope (GLAST) \\cite{glastref} was successfully launched on June 11, 2008. The main instrument onboard GLAST, the Large Area Telescope (LAT) represents an improvement by more than one order of magnitude over the sensitivity of its predecessor EGRET \\cite{egretref} in the energy range from 20 MeV to 10 GeV, and extends the high energy coverage to about 300 GeV. These features make the LAT a tremendous tool for the indirect search for particle dark matter \\cite{Baltz:2008wd}, expected in the best motivated particle models to be a weakly interacting massive particle (WIMP) with a mass in the 10-1000 GeV range \\cite{dmreviews,kkdm}. WIMPs occasionally pair-annihilate in dark matter halos, producing, as a result, Standard Model particles. Prompt production as well as decays, hadronization and radiative processes associated to the annihilation products give rise to stable species including gamma rays, in an energy range extending up to the WIMP mass. Secondary radiation from the energetic electrons and positrons produced in annihilation events can also be used as an indirect detection diagnostic, e.g. in X-ray and radio wavelengths \\cite{multiw}.  The signal from dark matter annihilation is proportional to the product of the particle pair-annihilation rate and to the particle number density squared along the line of sight. The latter quantity provides a guideline as to which regions of the gamma-ray sky are the most promising for the detection of a signal from dark matter annihilation. As discussed e.g. in \\cite{Baltz:2008wd}, targets include the center of our own Galaxy, satellite galaxies of the Milky Way, nearby massive galaxies and clusters, as well as diffuse emission from the entire Galactic halo and the summed signal from unresolved sources at all redshifts in the extra-galactic diffuse gamma-ray flux. Numerous studies have addressed in detail all of these possibilities (for reviews, see e.g. \\cite{dmreviews,kkdm}).  With the dawn of the GLAST era, theoretical speculations on indirect WIMP detection with gamma rays will soon face high-quality data and the challenge of the corresponding analysis. The main purpose of the present study is to introduce DMFIT, a numerical package that, interfaced with any spectral fitting routine, allows one to fit gamma-ray data with the expected emission from the annihilation of generic WIMP models. For the purpose of illustration, we use here DMFIT in conjunction with XSPEC \\cite{xspec}. We show fits to gamma-ray spectra simulated using the experimental response function of the LAT and the software analysis tools that will be used for the actual GLAST data analysis. We apply DMFIT to the case of an ``isolated'' gamma-ray source, such as might be associated to dark matter annihilation in Galactic dark matter clumps or local satellite galaxies, as well as to the complex case of the Galactic center region, where background modeling and an accurate understanding of conventional gamma-ray sources play a critical role.  The Galactic center region has long been considered a promising target to look for a signature from particle dark matter annihilation in gamma rays. First estimates of the detectability of a dark matter pair-annihilation signal from the center of the Galaxy date back 30 years \\cite{Gunn:1978gr,Stecker:1978du}. An incomplete list of recent references that discussed the detection of dark matter pair-annihilation in the Galactic center with gamma-ray observations includes \\cite{Cesarini:2003nr,Stoehr:2003hf,Aloisio:2004hy,Bloom:2004aq,Pieri:2005vp,Mambrini:2005vk,Profumo:2005xd,Jacholkowska:2005nz,Hall:2006na,Zaharijas:2006qb,Bergstrom:2006ny,Aharonian:2006wh,Ahn:2007ty,Morselli:2007ch,Hooper:2007gi,Dodelson:2007gd,Regis:2008ij,Serpico:2008ga,Baltz:2008wd}. In particular, ref.~\\cite{Cesarini:2003nr} discussed the dark matter interpretation of the gamma-ray excess reported by EGRET from the Galactic center \\cite{MayerHasselwander:1998hg}; ref.~\\cite{Profumo:2005xd} and \\cite{Aharonian:2006wh} discussed a similar possibility for the high energy gamma-ray flux detected by HESS \\cite{Aharonian:2004wa}; ref.~\\cite{Zaharijas:2006qb,Dodelson:2007gd} presented studies where both the EGRET and the HESS data were simultaneously taken into account as a background to dark matter searches.  In addition, recently, the GLAST collaboration gave in ref.~\\cite{Baltz:2008wd} a comprehensive and updated overview of the GLAST sensitivity to dark matter annihilation signals for several possible sources inside and outside the Galaxy, using the Collaboration's current state of the art Monte Carlo and event reconstruction software.  As far as the Galactic center region is concerned, in the present study we include all gamma-ray sources known to-date within an angle of 4 degrees from the Galactic center, in addition to the diffuse Galactic gamma-ray emission. Of special importance is the question of how to properly model the innermost sources, associated to the Sag A${^*}$ region: we model those with three different scenarios, described in detail in sec.~\\ref{sec:saga}. We introduce three reference particle dark matter setups, which are particularly illustrative for the scope of the present analysis, chosen to be theoretically well motivated, phenomenologically viable and producing the correct thermal relic neutralino abundance. We carry out complete one year GLAST simulations, making use of an up-to-date LAT instrumental response function and pointing mode setup, as well as of the software analysis tools that will be used to analyze the actual GLAST data.  The main scopes of the present study are to: \\begin{enumerate} \\item present the DMFIT tool, and show examples of its application; \\item provide an updated template for the gamma-ray sources potentially relevant to dark matter searches in the Galactic center region; \\item assess the capabilities of GLAST to provide information on particle dark matter properties such as the mass and the dominant annihilation mode, both in virtually ``background-free'' setups and in the complex gamma-ray environment of the Galactic center; \\item study the theoretical bias that background modeling and theoretical priors (such as the dominant annihilation mode) produce in the estimation of the fundamental particle properties of dark matter from gamma-ray data. \\end{enumerate}  The paper is organized as follows: sec.~\\ref{sec:dmfit} introduces the DMFIT tool and \\ref{sec:glastsim} provides details on the GLAST simulations; sec.~\\ref{sec:sources} discusses the gamma-ray sources included in the simulations, and \\ref{sec:dm} describes the particle dark matter models; sec.~\\ref{sec:dmfitdmonlly} shows examples of applications of DMFIT to gamma-ray spectra produced by dark matter annihilation only.  Finally, sec.~\\ref{sec:gc} addresses the Galactic center region: we discuss there the optimal energy and angular regions for GLAST observations and the performance of GLAST at inferring particle dark matter properties from the gamma-ray spectrum from the center of the Galaxy.  Our summary and conclusions are given in sec.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl} With the detailed analysis of a specific example, the Galactic center region, this study reaffirms and reinforces the point that GLAST has the potential to play a pioneering role in the race towards the discovery of the fundamental nature of dark matter. While eagerly awaiting real data to go beyond simulations, the present theoretical study highlighted that: \\begin{itemize} \\item the nature of the EGRET source in the Galactic center, and its possible association to dark matter annihilation, will be conclusively probed by GLAST with less than one year of data; \\item the determination of the dark matter particle properties such as the mass and the dominant and sub-dominant annihilation modes is possible to a varying degree of accuracy, depending on the brightness of the source and on the knowledge of the background spectrum; \\item for realistic particle dark matter models, GLAST has the potential to (i) firmly establish the presence of a dark matter source in the Galactic center, even when it is not the brightest source in the region, (ii) estimate the dark matter particle mass to better than 10\\%, and (iii) to disentangle the occurrence of more than one annihilation mode to 99\\% confidence level; \\item if the EGRET source at the Galactic center is associated to dark matter annihilation, we showed that the optimal energy range includes all photons above 0.1 GeV within an angular region of $0.5^\\circ$; otherwise, the optimal energy range for dark matter searches is above 1 GeV, and the optimal angular region goes from $1^\\circ$ for steep dark matter profiles, to $\\sim10^\\circ$ for shallower profiles; \\item the finite spatial extent of the dark matter source at the Galactic center amounts to a systematic effect in the estimate of the normalization of the dark matter signal by a factor $\\sim 1.6$ compared to the point source approximation, even for very steep profiles; \\item the estimate of the dark matter particle properties is affected by two sources of bias: (i) the background model and (ii) assumptions on the dark matter annihilation mode; specifically, we showed that different backgrounds can lead to both under- and over-estimates of the dark matter mass, and that the inclusion of annihilation modes with hard spectra biases estimates of the mass towards lower values. \\end{itemize} We introduced and showed applications of DMFIT, a numerical package that, interfaced with any spectral fitting routine, allows one to fit gamma-ray spectra to the emission of generic WIMP models. DMFIT  includes the $e^+e^-$  annihilation mode, relevant e.g. for Kaluza-Klein dark matter, and is able to fit for low neutralino masses. In addition, we provided an overview and a template of all known gamma-ray sources relevant for dark matter searches in the Galactic center region. While we do not claim here that the Galactic center is the most promising site to search for a signal from dark matter, this study indicates that GLAST data from that region can potentially be of tremendous relevance in the quest for dark matter."
    },
    "0808/0808.0472_arXiv.txt": {
        "abstract": "The first phase of stellar evolution in the history of the universe may be Dark Stars, powered by dark matter heating rather than by fusion. Weakly interacting massive particles, which are their own antiparticles, can annihilate and provide an important heat source for the first stars in the the universe.  This talk presents the story of these Dark Stars. We make predictions that the first stars are very massive ($\\sim 800 M_\\odot$), cool (6000 K), bright ($\\sim 10^6 L_\\odot$), long-lived ($\\sim 10^6$ years), and probable precursors to (otherwise unexplained) supermassive black holes. Later, once the initial DM fuel runs out and fusion sets in, DM annihilation can predominate again if the scattering cross section is strong enough, so that a Dark Star is born again. ",
        "introduction": "In October 2006, as guests of the Galileo Galilei Institute in Florence, two of us began a new line of research: the effect of Dark Matter particles on the very first stars to form in the universe. We found a new phase of stellar evolution: the first stars to form in the universe may be ``Dark Stars:'' dark matter powered rather than fusion powered. We first reported on this work in a paper (\\cite[Spolyar, Freese, \\& Gondolo 2008]{SpolyarFreeseGondolo08}) submitted to the arxiv in April 2007 (hereafter Paper I). When we presented this work at The First Stars conference in Santa Fe soon after (\\cite[Freese,Gondolo \\& Spolyar 2008]{FreeseGondoloSpolyar08}), many questions were raised, which we have addressed in the subsequent year. In this talk I review the basic ideas as well as report on the followup work we performed over the past year.  The Dark Matter particles considered here are Weakly Interacting Massive Particles (WIMPs) (such as the Lightest Supersymmetric Particle), which are one of the major motivations for building the Large Hadron Collider at CERN that will begin taking data very soon. These particles are their own antiparticles; they annihilate among themselves in the early universe, leaving the correct relic density today to explain the dark matter in the universe.  These particles will similarly annihilate wherever the DM density is high. The first stars are particularly good sites for annihilation because they form at high redshifts (density scales as $(1+z)^3$) and in the high density centers of DM haloes. The first stars form at redshifts $z \\sim 10-50$ in dark matter (DM) haloes of $10^6 M_\\odot$ (for reviews see e.g. \\cite[Ripamonti \\& Abel 2005]{RipamontiAbel05}, \\cite[Barkana \\& Loeb 2001]{BarkanaLoeb01}, \\cite[Bromm \\& Larson 2003]{BrommLarson03}; see also \\cite[Yoshida et al. 2006]{Yoshida_etal06}.) One star is thought to form inside one such DM halo.  As our canonical values, we will use the standard annihilation cross section, $\\langle \\sigma v \\rangle = 3 \\times 10^{-26} {\\rm cm^3/sec}$, and $m_\\chi = 100$ GeV for the particle mass; but we will also consider a broader range of masses and cross-sections. Paper I  found that DM annihilation provides a powerful heat source in the first stars, a source so intense that its heating overwhelms all cooling mechanisms; subsequent work has found that the heating dominates over fusion as well once it becomes important at later stages (see below). Paper I (\\cite[Spolyar, Freese, \\& Gondolo 2008]{SpolyarFreeseGondolo08}) suggested that the very first stellar objects might be ``Dark Stars,'' a new phase of stellar evolution in which the DM -- while only a negligible fraction of the star's mass -- provides the power source for the star through DM annihilation.  \\begin{figure}[t] \\centerline{\\includegraphics[width=0.5\\textwidth]{sellfig.pdf}} \\caption{Adiabatically contracted DM profiles in the first protostars for an initial NFW profile (dashed line) using (a) the Blumenthal method (dotted lines) and (b) an exact calculation using Young's method (solid lines), for $M_{\\rm vir}=5 \\times 10^7 M_\\odot$, $c=2$, and $z=19$.  The four sets of curves correspond to a baryonic core density of $10^4, 10^8, 10^{13},$ and $10^{16}{\\rm cm}^{-3}$.  The two different approaches to obtaining the DM densities find values that differ by less than a factor of two. } \\end{figure} ",
        "conclusions": "The line of research we began in Florence almost two years ago is reaching a very fruitful stage of development. Dark matter can play a crucial role in the first stars.  The first stars to form in the universe may be Dark Stars: powered by DM heating rather than by fusion.  Our work indicates that they may be very large ($850 M_\\odot$ for 100 GeV mass WIMPs). The connections between particle physics and astrophysics are ever growing!"
    },
    "0808/0808.2194_arXiv.txt": {
        "abstract": "We present the results of extensive multi-waveband monitoring of the blazar 3C~279 between 1996 and 2007 at X-ray energies (2-10 keV), optical R band, and 14.5 GHz, as well as imaging with the Very Long Baseline Array (VLBA) at 43 GHz. In all bands the power spectral density corresponds to ``red noise'' that can be fit by a single power law over the sampled time scales. Variations in flux at all three wavebands are significantly correlated. The time delay between high and low frequency bands changes substantially on time scales of years. A major multi-frequency flare in 2001 coincided with a swing of the jet toward a more southerly direction, and in general the X-ray flux is modulated by changes in the position angle of the jet near the core. The flux density in the core at 43 GHz---increases in which indicate the appearance of new superluminal knots---is significantly correlated with the X-ray flux.  We decompose the X-ray and optical light curves into individual flares, finding that X-ray leads optical variations (XO) in 6 flares, the reverse occurs in 3 flares (OX), and there is essentially zero lag in 4 flares. Upon comparing theoretical expectations with the data, we conclude that (1) XO flares can be explained by gradual acceleration of radiating electrons to the highest energies; (2) OX flares can result from either light-travel delays of the seed photons (synchrotron self-Compton scattering) or gradients in maximum electron energy behind shock fronts; and (3) events with similar X-ray and optical radiative energy output originate well upstream of the 43 GHz core, while those in which the optical radiative output dominates occur at or downstream of the core. ",
        "introduction": "Blazars form a subclass of active galactic nuclei (AGN) characterized by violent variability on time scales from hours to years across the electromagnetic spectrum. It is commonly thought that at radio, infrared, and optical frequencies the variable emission of blazars originates in relativistic jets and is synchrotron in nature \\citep{imp88,mar98}. The X-ray emission is consistent with inverse Compton (IC) scattering of these synchrotron photons, although seed photons from outside the jet cannot be excluded \\citep{mau96,rom97,cop99,bla00,sik01,chi02,arb05}. The models may be distinguished by measuring time lags between the seed-photon and Compton-scattered flux variations. Comparison of the amplitudes and times of peak flux of flares at different wavebands is another important diagnostic. For this reason, long-term multi-frequency monitoring programs are crucially important for establishing a detailed model of blazar activity and for constraining the physics of relativistic jets. Here we report on such a program that has followed the variations in emission of the blazar 3C~279 with closely-spaced observations over a time span of $\\sim$10 years.  The quasar 3C~279 \\citep[redshift=0.538;][]{bur65} is one of the most prominent blazars owing to its high optical polarization and variability of flux across the electromagnetic spectrum. Very long baseline interferometry (VLBI) reveals a one-sided radio jet featuring bright knots (components) that are ``ejected'' from a bright, presumably stationary ``core'' \\citep{jor05}. The measured apparent speeds of the knots observed in the past range from $4c$ to $16c$ \\citep{jor04},  superluminal motion that results from relativistic bulk velocities and a small angle between the jet axis and line of sight. Relativistic Doppler boosting of the radiation increases the apparent luminosity to $\\sim 10^4$ times the value in the rest frame of the emitting plasma.  Characterization of the light curves of 3C~279 includes the power spectral density (PSD) of the variability at all different wavebands.The PSD corresponds to the power in the variability of emission as a function of time scale. \\citet{law87} and \\citet{mch87} have found that the PSDs of many Seyfert galaxies, in which the continuum emission comes mainly from the central engine, are simple ``red noise'' power laws, with slopes between $-1$ and $-2$. More recent studies indicate that some Seyfert galaxies have X-ray PSDs that are fit better by a broken power law, with a steeper slope above the break frequency \\citep{utt02, mch04, mar03, ede99, pou01}. This property of Seyferts is similar to that of Galactic black hole X-ray binaries (BHXRB), whose X-ray PSDs contain one or more breaks \\citep{bel90,now99}. However, in blazars most of the X-rays are likely produced in the jets rather than near the central engine as in BHXRBs and Seyferts. It is unclear {\\it a priori} what the shape of the PSD of nonthermal emission from the jet should be, a question that we answer with the dataset presented here.  The raw PSD calculated from a light curve combines two aspects of the dataset: (1) the intrinsic variation of the object and (2) the effects of the temporal sampling pattern of the observations. In order to remove the latter, we apply a Monte-Carlo type algorithm based on the ``Power Spectrum Response Method'' (PSRESP) of \\citet{utt02} to determine the intrinsic PSD (and its associated uncertainties) of the light curve of 3C~279 at each of three wavebands. Similar complications affect the determination of correlations and time lags of variable emission at different wavebands. Uneven sampling, as invariably occurs, can cause the correlation coefficients to be artificially low. In addition, the time lags can vary across the years owing to physical changes in the source. In light of these issues, we use simulated light curves, based on the underlying PSD, to estimate the significance of the derived correlation coefficients.  In {\\S}2 we present the observations and data reduction procedures, while in {\\S}3 we describe the power spectral analysis and its results. In {\\S}4 we cross-correlate the light curves to determine the relationship of the emission at different wavebands. Finally, in {\\S}5 we discuss and interpret the results, focusing on the implications regarding the location of the nonthermal radiation at different frequencies, as well as the physical processes in the jet that govern the development of flux outbursts in blazars. ",
        "conclusions": "This paper presents well-sampled, decade-long light curves of 3C~279 between 1996 and 2007 at X-ray, optical, and radio wavebands, as well as monthly images obtained with the VLBA at 43 GHz. We have applied an algorithm based on a method by \\citet{utt02} to obtain the broadband PSD of nonthermal radiation from the jet of 3C~279. Cross-correlation of the light curves allows us to infer the relationship of the emission across different wavebands, and we have determined the significance of the correlations with simulated light curves based on the PSDs. Analysis of the VLBA data yields the times of superluminal ejections and reveals time variations in the position angle of the jet near the core. We have identified 13 associated pairs of X-ray and optical flares by decomposing the light curves into individual flares. Comparison of the observed radiative energy output of contemporaneous X-ray and optical flares with theoretical expectations has provided a quantitative evaluation of synchrotron and SSC models. We have discussed the results by focusing on the implications regarding the location of the nonthermal radiation at different frequencies, physical processes in the jet, and the development of disturbances that cause outbursts of flux density in blazars. Our main conclusions are as follows: \\\\ (1) The X-ray, optical, and radio PSDs of 3C~279 are of red noise nature, i.e., there is higher amplitude variability at longer time scales than at shorter time scales. The PSDs can be described as power laws with no significant break, although a break in the X-ray PSD at a variational frequency $\\lesssim 10^{-8}$ Hz cannot be excluded at a 95\\% confidence level.\\\\ (2) X-ray variations correlate with those at optical and radio wavebands, as expected if nearly all of the X-rays are produced in the jet. The X-ray flux correlates with the projected jet direction, as expected if Doppler beaming modulates the mean X-ray flux level.\\\\ (3) X-ray flares are associated with superluminal knots, with the times of the latter indicated by increases in the flux of the core region in the 43 GHz VLBA images. The correlation has a broad peak at a time lag of $130^{+70}_{-45}$ days, with X-ray variations leading.\\\\ (4) Analysis of the X-ray and optical light curves and their interconnection indicates that the X-ray flares are produced by SSC scattering and the optical flares by the synchrotron process. Cases of X-ray leading the optical peaks can be explained by an increase in the time required to accelerate electrons to the high energies needed for optical synchrotron emission. Time lags in the opposite sense can result from either light-travel delays of the SSC seed photons or gradients in maximum electron energy behind the shock fronts. \\\\ (5) The switch to optical-leading flares during the major multi-frequency outburst of 2001 coincided with a decrease in the apparent speeds of knots from 16-17$c$ to 4-7$c$ and a swing toward the south of the projected direction of the jet near the core. This behavior, as well as the high amplitude of the outburst, can be explained if the redirection of the jet (only a 1$\\degr$-2$\\degr$ change is needed) caused it to point closer to the line of sight than was the case before and after the 2001-02 outburst. \\\\ (6) Contemporaneous X-ray and optical flares with similar radiative energy output originate closer to the base of the jet, where the cross-section of the jet is smaller, than do flares in which the optical energy output dominates. This is supported by the longer time delay in the latter case. This effect is caused by the lower electron density and magnetic field and larger cross-section of the jet as the distance from the base increases.  Further progress in our understanding of the physical structures and processes in compact relativistic jets can be made by increasing the number of wavebands subject to intense monitoring. Expansion of such monitoring to a wide range of $\\gamma$-ray energies will soon be possible when GLAST scans the entire sky several times each day. When combined with similar data at lower frequencies as well as VLBI imaging, more stringent tests on models for the nonthermal emission from blazars will be possible."
    },
    "0808/0808.2477_arXiv.txt": {
        "abstract": "It has been argued that low-luminosity dwarf galaxies are the dominant source of ionizing radiation during cosmological reionization.  The fraction of ionizing radiation that escapes into the intergalactic medium from dwarf galaxies with masses less than $\\sim$$10^{9.5}$ solar masses plays a critical role during this epoch.  Using an extensive suite of very high resolution (0.1 pc), adaptive mesh refinement, radiation hydrodynamical simulations of idealized and cosmological dwarf galaxies, we characterize the behavior of the escape fraction in galaxies between $3 \\times 10^6$ and $3 \\times 10^9$ solar masses with different spin parameters, amounts of turbulence, and baryon mass fractions.  For a given halo mass, escape fractions can vary up to a factor of two, depending on the initial setup of the idealized halo.  In a cosmological setting, we find that the time-averaged photon escape fraction always exceeds 25\\% and reaches up to 80\\% in halos with masses above $10^8$ solar masses with a top-heavy IMF.  The instantaneous escape fraction can vary up to an order of magnitude in a few million years and tend to be positively correlated with star formation rate.  We find that the mean of the star formation efficiency times ionizing photon escape fraction, averaged over all atomic cooling ($T_{\\rm vir} \\ge 8000~$K) galaxies, ranges from 0.02 for a normal IMF to 0.03 for a top-heavy IMF, whereas smaller, molecular cooling galaxies in minihalos do not make a significant contribution to reionizing the universe due to a much lower star formation efficiency.  These results provide the physical basis for cosmological reionization by stellar sources, predominately atomic cooling dwarf galaxies. ",
        "introduction": "In most calculations of cosmological reionization, it has to be assumed that the product of star formation efficiency and hydrogen ionizing photon escape fraction is $\\ge 0.01$ \\citep[e.g.][]{Gnedin00, Cen03a, Cen03b, Wyithe03a, Wyithe03b, Venkatesan03, Somerville03, Chiu03, Haiman03, Ciardi03, Sokasian03, Sokasian04, Wyithe07, Srbinovsky08} in order to reionize the universe early enough to be consistent with the Wilkinson Microwave Anisotropy Probe (\\textit{WMAP}) observations \\citep{Spergel07, Komatsu08}, if stars produce the majority of ionizing photons.  \\citet{Gnedin00} suggested that, based on Local Group dwarf galaxies, the star formation efficiency at high redshift is $\\sim 4\\%$, consistent with theoretical works \\citep[e.g.][]{Krumholz05, Krumholz07}.  Thus, unless star formation efficiency is unusually high (i.e., $\\ge 10\\%$) at high redshift, this requirement seems to suggest a high ionizing photon escape fraction from high redshift galaxies, $f_{\\rm esc} \\ge 10\\%$, may be necessary in the context of stellar reionization within the standard cold dark matter model.  However, $f_{\\rm esc} \\ge 10\\%$ is by no means the norm, at least from observations at lower redshifts.  Approximately 6\\% of ionizing radiation escape from the Milky Way \\citep{Bland99, Putman03}.  For local starburst galaxies \\citet{Hurwitz97} gave $f_{\\rm esc}\\le 0.032,0.052, 0.11$ ($2\\sigma$) for Mrk 496, Mrk 1267, and IRAS 08339+6517 ($\\le 0.57$ in the case of Mrk 66).  \\citet{Deharveng01} gave an escape fraction of $f_{\\rm esc}<0.062$ for Mrk 54. \\citet{Heckman01} found $f_{\\rm esc} \\le 0.06$.  \\citet{Bergvall06} found $f_{\\rm esc}\\sim 0.04-0.1$ for a local extreme starburst dwarf, the Blue Compact Galaxy Haro 11.  \\citet{Chen07} placed a 95\\% upper limit of $f_{\\rm esc}\\le 0.075$ for star forming regions hosting gamma-ray bursts at $z\\ge 2$.  For Lyman break galaxies at $z\\sim 3$, \\citet{Shapley06} found $f_{\\rm esc}\\sim 0.03$.  \\citet{Siana07} examined starburst galaxies at $z \\sim 1.3$ and found that less than 20\\% have relative escape fractions% \\footnote{This quantity is the absolute escape fraction that is corrected for dust attenuation at 1500 \\AA~and is a common measure in observations of escaping Lyman continuum.} near unity and a global absolute escape fraction of $\\lsim 0.04$.  However in this study, some galaxies have upper limits of $f_{\\rm esc, rel}$ as low as 0.08.  \\citet{Inoue06} recently compiled available observations over a wide range of redshift and come to the conclusion that there is a trend of increasing $f_{\\rm esc}$ with redshift from $\\le 0.01$ at $z\\le 1$ to $\\sim 0.1$ at $z\\ge 4$.  Some theoretical works also seem to point to low values for $f_{\\rm esc}$.  Theoretical models of our own Galaxy with realistic star formation mode by \\citet{Dove00} give an estimate of $f_{\\rm esc}\\sim 0.06$ or lower.  \\citet{Ciardi02} studied the effect of gas inhomogeneities on $f_{\\rm esc}$ for Milky Way like galaxies using simulations with radiative transfer and concluded that $f_{\\rm esc}$ depends on the density structure as well as star formation rate (SFR). In addition, porosity in the interstellar medium (ISM) that is caused by SN mechanical feedback may provide additional channels in which radiation could escape, thus increasing the UV escape fraction \\citep{Clarke02}.  \\citet{Wood00} argued that $f_{\\rm esc}\\le 0.01$ for galaxies at $z\\sim 10$.  \\citet{Ricotti00} found that $f_{\\rm esc}\\ge 0.1$ only for halos of total mass less than $\\sim 10^7\\msun$ at $z\\ge 6$ with $f_{\\rm esc}$ dropping precipitously for larger halos.  \\citet{Fujita03}, using ZEUS-3D simulations, found that $f_{\\rm esc}\\le 0.1$ from disks of dwarf starburst galaxies of total mass $M\\ge 10^{8-10}\\msun$.  \\citet[][hereafter RS06 and RS07]{Razoumov06, Razoumov07} concluded that $f_{\\rm esc}$ = 0.01--0.1 in several young galaxies with $\\mvir = 10^{12-13} \\Ms$ at $z \\sim 3$ in their simulations with radiative transfer using adaptive ray tracing, agreeing with the observations presented in \\citet{Inoue06}. \\citet[][hereafter GKC08]{Gnedin08}, using detailed hydrodynamic simulations with radiative transfer, found $f_{\\rm esc}\\sim 0.01-0.03$ for galaxies of total mass $M\\ge 10^{11}\\msun$ at $z=3-5$.  In the relevant redshift range for cosmological reionization, $z=6-15$, most of the ionizing radiation due to stars comes from dwarf galaxies of $M\\le 10^{8-9}\\msun$ in the standard cold dark matter model \\citep{Barkana01}.  This purpose of this paper is to investigate how $f_{\\rm esc}$ depends on various physical parameters expected in realistic cosmological settings using detailed adaptive mesh refinement (AMR) simulations coupled with adaptive 3D ray-tracing radiative transfer, focusing on galaxies with total mass $M=10^{6.5}-10^{9.5}\\msun$ at $z\\ge 6$.  We study the dependence of $f_{\\rm esc}$ on four physical parameters: mass of the galaxy, spin parameter, baryonic mass fraction and turbulent energy.  Our work is an extension of GKC08 and \\citet{Fujita03}, by exploring the dependence of $f_{\\rm esc}$ on some important physical parameters aforementioned, by including lower mass galaxies and by employing very high resolution ($0.1$pc) simulations.  In addition, importantly, we allow for disparate lifetimes of stars of different masses, of their subsequent explosive energy (for those stars that become supernovae) input into the ISM.  We first describe our radiation hydrodynamics simulations of isolated and cosmological halos in \\S\\ref{sec:sims}.  There we also describe our algorithm for star formation and feedback that considers a multi-phase ISM and resolved molecular clouds.  In \\S\\ref{sec:sf} and \\S\\ref{sec:fesc}, we present the star formation rates and history and the resulting escape fraction of UV radiation, respectively.  Next in \\S\\ref{sec:disc}, we further compare our results with previous work and discuss effects from any neglected physical processes in our simulations and the implications of our results on reionization scenarios.  We summarize our work in the last section. ",
        "conclusions": "\\label{sec:disc}  The fraction of ionizing radiation that escapes high-redshift galaxies is an important value to quantify in order to characterize cosmological reionization.  In this section, we first discuss the differences and similarities, along with their respective causes, between our results and previous work.  Next we discuss any implications and agreements with semi-analytic reionization models. We last elaborate on possible influences from neglected physical processes.  \\subsection{Possible Physical and Computational Dependencies}  Our results suggest that $f_{\\rm esc}$ is higher than 0.1 in such galaxies, but we must first understand why our $f_{\\rm esc}$ values differ from the more massive galaxies presented in the simulations of \\citet{Razoumov06, Razoumov07} and \\citet{Gnedin08}.  Below we discuss five possible causes.  \\textit{Galaxy morphology and environment}--- The galaxies simulated here exhibit an irregular morphology and did not form a rotationally-supported disk, whereas the simulations of RS06, RS07, and GKC08 all studied disk galaxies.  As mentioned in \\citet{Fujita03}, a turbulent and clumpy ISM may allow for radiation to escape more easily into the IGM.  Compared to a disk configuration, there are more low-density escape routes, i.e. porosity, for the radiation because of the clumpy nature of the gas \\citep{Clarke02}. Furthermore, the stellar orbits are not confined within a disk and can reach the outskirts of the galaxy at apocenter, where gas densities are much lower than the galactic center and ionizing radiation can escape into the IGM at a much greater rate.  Perhaps this irregular morphology is only initially present in dwarf galaxies because radiative feedback shifts the angular momentum distribution to higher values \\citep{Wise08a}.  A fraction of this material will return to the galaxy and aid in disk formation in higher mass galaxies, such as the ones studied in RS06, RS07, and GKC08.  Furthermore in rare halos, the intersecting filaments tend to be dense and thin, where a halo with equal mass at lower redshift would be contained in a large filament \\citep{Ocvirk08}.  This redshift dependency on halo environment could affect \\fesc~at $r > r_{\\rm vir}$, but should not be apparent at \\rvir.  \\textit{Galactic mass}--- GKC08 find that $f_{\\rm esc}$ increases as much as 3 orders of magnitude from a halo mass of \\tento{10}~to \\tento{11}\\Ms.  The cold \\ion{H}{1} disk in their lower mass galaxies tend to be more vertically extended, which brings about this precipitous drop.  If this trend continues to lower masses, the values of $f_{\\rm esc}$ would be negligible in our galaxies, but we see no such continuation.  Our results suggest that the radiative feedback in dwarf galaxies with $V_c \\lsim 35 \\kms$ have a profound impact on the ISM by driving outflows and preventing any disk formation, increasing \\fesc.  Perhaps this circular velocity is a critical turning point below which $f_{\\rm esc}$ is large ($>0.1$) because of dynamical effects from radiative feedback.  In more massive galaxies, SN feedback provides the main impetus for driving outflows.  \\textit{Star formation rates}--- One can argue that our simulations only capture the initial collapse and an artificially strong starburst in high-redshift galaxies; however, after $\\sim$20 Myr of star formation, this collapse is reversed by stellar feedback and SFRs stabilize.  In addition our most massive galaxy ($\\mvir = 4 \\times 10^9 \\Ms$) has an average SFR of 0.2 \\Ms~yr$^{-1}$, which is comparable to the SFRs found in the $\\mvir \\sim 10^{10} \\Ms$ galaxies in GKC08.  The star formation efficiencies ($M_\\star/M_{\\rm{gas}}$) of our galaxies range from 5--10\\%, also agreeing with previous galaxy formation simulations.  Hence we do not think our treatment of star formation affects the SFRs and resulting UV escape fraction.  \\textit{Stellar IMF}--- We experimented with both normal ($N_\\gamma = 2,600$) and top-heavy ($N_\\gamma = 26,000$) IMFs and found that a normal IMF causes $f_{\\rm esc}$ values to be lower by 10--75\\%.  The differences in how the gas and ensuing star formation reacts to radiative feedback most likely causes this spread.  The normal IMF is similar to the ones used in all previously quoted theoretical works, and thus we do not expect our choice in IMF to cause our high UV escape fractions.  One minor shortcoming of our simulations is that we keep the SN feedback strength equal in our normal and top-heavy IMFs. Although this helps us localize the cause of any changes to a different specific luminosity, the strength of SN feedback should also affect the gas distribution, SFR, and \\fesc.  \\textit{Resolution}--- We can afford a spatial resolution of $\\sim$0.1 physical pc in all of our simulations.  However in larger galaxies, it becomes increasingly difficult to maintain this resolution, especially in cosmological simulations.  The highest resolution simulation in RS06 and RS07 use a gravitational softening length of 330 $h^{-1}$ pc, and their radiative transfer grid is refined so that 10 particles occupy each cell.  GKC08 use AMR simulations with a maximum resolution of 50 physical pc.  \\citet{Fujita03} have a fixed resolution of 0.28 pc in their $z = 8, M_d = 10^8 \\Ms$ simulation, similar to our resolution.  The $f_{\\rm esc}$ values of our simulations and \\citeauthor{Fujita03} are in agreement.  Regular and giant molecular clouds have radii of 2--20 and 10--60 pc \\citep[see][for a review]{MacLow04}, respectively, and are usually accounted with subgrid physics in cosmological simulations.  However our simulations capture the turbulent and clumpy nature of the ISM, while resolving such molecular clouds.  As discussed before, the turbulent ISM and its porosity may play a key role in allowing radiation to escape \\citep{Clarke02}.  In addition, high-resolution simulations with radiative transfer can model the fragmentation of ionization fronts \\citep[e.g.][]{Fujita03, Whalen08a}, which can further assist in boosting UV escape fractions.  \\subsection{Impact on reionization models} \\label{sec:models}  In reionization models, whether it be based on Press-Schetcher formalism or N-body simulations, the product $\\fescs \\equiv f_\\star \\times \\fesc$ ultimately dictates the absolute number of photons that escape from each halo into the IGM.  When calibrating these models against \\textit{WMAP} observations \\citep{Spergel07, Komatsu08} and \\lya~transmission from $z \\sim 6$ quasars \\citep[e.g.][]{Bolton07, Srbinovsky08}, the values of $\\ge 0.01$ are usually required. Recently, \\citeauthor{Srbinovsky08} found that \\fesc~must lie within the range 0.1--0.25 with a best-fit of $f_\\star = 0.11$ if all halos with masses $\\gsim 10^9 \\Ms$ contribute to reionization.  Furthermore, they conclude that even the smallest atomic line cooling halos ($\\mvir \\gsim 10^8 \\Ms$) have $\\fesc \\sim 0.05$.  Our results support this scenario where low-luminosity galaxies are the main contributors to reionization.  We plot \\fescs~from idealized and cosmological halos in Figure~\\ref{fig:fesc_cstar}.  We calculate $f_\\star$ using the initial gas mass in halo at the initial time; thus we are overestimating $f_\\star$ up to a factor of $\\sim$1.5 because the halos experience mass accretion in the 100 Myr that we have simulated.  Recall that none of the halos presented here undergo a major merger.  In our least massive cosmological halos with $\\mvir \\le 3 \\times 10^7 \\Ms$, \\fescs~increases rapidly with respect to halo mass from \\tento{-3}~to \\tento{-2} because atomic hydrogen line cooling becomes efficient in these halos.  \\fescs~is also an order of magnitude higher than the idealized cases.  As discussed in \\S\\ref{sec:sfr}, the additional gas accretion from filaments results in higher SFRs or equivalently $f_\\star$.  In atomic cooling halos ($\\tvir > 8000$ K), we find \\fescs~to always be $\\ge 0.01$, sufficiently high enough to meet the ``critical'' value to match reionization constraints, with an average of 0.033 and 0.026 for top-heavy and normal IMFs, respectively.  When integrated over a luminosity function with a faint-end slope of --1.7 \\citep[e.g.][]{Bouwens07}, the average of \\fescs~in atomic cooling halos is 0.029 and 0.021 for top-heavy and normal IMFs, respectively. If we include all halos, this average drops to $4.1 \\times 10^{-3}$ because the low-mass halos are more abundant and cannot cool and form stars as efficiently.  In Figure~\\ref{fig:fesc_cstar_trends}, we show the behavior of \\fescs~with halo mass in idealized halos grouped by halo parameters, similar to Figure~\\ref{fig:trends}.  There is no apparent trends with respect to turbulent energy; however, halos with higher spin parameters result in smaller values of \\fescs~in high mass halos that form a rotationally supported disk.  The halos with $f_b = 0.1$ have the highest values of \\fescs, and gas-poor halos are consistently lower in all halo masses.  \\begin{figure}[t] \\begin{center} \\epsscale{1.15} \\plotone{f11_color} \\caption{\\label{fig:fesc_cstar} The product of stellar mass fraction ($M_{\\rm star}/M_{\\rm gas}$) and escape fraction of ionizing photons ($f_{\\rm esc}$) from idealized halos and cosmological halos with a top-heavy IMF (filled diamonds) and normal IMF (open diamonds) as a function of halo mass. Idealized halos with a normal IMF are plotted as large circles. This product is a key quantity in semi-analytical reionization models in determining the amount of escaping radiation per collapsed gas fraction.} \\end{center} \\end{figure} \\begin{figure*}[t] \\begin{center} \\epsscale{1.0} \\plotone{f12_color} \\caption{\\label{fig:fesc_cstar_trends} Relative changes in the product of stellar mass fraction ($M_{\\rm star}/M_{\\rm gas}$) and escape fraction of ionizing photons ($f_{\\rm esc}$) from all idealized models grouped by initial halo parameters: turbulent energy (\\textit{left}), spin parameter (\\textit{middle}), and baryon mass fraction (\\textit{right}) when compared to the ``control halos'' with ($f_{\\rm turb}, \\lambda, f_b$) = (0.25, 0.04, 0.1).} \\end{center} \\end{figure*}  \\subsection{Effects from the Ultraviolet Background}  Our simulations accurately track the evolution of the \\ion{H}{2} regions and how they breakout into the IGM.  Here we discuss an important process that was neglected in our simulations, external feedback from the UV background (UVB).  Our models sample halos that are primarily dependent on \\hh~cooling ($\\tvir \\lsim 8000$ K) and atomic line cooling in more massive halos. Gas condensation in the lower mass \\hh~cooling halos can be delayed by a \\hh~dissociating UV (Lyman-Werner) background between 11.2 and 13.6 eV \\citep{Machacek01, Wise07b, OShea08}.  The Lyman-Werner background thus increases cooling times in the centers of such halos.  As a result, the minimum mass of a star-forming halo increases with the Lyman-Werner background intensity.  This will not affect the global halo properties but may suppress the SFR in these halos.  The Lyman-Werner background becomes less of an issue in atomic line cooling halos as \\lya~cooling provides ample amounts of free electrons for \\hh~cooling, and they become self-shielding to this radiation \\citep{Wise08a}.  We now consider any photo-heating from a hydrogen ionizing UVB, which can partially suppress gas accretion and thus star formation in dwarf galaxies \\citep[e.g.][]{Efstathiou92, Shapiro94, Thoul96}.  Here the Jeans ``filtering mass'' \\citep{Gnedin98} accurately describes the minimum halo mass that can cool and collapse given an IGM thermal history \\citep{Gnedin00, Wise08b}.  The Jeans filtering mass is calculated by analyzing the linear evolution of overdensities in the presence of Jeans smoothing.  Because the low-density IGM has a dynamical time on the order of a Hubble time, it slowly reacts to any photo-heating, and the filtering mass can be thought of some running time-average of the Jeans mass.  Thus at some redshift, the Jeans and filtering mass may differ substantially.  \\citet{Thoul96} originally showed photo-heating could totally suppress any star formation in low-redshift dwarf galaxies with $V_c \\lsim 35 \\kms$ and somewhat lower SFRs in galaxies up to 100 \\kms.  At $z \\gsim 3$, halos are less susceptible to negative feedback from photo-heating because of their higher average densities and thus cooling rates \\citep{Dijkstra04}. Closer inspection of collapsing halos in radiation hydrodynamics simulations also shows high-redshift, low-mass halos can collapse and is indeed regulated by the filtering mass.  Prior to reionization, \\citet{Gnedin00} showed that the filtering mass smoothly increases from $\\sim$\\tento{6}~\\Ms~before any star formation occurs to \\tento{9}\\Ms~at $z = 6$.  \\citet{Wise08b} calculated the filtering mass in regions that were ionized by Population III stars and found that it gradually increases the filtering mass to $\\sim$$3 \\times 10^7 \\Ms$ at $z \\sim 15$ around biased regions.  Recently, \\citet{Mesigner08} studied how photo-heating from an inhomogeneous UVB suppresses gas cooling and star formation.  They found that at $z = 10$ an UVB intensity of $J_{21} \\sim 0.1$ is necessary to completely suppress gas collapse in $\\mvir = 10^8 \\Ms$ halos, where $J_{21}$ is in units of \\tento{-21}~\\emis.  In their semi-numerical simulations at $z = 10$, they showed this intensity occurs in a volume fraction of $<1\\%$, suggesting that star formation is widespread in these low-mass halos even in the advanced stages of reionization.  They also illustrate how the collapsed gas fraction steadily decreases with increasing UVB.  Even at $z = 7$ in a \\tento{8}\\Ms~halo, approximately 40\\% of the gas is retained when compared to the no feedback case.  This gas mass loss will directly affect star formation and perhaps the subsequent UV escape fraction. Using the results of \\citeauthor{Mesigner08} and the dependence of \\fescs~on gas fraction (see Fig. \\ref{fig:fesc_cstar_trends}), our results from cosmological halos can be adjusted to account for external feedback from the UVB.  For example, at $z = 10$ with their fiducial model, their semi-numerical simulations show that $J_{21} \\sim 0.01$ is most common, which decreases the gas fraction of \\tento{8}, \\tento{8.5}, \\tento{9}\\Ms~halos to 0.2, 0.6, 0.85, respectively, of the gas fraction without any UVB.  From Fig. \\ref{fig:fesc_cstar_trends}, we can see that \\fescs~drops by $\\sim$50\\% when the gas fraction is $\\le$0.075 in idealized halos.  We would expect a similar decrease in halos that lose gas from UVB feedback, resulting in $\\fescs \\sim 0.02$ for halos in this mass range.   \\subsection{Prior Star Formation Episodes}  We assumed that the halos were unaffected by any prior stellar feedback, but in this paper, we have also shown that low-mass halos are susceptible to gas ``blow-out'' caused by radiative feedback. This effect should be evident starting with Population III star formation \\citep[e.g.][]{Yoshida07, Wise08a} in halos with circular velocities up to $V_c \\sim 35 \\kms$.  For example, \\citeauthor{Wise08a} found that gas fractions of dwarf galaxies with $\\mvir \\sim 3 \\times 10^7 \\Ms$ at redshift 15 have been decreased to $\\le0.1$.  The cosmological halos presented here, which have $f_b \\sim 0.13$, have $f_{\\rm esc} \\sim 0.5$ with a top-heavy IMF; however if $f_b$ were lower from previous gas ``blow-out'', the trends shown in Figure \\ref{fig:trends} suggest that the escape fraction in these halos should be even higher, but their values of \\fescs~(Fig. \\ref{fig:fesc_cstar_trends}) will decrease because of lower SFRs.  To further understand and quantify the consequences of these neglected processes, semi-analytical reionization models, semi-numerical or N-body + radiative transfer simulations are needed, as discussed in \\S\\ref{sec:models}.  In future work, we plan to study this in detail by following the hierarchical assembly of high-redshift dwarf galaxies in cosmological simulations, including both metal-free and metal-enriched star formation and feedback.    We presented results from an extensive suite of very high resolution ($0.1~$pc) AMR radiation hydrodynamics simulations that focus on the UV escape fraction from isolated halos and cosmological halos.  The latter cases were extracted from two large-scale cosmological simulations.  These simulations accurately track the evolution of \\ion{H}{2} regions while including radiative and SNe feedback on the surrounding medium.  The halo masses studied in this work span from $3 \\times 10^6$ to $3 \\times 10^9 \\Ms$.  We used the isolated halos to gauge if the escape fraction depends on the turbulent fractional energy, spin parameter, and baryon mass fraction of the halo.  The cosmological halos provide a good estimate of realistic escape fractions of high-redshift dwarf galaxies.  We also investigated how a top-heavy IMF and a normal IMF affects the escape fraction.  Our main findings in this paper are  \\medskip  1. Radiative feedback from massive stars, primarily arising from D-type fronts, is most effective at ejecting gas from halos with masses less than $10^{8.5} \\Ms$ ($V_c = 35 \\kms$).  2. Radiation preferentially escapes through channels with low column densities, creating anisotropic \\ion{H}{2} and a highly varying $f_{\\rm esc}$ along different lines of sight, agreeing with previous work.  3. For a given halo mass, escape fractions in isolated halos have a standard deviation of 0.2, arising from differing physical halo parameters.  The largest effect comes from the baryon mass fraction, where $f_{\\rm esc}$ is lower in gas-rich halos with masses below \\tento{8} \\Ms.  In gas-rich halos with masses larger than \\tento{8} \\Ms, star formation is more intense, overcoming any large neutral fraction it must ionize and resulting in higher values of $f_{\\rm esc}$ than halos with smaller gas fractions.  4. High-redshift dwarf galaxies with $\\mvir > 10^7 \\Ms$, a top-heavy IMF, and irregular morphology have $0.25 \\le f_{\\rm esc} \\le 0.8$, which we determined from our simulations of cosmological halos.  A normal IMF decreases $f_{\\rm esc}$ to 0.05--0.1 in halos with $\\mvir < 10^{7.5} \\Ms$ and $f_{\\rm esc} \\sim 0.4$ in more massive halos, which have maximum SFRs spanning almost 2 orders of magnitude.  5. Escape fractions are dependent not only on the current SFR but on the photo-heating and dispersion of gas, following feedback from previous episodes of star formation.  Values of $f_{\\rm esc}$ can vary up to an order of magnitude in a few million years, i.e. the dynamical time of a molecular cloud on which variations in SFRs can occur.  6. The mean of the product of star formation efficiency and ionizing photon escape fraction, averaged over all atomic cooling ($\\tvir \\ge 8000~$K) galaxies, ranges from $0.02$ for a normal IMF to $0.03$ for a top-heavy IMF.  Smaller, molecular cooling galaxies in minihalos are not significant contributors to reionizing the universe primarily because of a much lower star formation efficiency in minihalos than in atomic cooling halos.  \\medskip  The high escape fraction of UV radiation has important implications on reionization, allowing a large amount of radiation from low-luminosity dwarf galaxies to freely propagate into the IGM.  Although our simulations miss the entire assembly of halos and their complete star formation history, they are robust in following the dynamics of star formation and feedback during the 100 Myr studied and suggest that escape fraction in these galaxies are larger than previously assumed."
    },
    "0808/0808.4121_arXiv.txt": {
        "abstract": "The recent {\\it Spitzer} detections of the 9.7$\\mum$ Si--O silicate emission in type 1 AGNs provide support for the AGN unification scheme. The properties of the silicate dust are of key importance to understanding the physical, chemical and evolutionary properties of the obscuring dusty torus around AGNs. Compared to that of the Galactic interstellar medium (ISM), the 10$\\mum$ silicate emission profile of type 1 AGNs is broadened and has a clear shift of peak position to longer wavelengths. In literature this is generally interpreted as an indication of the deviations of the silicate composition, size, and degree of crystallization of AGNs from that of the Galactic ISM. In this {\\it Letter} we show that the observed peak shift and profile broadening of the 9.7$\\mum$ silicate emission feature can be explained in terms of porous composite dust consisting of ordinary interstellar amorphous silicate, amorphous carbon and vacuum. Porous dust is naturally expected in the dense circumnuclear region around AGNs, as a consequence of grain coagulation. ",
        "introduction": "Dust is the cornerstone of the unification theory of active galactic nuclei (AGNs). This theory proposes that all AGNs are essentially ``born equal'' -- all types of AGNs are surrounded by an optically thick dust torus and are basically the same object but viewed from different lines of sight (see e.g. Antonucci 1993; Urry \\& Padovani 1995): type 1 AGNs, which always display broad hydrogen emission lines in the optical and have no obvious obscuring effect, are viewed face-on which allows a direct view of the central nuclei, while type 2 AGNs are viewed edge-on with most of the central engine and broad line regions being hidden by the obscuring dust.  Silicate dust, a major solid species in the Galactic interstellar medium (ISM) as revealed by the strong 9.7$\\mum$ and 18$\\mum$ bands (respectively ascribed to the Si--O stretching and O--Si--O bending modes in some form of silicate material, e.g. olivine Mg$_{2x}$Fe$_{2-2x}$SiO$_4$), has also been detected in AGNs both in {\\it emission} and in {\\it absorption} (see Li 2007 for a review).  The first detection of the silicate {\\it absorption} feature in AGNs was made at 9.7$\\mum$ for the prototypical Seyfert 2 galaxy NGC\\,1068 (Rieke \\& Low 1975; Kleinmann et al.\\ 1976), indicating the presence of a large column of silicate dust in the line-of-sight to the nucleus. It is known now that most of the type 2 AGNs display silicate {\\it absorption} bands (e.g. see Roche et al.\\ 1991, Siebenmorgen et al.\\ 2004, Hao et al.\\ 2007, Spoon et al.\\ 2007, Roche et al.\\ 2007) which is expected from the AGN unified theory -- for a centrally heated optically thick torus viewed edge-on, the silicate features should be in absorption.  For type 1 AGNs viewed face-on, one would expect to see the silicate features in {\\it emission} since the silicate dust in the surface of the inner torus wall will be heated to temperatures of several hundred kelvin to $\\simali$1000$\\K$ by the radiation from the central engine, allowing for a direct detection of the 9.7$\\mum$ and 18$\\mum$ silicate bands emitted from this hot dust. However, their detection has only very recently been reported in a number of type 1 AGNs covering a broad luminosity range, thanks to {\\it Spitzer} (Hao et al.\\ 2005, Siebenmorgen et al.\\ 2005, Sturm et al.\\ 2005, Weedman et al.\\ 2005, Shi et al.\\ 2006, Schweitzer et al.\\ 2008). It is worth noting that the silicate emission features have recently also been detected in type 2 QSOs (Sturm et al.\\ 2006, Teplitz et al.\\ 2006).  \\begin{figure} \\begin{center} \\psfig{figure=f1.cps, width=2.5in, angle=270} \\end{center} \\caption{\\label{fig:sil_obs_spec} Comparison of the silicate emission features of the quasar 3C\\,273 (Hao et al.\\ 2005) and the low-luminosity AGN NGC\\,3998 (Sturm et al.\\ 2005) with the silicate absorption feature of the ISM toward the Galactic center object Sgr A$^{\\ast}$ (Kemper et al.\\ 2004). Most notably is the long wavelength shift of the peak positions of 3C\\,273 and NGC\\,3998 in comparison with the ISM profile. All profiles are normalized to their peak values. } \\end{figure}  Compared to that of the Galactic ISM, the 9.7$\\mum$ silicate emission profiles of some AGNs appear ``anomalous''. As illustrated in Figure \\ref{fig:sil_obs_spec}, both quasars (high luminosity counterparts of Seyfert 1 galaxies; Hao et al.\\ 2005, Siebenmorgen et al.\\ 2005) and the low-luminosity AGN NGC\\,3998 (Sturm et al.\\ 2005) exhibit silicate emission peaks at a much longer wavelength ($\\simali$10--11.5$\\mum$), inconsistent with the ``standard'' silicate ISM dust (which peaks at $\\simali$9.7$\\mum$). The 9.7$\\mum$ feature of NGC\\,3998 is also much broader than that of the Galactic ISM (Sturm et al.\\ 2005).\\footnote{% We should note that the width of the interstellar 9.7$\\mum$ Si--O absorption feature is not universal, but varies from one sightline to another (see Draine 2003). Generally speaking, it is relatively narrow in diffuse clouds and broad in molecular clouds (Bowey et al.\\ 1998). In contrast, its peak wavelength is relatively stable. }  The deviations of the silicate emission profiles of type 1 AGNs from that of the Galactic ISM dust are generally interpreted as an indication of the differences between AGNs and the Galactic ISM in the composition, size distribution, and degree of crystallization of the silicate dust (Sturm et al.\\ 2005). However, we show in this {\\it Letter} that the observed peak shift and profile broadening of the 9.7$\\mum$ silicate emission features of AGNs can be explained in terms of porous composite dust consisting of ordinary interstellar amorphous silicate, amorphous carbon\\footnote{% There must exist a population of carbonaceous dust in the AGN torus, as revealed by the detection of the 3.4$\\mum$ absorption feature, attributed to the C--H stretching mode in saturated aliphatic hydrocarbon dust (see Li 2007 for a review). Whether the bulk form of the carbonaceous component in AGNs is amorphous carbon or hydrogenated amorphous carbon is not clear. } and vacuum, without invoking an exotic dust composition or size distribution. In the dense circumnuclear region around AGNs, a porous structure is naturally expected for the dust formed through the coagulation of small silicate and carbonaceous grains. ",
        "conclusions": "} It is well-known theoretically that as the size of a silicate grain increases, the 9.7$\\mum$ Si--O feature becomes wider with its peak shifted to longer wavelengths (see Fig.\\,9 of Dorschner et al.\\ 1995, Fig.\\,1 of Voshchinnikov \\& Henning 2008). But for compact silicate spheres to have its 9.7$\\mum$ Si--O feature peaking at $\\lambda_{\\rm peak} >10.5\\mum$, their size needs to exceed $\\simali$2$\\mum$. However, for these large grains the 9.7$\\mum$ Si--O feature fades away (see Fig.\\,6 of Greenberg 1996, Fig.\\,1 of Voshchinnikov \\& Henning 2008). Therefore, it is difficult to account for the ``anomalous'' silicate emission features of AGNs just in terms of an increase in grain size.  A non-spherical grain shape or shape distribution can also broaden the 9.7$\\mum$ Si--O feature and red-shift its peak wavelength (see Li 2008). But to account for the peak wavelengths of $\\lambda_{\\rm peak}$\\,$\\simali$10--11.5$\\mum$ observed in some AGNs, the dust has to be extremely elongated which is unrealistic (e.g. spheroidal dust needs to have an elongation $>$6).  The peak wavelength of the 9.7$\\mum$ Si--O feature also depends on the silicate mineralogy. But we are not aware of any amorphous silicate species that could give rise to a Si--O feature peaking at $\\lambda >10\\mum$.  Crystalline olivine has its Si--O feature peaks at $\\simali$11.2$\\mum$ (see Yamamoto et al.\\ 2008). But this feature is too sharp compared to the broad Si--O emission features of AGNs. Within the signal-to-noise ratio limits, we see no clear evidence for crystalline silicates in 3C\\,273 and NGC\\,3998 (see Fig.\\,\\ref{fig:sil_obs_spec}). So far, the detection of crystalline silicate dust in AGNs has only been reported in the BAL (broad absorption line) quasar PG\\,2112+059 (Markwick-Kemper et al.\\ 2007).\\footnote{% Spoon et al.\\ (2006) reported firstly the detection of narrow absorption features of crystalline silicates in ultraluminous IR galaxies (ULIRGs).  } The 9.7$\\mum$ Si--O feature emission features of the protoplanetary disks around T Tauri and Herbig Ae/Be stars are often also much broader than that of the ISM (e.g. see Bouwman et al.\\ 2001, Forrest et al.\\ 2004). This is usually interpreted as an indicator of grain growth (e.g. see Natta et al.\\ 2007). The porous composite dust model was recently used by Voshchinnikov \\& Henning (2008) to demonstrate that a similar behavior of the feature shape occurs when the porosity of fluffy dust varies. Kr\\\"ugel \\& Siebenmorgen (1994) have also demonstrated that a porous structure results in the broadening, weakening and redshifting of the 9.7$\\mum$ Si--O feature. Fluffy dust aggregates are also naturally expected in protoplanetary disks as a result of grain coagulation.  We plot in Figure \\ref{fig:agn_lambdap_porosity} as a function of porosity $P$ the peak wavelengths $\\lambda_{\\rm peak}$ of the 9.7$\\mum$ Si--O feature of the absorption efficiency profiles $\\Qabs(\\lambda)$ and the emission profiles obtained by folding $\\Qabs(\\lambda)$ with the Planck function $B_\\lambda(T)$ at various temperatures. It is seen that for 3C\\,273 and NGC\\,3998, to account for the observed $\\lambda_{\\rm peak}\\approx 10.6\\mum$, one requires either cool dust (with $T$\\,$\\simali$140--150$\\K$) with a low porosity ($P<0.3$), or warm dust (with $T$\\,$\\simali$170--220$\\K$) with a high porosity ($P>0.6$). The 18$\\mum$ O--Si--O emission feature further constrains that cool dust is preferred in 3C\\,273 while NGC\\,3998 appears to favour warm dust.  Also seen in Figure \\ref{fig:agn_lambdap_porosity} is that the range of the peak wavelengths $10\\mum \\simlt \\lambda_{\\rm peak}\\simlt 11.5\\mum$ of the 9.7$\\mum$ Si--O emission features of the five PG quasars of Hao et al.\\ (2005) falls within the predictions of the porous composite dust model. This indicates that the long-wavelength shifted 9.7$\\mum$ Si--O features of these quasars can be accounted for by the combined effects of porosity and temperature (i.e. either highly porous warm dust or cold dust with a lower porosity). The 18$\\mum$ O--Si--O emission feature will allow us to break this degeneracy. Finally, for those AGNs whose 9.7$\\mum$ Si--O absorption or emission features do not exhibit any long-wavelength shift, one may resort to dust with a low porosity $P<0.3$ (and $T>220\\K$ if they are in emission).  We should stress that although the combined effects of porosity and temperature suffice by themselves to explain the general trend of broadening the 9.7$\\mum$ Si--O emission features of AGNs and shifting their peak wavelengths to longer wavelengths, we by no means exclude the effects of other factors such as grain size,\\footnote{% The grain size effect is probably (at least in part) responsible for the nondetection of the 9.7$\\mum$ and 18$\\mum$ silicate emission features in many type 1 AGNs (e.g. see Hao et al.\\ 2007). The silicate size distribution in some AGNs might be dominated by large grains ($>$ a few $\\mum$; see Maiolino et al.\\ 2001a,b) or small silicate grains are depleted (e.g. see Laor \\& Draine 1993, Granato \\& Danese 1994). Alternative explanations include sophisticated torus geometries (e.g. tapered disk configurations [Efstathiou \\& Rowan-Robinson 1995], clumpy torus models [Nenkova et al.\\ 2002, 2008; but see Dullemond \\& van Bemmel 2005]) and an assumption of strong anisotropy of the source radiation (Manske et al.\\ 1998). } shape, and mineralogy. It is very likely that all these factors act together to produce the anomalous silicate emission features seen in AGNs.  Finally, we emphasize that although a steeply rising (cold) Planck function could redshift the 9.7$\\mum$ Si--O emission feature, the observed shift of the peak wavelength of this emission feature in AGNs cannot be purely a temperature effect since the silicate absorption profiles of some AGNs also appear anomalous; e.g., the ``9.7$\\mum$'' silicate feature of Mkn 231, a peculiar type 1 Seyfert galaxy, is seen in absorption peaking at $\\simali$10.5$\\mum$ (Roche et al.\\ 1983); Jaffe et al.\\ (2004) found that the 9.7$\\mum$ silicate absorption spectrum of NGC\\,1068 shows a relatively flat profile from 8 to 9$\\mum$ and then a sharp drop between 9 and 10$\\mum$; in comparison, the Galactic silicate absorption profiles begin to drop already at $\\simali$8$\\mum$.   To conclude, we have explored in a quantitative way on the effects of grain porosity on the silicate Si--O stretching feature using the multi-layered sphere model. It is found that the Si--O feature broadens and shifts to longer wavelengths with the increasing of dust porosity. We conclude that the combined effects of dust porosity and cool temperature of $T<$\\,200$\\K$ (which further redshifts the silicate feature) could explain the observed broadening and longer-wavelength shifting of the 9.7$\\mum$ Si--O feature of AGNs."
    },
    "0808/0808.1575_arXiv.txt": {
        "abstract": "We report the discovery of a wide ($135\\pm25$~AU), unusually blue L5 companion 2MASS J17114559+4028578 to the nearby M4.5 dwarf G 203-50 as a result of a targeted search for common proper motion pairs in the Sloan Digital Sky Survey and the Two Micron All Sky Survey.  Adaptive Optics imaging with Subaru indicates that neither component is a nearly equal mass binary with separation $> 0.18\\arcsec$, and places limits on the existence of additional faint companions.  An examination of TiO and CaH features in the primary's spectrum is consistent with solar metallicity and provides no evidence that G 203-50 is metal poor.  We estimate an age for the primary of 1-5 Gyr based on activity.  Assuming coevality of the companion, its age, gravity and metallicity can be constrained from properties of the primary, making it a suitable benchmark object for the calibration of evolutionary models and for determining the atmospheric properties of peculiar blue L dwarfs. The low total mass ($M_{tot}=0.21\\pm0.03$~\\msun), intermediate mass ratio ($q=0.45\\pm0.14$), and wide separation of this system demonstrate that the star formation process is capable of forming wide, weakly bound binary systems with low mass and BD components.  Based on the sensitivity of our search we find that no more than $2.2\\%$ of early-to-mid M dwarfs ($9.0 <M_V < 13.0$) have wide substellar companions with $m>0.06$~\\msun. ",
        "introduction": "\\label{sect:intro}  Although star birth is a complex process, the observation of binary systems--frequencies, mass ratios, and separations--can provide insight into the formation process as well as constraints for theoretical models. The formation of brown dwarfs (BDs) is particularly challenging since their masses are an order of magnitude smaller than the typical Jeans mass in molecular clouds.  Whether BDs form similarly to their more massive stellar counterparts or require additional mechanisms is currently an open question.  The answer may lie in the multiplicity properties of these substellar objects. Whereas free-floating BDs are observed in abundance, finding BDs as companions to stars has proved more difficult.  A ``brown dwarf desert'' ($\\lesssim0.5\\%$ companion fraction) is observed at close separations ($<3$~AU) to main sequence stars,  in comparison to a significant number of both planetary and stellar mass companions seen at similar separations \\citep{marcy00}.  It has recently been determined that this desert does not extend out to larger separations for solar analogs (F,G,K stars), $\\sim7\\%$ of which are found to harbor substellar companions at separations greater than 30~AU \\citep{met05}. However, searches for substellar companions to M dwarfs at large separations ($\\gtrsim$ 40 AU) have yielded mostly null results \\citep[e.g.][]{allen08,mccarthy04,hinz02,daemgen07} or sparse results \\citep[e.g.][]{oppen01}.  There are only 5 known BD companions to stellar M dwarf primaries at separations greater than 40AU: TWA 5b,c \\citep{lowrance99}, G 196-3B \\citep{rebolo98}, GJ 1001B \\citep{goldman99}, Gl 229B \\citep{nakajima95}, and LP 261-75B \\citep{burgasser05,reid06}.  For thoroughness we note that the L2.5 companion GJ 618.1B \\citep{wilson01} may also fall into this category, however it is more likely stellar.  In the VLM regime ($M_{1}<0.1$~\\msun) surveys have found that no more than $\\sim1\\%$ of stars have wide companions, including stellar ones \\citep{burgasser07a}.  Additionally, VLM binaries are found to be on average 10-20 times more tightly bound than their stellar counterparts, hinting that disruptive dynamical interactions may play an important role in their formation \\citep{close03}. These observations have been cited as evidence in favor of the ejection hypothesis \\citep{rep01, bate05} where BDs and VLM stars are thought to be stellar embryos formed by the fragmentation of a more massive pre-stellar core, then prematurely ejected from their birth environments.  However, BD companions to more massive stars do not tend to form harder binaries than stellar systems of similar total mass \\citep[e.g.][]{reid01,met05}.  While this is potentially evidence that BDs can form similarly to stars via turbulent fragmentation within molecular clouds \\citep{padoan02}, it is also consistent with simulations of disk instabilities \\citep[e.g.][]{stamatellos07,boss00}, which are capable of producing substellar companions around more massive primaries.  While a significant fraction of solar mass stars may retain wide BD companions, this does not seem to hold true for lower mass stars. As a result, very few wide BD companions to low mass stars are known. The discovery and characterization of these systems, especially in the intermediate range between the solar analog and VLM regimes, will help complete the emerging  picture of BD multiplicity at wide separations. Additionally, wide BD companions to stars make suitable ``benchmark'' objects, as their properties can be inferred from those of the primary \\citep{pinfield06}.  This is important for the calibration of BD evolutionary models, which requires independent age estimates.  Here we present the discovery of a wide substellar companion to a nearby M4.5 star.  The search, discovery and followup observations are outlined in \\S\\ref{sect:search}, while the physical properties of the system and its components are given in \\S \\ref{sect:properties}.  In \\S\\ref{sect:discussion} we discuss the companion's unusual NIR colors, possible formation scenarios, and the sensitivity of our search.  Given a space density for M dwarfs we make a crude estimate of how rare such systems may be. A brief summary and outlook are presented in \\S \\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:discussion}  \\subsection{The blue NIR colors of 2M1711+4028}\\label{sect:blue} The NIR colors of mid-late L dwarfs vary significantly within a single spectral type.  For L5 dwarfs there is a spread of $\\sim$0.7 magnitudes in $J-K_s$ \\citep{kirkpatrick08,cushing08}.  Although surface gravity and metallicity play a role, comparison of atmospheric models to actual spectra \\citep[e.g.][]{knapp04, cushing08, burgasser08} suggests that large variations in the NIR colors of L dwarfs are primarily related to the properties of condensate clouds in their atmospheres, with unusually red SEDs arising from thick clouds and blue ones from thin or large-grained clouds. Common to the known peculiar blue L dwarfs is exaggerated H2O absorption and diminished CO, as seen in the spectrum of 2M1711+4028. As discussed in \\S \\ref{sect:secondary} the discrepancy between the late-type H20 indices and the earlier type FeH and K1 indices, along with its unusually blue NIR colors are indications that 2M1711+4028 falls into this category.  For comparison we overplot 2M1711+4028's spectrum with the spectra of the very red L4.5 dwarf 2MASS J22244381-0158521 \\citep{kirkpatrick00}, and the relatively blue L5 optical standard 2MASS~J15074769-1627386 (see figure \\ref{fig:blue2}).  All spectra agree reasonably well in the J-band but diverge significantly at H and K, which may indicate differing properties of condensate clouds in their atmospheres. As a member of a wide binary system the surface gravity and metallicity of 2M1711+4028 can be constrained from properties of the primary, yielding an excellent laboratory for studying BD atmospheres.  The estimated age of 1-5~Gyr for the G 203-50 primary implies a relatively high surface gravity for the companion, possibly contributing the its blue NIR colors.  However surface gravity alone does not seem to be sufficient in explaining the NIR colors of peculiar blue L dwarfs \\citep{burgasser08}.  Additionally, the primary shows no signs of being metal poor, lending support to the hypothesis that unusually blue NIR colors can be primarily attributed to cloud properties. Higher resolution spectroscopy of the primary and a more precise determination of the metallicity is required in order to confirm this conclusion.  Another potential cause of unusually blue NIR colors is unresolved binarity.  This may explain the slight onset of CH4 absorption at 2.2 $\\mu$m in the spectrum of 2M1711+4028.  At least one of the known peculiar blue L dwarfs, 2MASS J08053189+4812330 \\citep{burgasser07b} is thought to be an unresolved binary with L4.5 and T5 components.  However, this system exhibits a pronounced dip in the 1.6 $\\mu$m CH4 feature due to the peaked shape of the T dwarf's SED, whereas the spectrum of 2M1711+4028 is relatively flat in that region.  Furthermore, the SED and colors of our companion are very similar to that of another blue L4.5 dwarf, 2MASS J11263991-5003550 (2M1126-5003 hereafter), discussed at length by \\citet{burgasser08}.  By constructing composite spectra using published L and T dwarf spectra from the SpeX prism library, \\citet{burgasser08} determined that no reasonable composite spectrum could be found that matched that of 2M1126-5003 in both the optical and NIR.  Without an optical spectrum for 2M1711+4028 we are limited in the conclusions we can draw, but its similarities to 2M1126-5003 may suggest that 2M1711+4028 is a single BD. Our AO images support this conclusion, indicating that the BD is not a near-equal mass binary with separation $> 0.18\\arcsec$.  \\subsection{Formation of G 203-50AB}\\label{sect:formation}  With a total mass of $\\sim$0.21~\\msun~G~203-50AB is slightly more massive than the rare wide VLM binaries, but much less massive than the solar analogues around which BDs are routinely found at wide separations (see \\S \\ref{sect:intro}, and figure \\ref{fig:sep}).  It is therefore of interest to consider how G~203-50AB may have formed. Could the secondary have formed through gravitational instability in a disk around the primary? Given the mass ratio of $q$=0.45, that would imply a $M_{disk} >$ 0.45$M_*$, whereas typical disks around low-mass stars contain a few percent of $M_*$ at $\\sim$1 Myr \\citep{scholz06}. Since it is unlikely that the entire disk would end up in the companion, the total disk mass, even in a conservative estimate, would have to be larger than the primary's own mass to start with. Thus, we conclude that formation of 2M1711+4028 in a protostellar disk around G~203-50 is implausible.  On the other hand, gravitational fragmentation of prestellar cores appears to be capable of forming a wide variety of binary systems, depending on the size, mass and angular momentum of the core \\citep[e.g.][]{bate00}. However, simulations usually have some difficulty producing binary stars with low component masses and wide separations \\citep[e.g.][]{bate03,goodwin04}. Some theoretical models invoke ejection from the parent cloud to halt further accretion that would otherwise lead to higher masses. Given its projected separation of 135$\\pm$25~AU, the G~203-50AB binary has a binding energy of 12.6$\\pm3.8~\\times$~$10^{-41}$~erg, placing it below the empirical ``minimum'' noted by \\citet{close03} and \\citet{burgasser07a}.  Thus, it is unlikely to have survived such an ejection.  We suggest that G~203-50AB most likely formed via fragmentation of an isolated core and did not suffer strong dynamical interactions during the birth process or subsequently.   \\subsection{Search Sensitivity}\\label{sect:sensitivity} In order to assess the sensitivity of our search we simulated proper motion distributions of M dwarfs distributed uniformly in a spherical volume out to 25 pc, with tangential space velocities drawn from the distribution of \\cite{schmidt07}.  To be sensitive to a particular M dwarf primary, its displacement between the 2MASS and SDSS surveys had to be greater than the $3\\sigma$ dispersion of all other stars in the 4~${\\rm deg^2}$ section of the sky in which it was found.  For each such section of the sky, the time baseline between the surveys was computed and used to determine the minimum proper motion required for a detection.  We assumed that the population of M dwarfs within 25 pc was uniformly distributed and assigned equal weight to each 4 ${\\rm deg^2}$ area of the sky.  Using the simulated proper motion distributions the fraction of M dwarfs whose proper motions we could have measured was determined for each section of the sky, giving an average fraction of 0.58.  Adopting an M dwarf space density ($9.0 < M_V < 13.0$ or roughly M0.5-M5.5) of $283.37 \\times10^{-4}~{\\rm pc^{-3}}$ \\citep{reid02}, and given a search area of 7668 ${\\rm deg^2}$ of the sky, we should have been sensitive to approximately $201\\pm12$ early-mid M dwarfs within 25 pc.  Although in some cases we were able to recover binary systems with separations $<$~$4\\arcsec$, we conservatively put a lower limit of $6\\arcsec$ on our sensitivity, ensuring that components are well separated.  The upper limit for separation is set by our search radius, which extended to $120\\arcsec$.  These limits correspond to projected separations of 30-600~AU at 5~pc and 150-3000~AU at 25~pc.  Our sensitivity to companions around each star was dictated by the mean 2MASS J-band limiting magnitude (S/N=10) of $\\sim16.5$, corresponding to a minimum mass of $\\sim0.06$~\\msun~at 25~pc, assuming an age of 1-5 Gyr.  Other factors preventing us from finding companions include poor astrometry due to saturation of the primary, or low S/N of the secondary.  To estimate the number of binaries missed we used SIMBAD and DwarfArchives\\footnote{http://DwarfArchives.org} to compile a list of 31 M-dwarfs and 38 BDs with previously measured proper motions large enough to pass our cuts, and tested whether we could measure the same proper motions using SDSS and 2MASS astrometry. We found that $91\\%$ of the time for M dwarfs, and $79\\%$ of the time for BDs  our measured proper motions agreed with the previously measured ones, using the same criteria as our matching algorithm described in \\S\\ref{sect:search}.  Therefore, we should have been capable of identifying approximately $72\\%$ of binaries with sufficiently high proper motions.  Correspondingly we adjust our sensitivity to $\\sim145\\pm9$ M dwarfs.  Adopting Poisson uncertainties on a $1\\sigma$ confidence interval for our single detection we roughly estimate that $0.7^{+1.5}_{-0.6}\\%$ of early-mid M dwarfs have substellar companions with masses greater than $\\sim0.06$~\\msun, at separations above $\\sim120$~AU."
    },
    "0808/0808.4159.txt": {
        "abstract": "\\noindent  We investigate further a model of the accreting millisecond X-ray pulsars we proposed earlier. In this model, the X-ray--emitting regions of these pulsars are near their spin axes but move. This is to be expected if the magnetic poles of these stars are close to their spin axes, so that accreting gas is channeled there. As the accretion rate and the structure of the inner disk vary, gas is channeled along different field lines to different locations on the stellar surface, causing the X-ray--emitting areas to move. We show that this ``nearly aligned moving spot model'' can explain many properties of the accreting millisecond X-ray pulsars, including their generally low oscillation amplitudes and nearly sinusoidal waveforms; the variability of their pulse amplitudes, shapes, and phases; the correlations in this variability; and the similarity of the accretion- and nuclear-powered pulse shapes and phases in some. It may also explain why accretion-powered millisecond pulsars are difficult to detect, why some are intermittent, and why all detected so far are transients. This model can be tested by comparing with observations the waveform changes it predicts, including the changes with accretion rate. ",
        "introduction": "\\label{sec:intro}  Highly periodic millisecond X-ray oscillations have been detected with high confidence in 22 accreting neutron stars in low-mass X-ray binary systems (\\mbox{LMXBs}), using the \\textit {Rossi X-ray Timing Explorer} (\\textit {RXTE}) satellite (see~\\citealt {lamb08a}). We refer to these stars as accreting millisecond X-ray pulsars (\\mbox{AMXPs}). Accretion-powered millisecond oscillations have so far been detected in 10 \\mbox{AMXPs}. They are always observable in seven \\mbox{AMXPs}, but are only intermittently detected in three others. Nuclear-powered millisecond oscillations have been detected with high confidence during thermonuclear X-ray bursts in 16 \\mbox{AMXPs}. Persistent accretion-powered millisecond oscillations have been detected in two \\mbox{AMXPs} that produce nuclear-powered millisecond oscillations; intermittent accretion-powered millisecond oscillations have been detected in two others.  The \\mbox{AMXPs} have several important properties:  \\textit {Low oscillation amplitudes}. The fractional amplitudes of the accretion-powered oscillations of most \\mbox{AMXPs} are often only $\\sim\\,$1\\%--2\\%.\\footnote{We characterize the strengths of oscillations by their rms amplitudes, because the rms amplitude can be defined for any waveform, is usually relatively stable, and is closely related to the power. We convert reported semi-amplitudes of purely sinusoidal oscillations or Fourier components to rms amplitudes by dividing by $\\sqrt2$.} Persistent accretion-powered oscillations with amplitudes $\\la\\,$1\\% are often detected with high confidence in IGR~J00291$+$5934 (\\citealt{gall05,patr08}) and XTE~J1751$-$305 (\\citealt{mark02, patr08}). Persistent accretion-powered oscillations with amplitudes as low as 2\\% are regularly seen in XTE~J1807$-$294 (\\citealt{zhan06, chou08, patr08}), XTE~J0929$-$314 (\\citealt{gall02}), and XTE~J1814$-$338 (\\citealt{chun08, patr08}). The amplitude of the accretion-powered oscillation seen in SWIFT~1756.9$-$2508 was $\\sim\\,$6\\% (\\citealt{krim07}). The intermittent accretion-powered oscillations detected in SAX~J1748.9$-$2021 (\\citealt{gavr07, alta08, patr08}), HETE~J1900.1$-$2455 (\\citealt{gall07}), and Aql~X-1 (\\citealt{case08}) have amplitudes $\\sim\\,$0.5\\%--3\\%.  \\textit {Nearly sinusoidal waveforms}. The waveforms (light curves) of the accretion-powered oscillations of most \\mbox{AMXPs} are nearly sinusoidal (see \\citealt{wijn06} and the references in the preceding paragraph). The amplitude of the first harmonic (fundamental) component is usually $\\ga\\,$10 times the amplitude of the second harmonic (first overtone) component, although the ratio can be as small as $\\sim\\,$3.5, as is sometimes the case in XTE~J1807$-$294 (\\citealt{zhan06}), or even $\\ga\\,$1, as is sometimes the case in SAX~J1808.4$-$3658 (see, e.g., \\citealt{hart08}).  \\textit {Highly variable oscillation amplitudes}. The fractional amplitudes of the accretion-powered oscillations of most \\mbox{AMXPs} vary in time by factors ranging from $\\sim\\,$2 to $\\sim\\,$10. Observed fractional amplitudes vary from 0.7\\% to 1.7\\% in SAX~J1748.9$-$2021, from 0.7\\% to 3.7\\% in XTE~J1751$-$305, from 3\\% to 7\\% in XTE~J0929$-$314, from 1\\% to 9\\% in IGR~J00291$+$5934, from 2\\% to 11\\% in XTE~J1814$-$338, from 1\\% to 14\\% in XTE~J1808$-$338, and from 2\\% to 19\\% in XTE~J1807$-$294 (see the references above).  \\textit {Highly variable pulse phases}. The phases of accretion-powered pulses have been seen to vary rapidly by as much as $\\sim\\,$0.3 cycles in several \\mbox{AMXPs}, including SAX~J1808.4$-$3658 \\citep{morg03, hart08} and XTE~J1807$-$294 \\citep{mark04}. Wild changes in the apparent pulse frequency have been observed with \\textit {both} signs at the \\textit{same} accretion rate in XTE~J1807$-$294 (see \\citealt{mark04}). If interpreted as caused by changes in the stellar spin rate, these phase variations would be more than a factor of 10 larger than expected for the largest accretion torques and smallest inertial moments thought possible for these systems (see \\citealt{ghos79b, latt01}).  \\textit {Undetected accretion-powered oscillations}. More than 80 accreting neutron stars in \\mbox{LMXBs} are known (\\citealt{chak05, liu07}), but accretion-powered millisecond X-ray oscillations have so far been detected in only 10 of them. Accretion-powered oscillations have not yet been detected even in 13 \\mbox{AMXPs} that produce periodic nuclear-powered millisecond oscillations, indicating that they have millisecond spin periods (\\citealt{lamb08a}); eight of these also produce kilohertz quasi-periodic oscillations (\\mbox{QPOs}) with frequency separations that indicate that they not only have millisecond spin periods but also have dynamically important magnetic fields (\\citealt{bout08a}).  \\textit {Intermittent accretion-powered oscillations}. Accretion-powered millisecond X-ray pulsations have been detected only occasionally in SAX~J1748.9$-$2021 (\\citealt{gavr07, alta08, patr08}), HETE~J1900.1$-$2455 (\\citealt{gall07}), and Aql~X-1 (\\citealt{case08}). When oscillations are not detected, the upper limits are typically $\\la0.5$\\%.  \\textit{Correlated pulse arrival times and amplitudes}. The phase residuals of the accretion-powered pulses of several \\mbox{AMXPs} appear to be anti-correlated with their fractional amplitudes, at least over some of the amplitude ranges observed. \\mbox{AMXPs} that show this type of behavior include XTE~J1807$-$294 and XTE~J1814$-$338 (\\citealt{patr08}).  \\textit{Similar accretion- and nuclear-powered pulses}. The shapes and phases of the nuclear-powered X-ray pulses of the \\mbox{AMXPs} SAX~J1808.4$-$3658 (\\citealt{chak03}) and XTE~J1814$-$338 (\\citealt{stro03}) are very similar to the  shapes and phases of their accretion-powered X-ray pulses.  \\textit {Concentration in transient systems of \\mbox{AMXPs} with accretion-powered oscillations}. The \\mbox{AMXPs} in which accretion-powered oscillations have been detected tend to be found in binary systems that have outbursts lasting about a month (but see \\citealt{gall08}) separated by quiescent intervals lasting years (\\citealt{chak05, rigg08}). The accretion rates of these neutron stars are very low.  In this paper we investigate further the ``nearly aligned moving spot'' model of \\mbox{AMXP} X-ray emission that we proposed previously \\citep{lamb06,lamb07,lamb08b}. This model has three main features:  \\begin{enumerate} \\item  The strongest poles of the magnetic fields of neutron stars with millisecond spin periods are located near---and sometimes very near---the stellar spin axis. This behavior is expected for several magnetic field evolution mechanisms.  \\item  The star's magnetic field channels accreting gas close to its spin axis, creating X-ray emitting areas there and depositing nuclear fuel there.\\footnote{\\cite{muno02} considered a single bright spot near the spin axis as well as a uniformly bright hemisphere and antipodal spots near the spin equator as possible reasons for the nearly sinusoidal waveforms of some X-ray burst oscillations, but did not consider accretion-powered oscillations or other consequences of emission from near the spin axis.}  \\item  The X-ray emitting areas on the stellar surface move, as changes in the accretion rate and the structure of the inner disk cause accreting gas to be channeled along different field lines to different locations on the stellar surface. (The magnetic field of the neutron star is fixed in the stellar crust on the timescales relevant to the phenomena considered here.) \\end{enumerate}  These features provide the basic ingredients needed to understand the \\mbox{AMXP} properties discussed above. This is the subject of the sections that follow. As a guide to these sections, we summarize our results here.  \\begin{enumerate} \\item  Emission from near the spin axis naturally produces weak modulation, regardless of the viewing direction. The reason is that uniform emission from a spot centered on the spin axis is axisymmetric about the spin axis and therefore produces no modulation. Emission from a spot close to the spin axis has only a small asymmetry and therefore produces only weak modulation.  \\item  Emission from near the spin axis also naturally produces a nearly sinusoidal waveform, because the asymmetry of the emission is weak and broad.  \\item  If the emitting area is close to the spin axis, even a small movement in latitude can change the oscillation amplitude by a substantial factor.  \\item  If the emitting area is close to the spin axis, movement in the longitudinal direction by a small distance can change the phase of the oscillation by a large amount.  Changes in the latitude and longitude of the emitting area are expected on timescales at least as short as the $\\sim\\,$0.1~ms dynamical time at the stellar surface and as long as the $\\sim\\,$10~d timescale of the variations observed in the mass accretion rate.  \\item  If the emitting area is very close to the spin axis and remains there, the oscillation amplitude may be so low that it is undetectable. The effects of rapid changes in the position of the emitting area---possibly in combination with other effects, such as reduction of the modulation fraction by scattering in circumstellar gas---may also play a role in reducing the detectability of accretion-powered oscillations in neutron stars with millisecond spin periods. These effects may explain the fact that accretion-powered X-ray oscillations have not yet been detected in many accreting neutron stars that are thought to have millisecond spin periods and dynamically important magnetic fields.  \\item  If the emitting area is very close to the spin axis, a small change in the accretion flow can suddenly channel gas farther from the spin axis, causing the emitting area to move away from the axis. This can make a previously undetectable oscillation become detectable. Temporary motion of the emitting area away from the spin axis may explain the intermittent accretion-powered oscillations of some \\mbox{AMXPs} (\\citealt{lamb09}).  \\item  If the pulse amplitude and phase variations observed in \\mbox{AMXPs} are caused by motion of the emitting area, they should be correlated. In particular, the pulse phase should be much more scattered when the pulse amplitude is very low. The reason is that changes in the longitudinal position of the emitting area by a given distance produce much larger phase changes when the emitting area is very close to the spin axis, which is also when the oscillation amplitude is very low.  The observational consequences discussed so far depend only on features~(2) and~(3) of the model, i.e., that the accretion-powered X-ray emission of \\mbox{AMXPs} comes from areas near their spin axes and that these areas move significantly on timescales of hours to days.  \\item  The picture of \\mbox{AMXP} X-ray emission outlined here suggests that the shapes and phases of the nuclear- and accretion-powered pulses are similar to one another in some \\mbox{AMXPs} because the nuclear- and accretion powered X-ray emission comes from approximately the same area on the stellar surface. The reason for this is that in some cases, the mechanism that concentrates the magnetic flux of the accreting neutron star near its spin axis, as it is spun up, will naturally produce magnetic fields strong enough to confine accreting nuclear fuel near the magnetic poles at least partially, even though the dipole component of the magnetic field is weak.  \\item  The picture of neutron star magnetic field evolution and \\mbox{AMXP} X-ray emission outlined here also suggests a possible explanation for why the \\mbox{AMXPs} in which accretion-powered oscillations have been detected are in transient systems. If most neutron stars in \\mbox{LMXBs} were spun up by accretion from a low spin rate to a high spin rate, their magnetic poles were forced very close to their spin axes, making accretion-powered oscillations difficult or impossible to detect. However, those stars that are now in compact transient systems now experience infrequent episodes of mass accretion and the accretion rate is very low. By now they have been spun down from their maximum spin rates, a process that could force their magnetic poles away from their spin axes enough to produce detectable accretion-powered oscillations. \\end{enumerate}  These last two observational consequences depend on feature~(1) of the model, i.e., on how the magnetic fields of neutron stars evolve as they are spun up and down by accretion and electromagnetic torques.  In the remainder of this paper we discuss in detail the features of the model and its observational implications. In Section~\\ref {sec:model}, we outline our approach, discussing our modeling of X-ray emission from the stellar surface, our computational and the code verification methods, and the pulse profile representation we use.  We present our results in Sections~\\ref {sec:amplitudes} and~\\ref {sec:variations}. These results are based on our computations of several hundred million waveforms for different emitting regions, beaming patterns, stellar models, and viewing directions. In Section~\\ref {sec:amplitudes}, we consider the shape and amplitude of X-ray pulses as a function of the size and inclination of the emitting areas, the compactness of the star, and the stellar spin rate. In Section~\\ref {sec:variations}, we consider the changes in the pulse amplitude and phase produced by various motions of the emitting regions on the stellar surface and explore the origins of correlated changes in the pulse amplitude and phase and the effects of rapid movement of the emitting areas. We also discuss why oscillations have not yet been detected in many accreting neutron stars in LMXBs and why accretion-powered oscillations are detected only intermittently in some \\mbox{AMXPs}.  In Section~\\ref {sec:discussion}, we summarize the results of our model calculations. We also discuss how the magnetic poles of most \\mbox{AMXPs} can be forced close to their spin axes, how such mechanisms may explain why the \\mbox{AMXPs} that produce accretion-powered millisecond oscillations are transient pulsars, the consistency of the model with the observed properties of rotation-powered millisecond pulsars, and possible observational tests of the model discussed here.  Further results of our investigation of the present model will be presented elsewhere (S. Boutloukos et al., in preparation). ",
        "conclusions": "\\label{sec:conclusions}  In previous sections we have explored in some detail the nearly aligned moving spot model of \\mbox{AMXP} X-ray emission. In this model the X-ray emitting regions are close to the stellar spin axis and move around on the stellar surface with time. Here we list our principal conclusions.  \\textit {Pulse amplitudes and shapes}. In Section~\\ref{sec:spot-inclination}, we investigated the amplitudes and shapes of the pulses produced by emitting spots on the stellar surface as a function of their inclination to the spin axis, for several X-ray beaming patterns and a range of stellar masses, compactnesses, and spin rates. We found that emitting areas on or near the stellar surface can produce fractional amplitudes as low as the 1\\%--2\\% values often observed only if they are located within a few degrees of the stellar spin axis. Regions near the spin axis also naturally produce nearly sinusoidal pulse profiles.  We explored effect of spot size on pulse amplitude in Section~\\ref{sec:spot-size}. We found that although the pulse amplitudes produced by large emitting areas tend to be smaller, this effect is weak. Unless almost the entire surface of the star is uniformly emitting, even large spots produce pulse amplitudes greater than those observed in the \\mbox{AMXPs}, unless the spots are centered close to the spin axis.  In Section~\\ref{sec:compactness}, we studied the effect of stellar compactness on pulse amplitudes. Although the pulse amplitudes produced by very compact neutron stars tend to be smaller than the amplitudes produced by less compact stars, we found that this effect is too weak to explain by itself pulse amplitudes as small as those observed in the \\mbox{AMXPs}. Stellar compactness clearly cannot explain why the fractional amplitudes of several \\mbox{AMXPs} are $\\sim\\,$1\\%--2\\% at some times but $\\sim\\,$15\\%--25\\% a few hours or days later, because the stellar compactness cannot change on such short timescales.  These results show that emission from the stellar surface can explain the low amplitudes and nearly sinusoidal waveforms typically observed in \\mbox{AMXPs} only if the emitting areas are located close to the stellar spin axis.  \\textit {Variability of pulse amplitudes, shapes, and arrival times}. In Section~\\ref{sec:amplitude-variations}, we investigated the amplitude changes that can be produced by motion of the emitting area on the stellar surface. We found that if the emitting area is close to the spin axis, even a small change in the latitude of the area can change the oscillation amplitude by a substantial factor. For example, changes in the inclination of the emitting area by $\\la\\,$10$\\arcdeg$ can explain the amplitude variations seen in the \\mbox{AMXPs} and the relatively large fractional amplitudes $\\sim\\,$15\\%--20\\% occasionally seen in some of them.  In Section~\\ref{sec:phase-variations}, we studied the changes in the arrival times (phases) of the harmonic components of the pulse caused when the emitting area moves around on the stellar surface. We found that changes in the latitude of the emitting area can shift the phases of the first and second harmonics by at least 0.15~cycles and by different amounts. We showed that if the emitting area is close to the spin axis and moves in the azimuthal direction by even a small distance, the phases of the first and second harmonics can easily shift by as much as \\mbox{$\\sim\\,$0.1}--0.4 cycles. If the emitting area loops the spin axis, the phases of the Fourier components will shift by more than one cycle.  Motion of the emitting area on the stellar surface generally produces both amplitude and phase variations. As discussed in Section~\\ref{sec:model}, the position of the emitting area is expected to reflect the accretion rate and structure of the inner disk, and is therefore expected to change on timescales at least as short as the $\\sim\\,$0.1~ms dynamical time at the stellar surface and as long as the $\\sim\\,$10~d timescale of the variations observed in the mass accretion rate. Changes in the position of the emitting area on timescales longer than the $\\sim\\,$$10^{2}$--$10^{3}$~s integration times required to construct a pulse profile will produce changes in the apparent pulse amplitude and phase.  These results show that if the emitting area is close to the spin axis, modest changes in its location can explain the rapidly varying harmonic amplitudes and phases of the \\mbox{AMXPs}.   \\textit {Correlated amplitude and phase variations}. We showed in Section~\\ref{sec:correlated-variations} that changes in the latitude and longitude of the emitting area tend to produce correlated changes in the amplitudes and phases of the harmonic components of the pulse. A strong expectation in the nearly aligned moving spot model is that the scatter in the pulse arrival times (i.e., the pulse time or phase residuals) should decrease steeply with increasing pulse amplitude. The residuals of several \\mbox{AMXPs}, including XTE~J1807$-$294, SAX~J1748.9-2021, and IGR~J00291$+$5934, behave in this way. The magnitudes of the phase residuals of these \\mbox{AMXPs} are consistent with the nearly aligned moving spot model if the emitting areas wander in the azimuthal direction by distances $\\sim\\,$$10^{-2}$--$10^{-3}$ times the stellar circumference.  A second expectation in the nearly aligned moving spot model is that the arrival times of pulses will form a track in the phase-residual versus pulse-amplitude plane, especially if the change in the pulse amplitude is large. Such a track will be formed if the emitting areas where the accreting matter impacts the stellar surface move repeatedly along a particular path in stellar latitude and longitude as the accretion rate and the structure of the inner disk change. This is expected because models of the flow of accreting matter from the inner disk to the stellar surface predict that matter will be guided along particular but different magnetic flux tubes as the accretion rate and the structure of the inner disk vary (see Section~\\ref {sec:modeling-emission}). As discussed in Section~\\ref{sec:correlated-variations}, tracks of this type are observed in plots of pulse arrival time versus pulse amplitude for XTE~J1814$-$338 and XTE~J1807$-$294.   \\textit {Accretion- and nuclear-powered oscillations}. The success of the emitting spot model discussed here supports magnetic field evolution models in which the magnetic flux of the accreting neutron star becomes concentrated near its spin axis as it is spun up. As discussed in Sections~\\ref{sec:accretion-nuclear} and~\\ref{sec:pole-movement}, these evolutionary models can produce magnetic fields strong enough to partially confine accreting nuclear fuel near the star's magnetic poles, even though the dipole component of the magnetic field is very weak. This picture of magnetic field evolution in turn suggests that the shapes and phases of the nuclear- and accretion-powered pulses are similar to one another in some \\mbox{AMXPs} because the nuclear- and accretion powered X-ray emission comes from approximately the same area on the stellar surface.   \\textit {Effects of rapid spot motions}. In Section~\\ref{sec:spot-movements}, we discussed the effects of spot movements on timescales shorter than the time required to construct a pulse profile. Such effects are expected, because emitting spots are likely to move on the stellar surface on timescales at least as short as the $\\sim\\,$0.1~ms dynamical time there whereas constructing a pulse profile usually requires folding $\\sim\\,$$10^{2}$--$10^{3}$~s of X-ray flux data. Rapid spot motions will produce X-ray flux variations on these same timescales. Our computations show that motion of the emitting area on the stellar surface on timescales longer than the spin period usually changes the amplitudes and the phases of the harmonic components of the theoretical pulse profile.  Variations of the X-ray flux on any timescales shorter than the time required to construct a pulse profile will appear in the analysis as noise in excess of the normal counting noise. This noise will reduce the measured amplitude of the oscillations at the spin frequency and its overtones for all spot locations, but its effect is likely to be stronger when the emitting area is near the spin axis because displacement of the emitting area by a given distance there produces a larger change in the pulse phase and, for many geometries, in the pulse amplitude. This effect will therefore tend to reduce the apparent pulse amplitude even further when the emitting area is close to the spin axis.   \\textit {Undetectable and intermittent pulsations}. In Section~\\ref {sec:spot-inclination}, we showed that if the emitting areas of some \\mbox{AMXPs} are very close to the spin axis and remain there, the amplitudes of the oscillations they would produce can be $\\sim\\,$0.5\\% or less, making them undetectable with current instruments. Rapid X-ray flux variations will make accretion-powered oscillations at the spin frequency more difficult to detect. Other effects, such as scattering of X-ray photons in circumstellar gas, may also play a role in reducing the detectability of such oscillations. These results show that the nearly aligned moving spot model may, possibly in combination with other effects, explain the nondetection of accretion-powered oscillations at the millisecond spin frequencies of some accreting neutron stars in which nuclear-powered oscillations have been detected. In Section~\\ref {sec:amplitude-variations}, we showed that if the emitting area is within a few degrees of the spin axis and moves toward the rotation equator by $\\sim\\,$10$\\arcdeg$, oscillations that were undetectable can become detectable. This may explain why accretion-powered oscillations appear only intermittently in some \\mbox{AMXPs} \\citep{lamb08a}. The model suggests that oscillations may also disappear intermittently in some \\mbox{AMXPs}.   \\textit {Evolution of \\mbox{AMXP} magnetic fields}. In Section~\\ref{sec:pole-movement}, we explained why the magnetic poles of most \\mbox{AMXPs} are expected to be very close to their spin axes. One consequence is that their magnetic fields will channel accreting gas to the stellar surface near the spin axis. Hence the star's accretion-powered X-ray emission will come from areas near the spin poles. A second consequence is that many \\mbox{AMXPs} may have surface magnetic fields as strong as $\\sim\\,$$10^{11}$--$10^{12}$~G, even though the dipole components of these fields are only $\\sim\\,$$10^{8}$--$10^{9}$~G. If so, many \\mbox{MRPs} may also have surface magnetic fields as strong as $\\sim\\,$$10^{11}$--$10^{12}$~G, even though the dipole components inferred from their spin-down rates are only $\\sim\\,$$10^{8}$--$10^{9}$~G.   \\textit {Transient nature of \\mbox{AMXPs}}. In Section~\\ref{sec:transients}, we noted that the picture of \\mbox{AMXP} magnetic field evolution just described suggests why the \\mbox{AMXPs} in which accretion-powered oscillations have been detected are in transient systems. If the magnetic poles of most neutron stars in \\mbox{LMXBs} were forced very close to their spin axes during the initial, persistent phase of mass transfer, accretion-powered X-ray oscillations would be difficult or impossible to detect. Later, when mass transfer is transient, the stars will spin down, causing their magnetic poles to move away from their spin axes and making accretion-powered X-ray oscillations detectable.   \\textit {Tests of the model}. The nearly aligned moving spot model leads to a number of expectations about \\mbox{AMXP} X-ray emission that can be tested (see Section~\\ref{sec:discussion} for details):  \\begin{enumerate} \\item  The amplitudes, harmonic content, and arrival times of pulses should be functions of the X-ray luminosity and spectrum of the pulsar.  \\item  The amplitudes, harmonic content, and arrival times of pulses should be correlated with one another.  \\item  The arrival times of pulses with low amplitudes should fluctuate much more than the arrival times of pulses with high amplitudes.  \\item  Pulse phase residuals are likely to form a track in the phase-residual versus pulse-amplitude plane. Such tracks are likely to be more evident if the range of the pulse amplitude variation is large. The position of pulses along such a track should be correlated with the X-ray luminosity and spectrum of the pulsar.  \\item  If the accretion- and nuclear-powered pulses of an \\mbox{AMXP} appear nearly identical, long-term (days-to-weeks) variations in the phase residuals of the two types of pulses should track one another.  \\item  The strength of the excess background noise produced by pulse shape fluctuations should be correlated with the amplitudes, harmonic content, and arrival times of pulses and the X-ray luminosity and spectrum of the pulsar.  \\item  The excess noise produced by motion of the emitting area should be stronger when the pulse amplitude is smaller and weaker when the pulse amplitude is larger.  \\item  If \\mbox{AMXPs} do have surface magnetic fields as strong as $\\sim\\,$$10^{11}$--$10^{12}$~G, their keV X-ray spectra may show strong-magnetic-field features. Such features are more likely in \\mbox{AMXPs} that show evidence of nuclear fuel confinement, such as SAX~J1808.4$-$3658 and XTE~J1814$-$338.  \\item  If \\mbox{AMXPs} have surface magnetic fields $\\sim\\,$$10^{11}$--$10^{12}$~G, their offspring should have millisecond spin periods and total magnetic fields of similar strength, but dipole fields $\\sim\\,$$10^{8}$--$10^{9}$~G. This should affect particle acceleration and $\\gamma$-ray emission by these neutron stars.  \\item  If the explanation of the transient nature of the \\mbox{AMXPs} suggested here is correct, they should be spinning down on long timescales. \\end{enumerate}   \\textit{Note added in manuscript:} This work was presented at the 2008 April Amsterdam Workshop on Accreting Millisecond Pulsars, where we emphasized a test of our nearly aligned moving spot model that could also have wider implications. We argued that the close similarity of the accretion-powered and burst oscillation waveforms in XTE~J1814$-$338 strongly suggests that the emitting regions that produce them are similar and collocated, and that if, as we had proposed previously, the phase wandering of the accretion-powered oscillations is caused by movement of the emitting region, then the phase of the burst oscillations should wander in the same way. After the Workshop, \\citet{watt08} investigated this possibility and found just such a correlation in the \\textit {RXTE} data on XTE~J1814$-$338. In particular, these authors found that during the main part of the two-month outburst of XTE~J1814$-$338, its burst oscillation not only has a waveform similar to that of the accretion-powered oscillation \\citep{stro03, watt05, watt06}, but is also phase-locked with it, and that the peak of the burst oscillation coincides with the peak of the soft component of the accretion-powered oscillation. This indicates that, as we had suggested, the accretion- and nuclear-powered emitting regions in this pulsar very nearly coincide, and that the simultaneous wandering of the arrival times of both oscillations by $\\sim\\,$1~ms ($\\sim\\,$0.3 in phase) during the outburst is due to wandering of the matter (and hence the fuel) deposition pattern on the stellar surface."
    },
    "0808/0808.0471.txt": {
        "abstract": "We have performed a census of circumstellar disks around brown dwarfs in the $\\sigma$~Ori cluster using all available images from the Infrared Array Camera onboard the {\\it Spitzer Space Telescope}. To search for new low-mass cluster members with disks, we have measured photometry for all sources in the {\\it Spitzer} images and have identified the ones that have red colors that are indicative of disks. We present five promising candidates, which may consist of two brown dwarfs, two stars with edge-on disks, and a low-mass protostar if they are bona fide members. Spectroscopy is needed to verify the nature of these sources. We have also used the {\\it Spitzer} data to determine which of the previously known probable members of $\\sigma$~Ori are likely to have disks. By doing so, we measure disk fractions of $\\sim40$\\% and $\\sim60$\\% for low-mass stars and brown dwarfs, respectively. These results are similar to previous estimates of disk fractions in IC~348 and Chamaeleon~I, which have roughly the same median ages as $\\sigma$~Ori ($\\tau\\sim3$~Myr). Finally, we note that our photometric measurements and the sources that we identify as having disks differ significantly from those of other recent studies that analyzed the same {\\it Spitzer} images. For instance, previous work has suggested that the T dwarf S~Ori~70 is redder than typical field dwarfs, which has been cited as possible evidence of youth and cluster membership. However, we find that this object is only slightly redder than the reddest field dwarfs in $[3.6]-[4.5]$ ($1.56\\pm0.07$ vs.\\ 0.93--1.46). We measure a larger excess in $[3.6]-[5.8]$ ($1.75\\pm0.21$ vs.\\ 0.87--1.19), but the flux at 5.8~\\micron\\ may be overestimated because of the low signal-to-noise ratio of the detection. Thus, the {\\it Spitzer} data do not offer strong evidence of youth and membership for this object, which is the faintest and coolest candidate member of $\\sigma$~Ori that has been identified to date. ",
        "introduction": "\\label{sec:intro}    The lifetime of an accretion disk around a young star represents a fundamental constraint on the amount of time available for the formation of giant planets. The typical lifetimes of disks are estimated by comparing the prevalence of disks among clusters that span a range of ages ($\\tau\\sim1$--10~Myr). By measuring disk fractions as a function of stellar mass and star-forming environment as well as age, the influence of these two factors on disk lifetimes can be characterized.  Disk fractions are usually measured through infrared (IR) photometry of young clusters ($\\lambda>3$~\\micron) and the identification of the stars that exhibit excess emission from cool dust. Observations of this kind were first performed with the {\\it Infrared Astronomical Satellite} and ground-based near-IR cameras \\citep{kh95,hai01} and have made rapid progress in recent years with the {\\it Spitzer Space Telescope} \\citep{wer04}. Because {\\it Spitzer} is exceptionally well-suited for detecting disks around members of young clusters \\citep{all04,gut04,meg04,muz04}, it has been used extensively for measuring the disk fractions of stars \\citep{meg05,lada06,car06,sic06,her07,her07b,her08,mue07,dahm07,bar07,dam07,luh08cha}. {\\it Spitzer}'s unique sensitivity to faint IR sources has enabled measurements of disk fractions well into the substellar regime in nearby star-forming regions \\citep[$d=150$--300~pc,][]{luh05frac,luh06tau2,luh08cha,gui07}.     The cluster of young stars associated with the O star $\\sigma$~Ori~A has been thoroughly searched for low-mass stars and brown dwarfs \\citep[e.g.,][]{bej01,bar02,mar01,zap02a,zap02b}. As a result, it is an attractive target for measuring the disk fraction among low-mass objects. Several recent studies have sought to do this with {\\it Spitzer}. \\citet{her07} performed {\\it Spitzer} imaging of most of the $\\sigma$~Ori cluster and measured the disk fraction as a function of mass down to $\\sim0.1$~$M_\\odot$ for a sample of $\\sim300$ probable members (see also \\citet{cab06}). \\citet{cab07} identified a sample of brown dwarf candidates from optical and near-IR data and used the {\\it Spitzer} images from \\citet{her07} to estimate a disk fraction for those sources, arriving at a value of $\\sim50$\\%. Through further analysis of those {\\it Spitzer} data, \\citet{zap07} found that 6 of 12 brown dwarf candidates exhibited excess emission at 8~\\micron. \\citet{sch08} obtained deeper {\\it Spitzer} images of 18 of the faintest candidate members and reported a disk fraction of 29\\%. In addition to measuring disk fractions, the {\\it Spitzer} data have been used to assess the youth and membership of the coolest candidate member of the cluster, the T dwarf S~Ori~70 \\citep{zap02a}. \\citet{zap08} found that it is redder than typical field dwarfs in $[3.6]-[4.5]$ in the images from \\citet{her07}. Similarly, \\citet{sch08} suggested that S~Ori~70 could be anomalously red in $[3.6]-[4.5]$ and $[3.6]-[5.8]$ based on their deeper images, which they attributed to a disk or low surface gravity. In either case, the apparent color excesses would comprise evidence of youth for this object, whose membership in the $\\sigma$~Ori cluster has been questioned \\citep{mar03,bur04}.     The substellar population of the $\\sigma$~Ori cluster extends down to and below the detection limits of the images that have been obtained by {\\it Spitzer}. Thus, determining whether these objects have disks is a challenging task. For instance, \\citet{zap07} and \\citet{sch08} disagree on whether the {\\it Spitzer} data show evidence of disks for half of the sources considered by both studies. In this paper, we seek to accurately characterize the disk population among low-mass stars and brown dwarfs in $\\sigma$~Ori through an analysis of all {\\it Spitzer} images of this cluster that includes careful treatment of errors and biases. We begin by summarizing the {\\it Spitzer} observations and our data reduction methods (\\S~\\ref{sec:obs}). We then use these data to search for new low-mass members of $\\sigma$~Ori that have disks (\\S~\\ref{sec:select}) and to estimate the disk fraction for the substellar members of the cluster (\\S~\\ref{sec:pop}). ",
        "conclusions": "\\label{sec:conc}   We have investigated the disk population among stars and brown dwarfs in the $\\sigma$~Ori cluster using mid-IR images obtained with IRAC onboard the {\\it Spitzer Space Telescope}. For this study, we have employed all available IRAC data for $\\sigma$~Ori, which consist of shallow images (80.4~s) from \\citet{her07} and deep images ($\\sim1100$~s) from \\citet{sch08}. We measured photometry for all sources detected in these images and searched the resulting data for new members of the cluster based on the red colors that are expected from stars with disks. The five most promising candidates have colors and magnitudes that are suggestive of edge-on disks, brown dwarfs, and a low-mass protostar. We then examined the IRAC colors for $\\sim300$ probable cluster members found in previous studies and identified the ones that are likely to have disks. In doing so, we have attempted to fully account for the errors and biases in the photometry of the brown dwarf candidates \\citep[e.g., flux overestimation at low SNR;][]{bei03}, some of which are near the detection limits of the IRAC data. S~Ori~60 is the faintest candidate member of $\\sigma$~Ori that exhibits significant IR excess emission. This object is comparable in luminosity to the faintest brown dwarf that shows evidence of a disk in other star-forming regions \\citep{luh08cha}. Using our classifications of the IRAC data, we computed the disk fraction as a function of $M_J$, which acts as a proxy for stellar mass. The disk fractions for low-mass stars (0.08--0.5~$M_\\odot$) and brown dwarfs (0.01--0.08~$M_\\odot$) are $\\sim40$\\% and $\\sim60$\\%, respectively, which are similar to the disk fractions derived from IRAC surveys of two other clusters near the same age, IC~348 and Chamaeleon~I \\citep{mue07,luh08cha}.  Although our disk fraction for brown dwarfs in $\\sigma$~Ori is similar to other estimates based on the IRAC data considered here \\citep{cab07,zap07,sch08}, our photometric measurements and our classifications of the IRAC colors (disk vs.\\ no disk) differ significantly from those in previous studies. For instance, \\citet{zap08} and \\citet{sch08} suggested that the T dwarf S~Ori~70 may exhibit color excesses in $[3.6]-[4.5]$ and $[3.6]-[5.8]$ relative to field dwarfs at the same spectral type. They cited these non-standard colors as possible evidence of youth, which would indicate that S~Ori~70 is a cluster member rather than a field dwarf.  However, we have found that this object is not significantly redder in $[3.6]-[4.5]$ than field dwarfs and that its SNR may be too low at 5.8~\\micron\\ for a useful measurement of $[3.6]-[5.8]$. Therefore, we conclude that the IRAC data do not provide firm constraints on the membership of S~Ori~70."
    },
    "0808/0808.3436_arXiv.txt": {
        "abstract": "We observed the brightest part of HESS J1825$-$137 with the Suzaku XIS, and found diffuse X-rays extending  at least up to $\\timeform{15'}$ ($\\sim 17$~pc) from the pulsar PSR~J1826$-$1334. The spectra have no emission line, and are fitted with an absorbed power-law model. The X-rays, therefore, are likely due to synchrotron emission from a pulsar wind nebula. The photon index near at the pulsar ($r\\leq1.5'$) is  $\\Gamma=1.7$ while those in $r=1.5-16'$ are nearly constant at $\\Gamma=2.0$. The spectral energy distribution of the Suzaku and H.E.S.S. results are naturally explained by a combined process; synchrotron X-rays and $\\gamma$-rays by the inverse Compton of the cosmic microwave photons by  high-energy electrons in a magnetic field of $\\sim 7~\\mu$G. If the electrons are accelerated at the pulsar, the electrons must be transported over 17~pc in the synchrotron life time of 1900~yr, with a velocity of $\\geq 8.8 \\times 10^3$~km~s$^{-1}$. ",
        "introduction": "PSR~J1826$-$1334 was discovered by the Jodrell Bank radio survey. The spin period ($P$) is 101~ms and the period derivative ($\\dot{P}$) is $7.5\\times10^{-14}~\\mathrm{s~s^{-1}}$, implying a spin-down luminosity ($\\dot{E}$) of $2.8 \\times 10^{36} \\mathrm{~erg~s^{-1}}$ and a characteristic age ($\\tau$) of 21~kyr~\\citep{Clifton1992}. The distance to PSR~J1826$-$1334 has been measured as $3.9 \\pm 0.1$~kpc by the dispersion measure of its radio pulses~\\citep{Taylor1993,Cordes2002}.  From the spin-down luminosity and characteristic age, PSR~J1826$-$1334 may be a member of the Vela-like pulsars, in the evolutionary stage of the high-energy emission processes dominance. There are pulsed non-thermal and blackbody emission from a pulsar, an extended nebulosity powered by pulsar wind  (Pulsar Wind Nebula; PWN), and thermal and/or non-thermal emission from a supernova remnant (SNR)~\\citep{Bec97,Gaensler2003PSRB1823}.  Extended X-rays around PSR~J1826$-$1334 had been found with the ROSAT~\\citep{Finley1996} and ASCA~\\citep{Sakurai2001} observations, possibly due to PWN, but no extended radio emission was found ~\\citep{Bra89,Gae00}.  XMM-Newton resolved the extended X-rays into two parts; \"core\", a bright and elongated emission in the east-west direction with an extent of \\timeform{30\"}, and  \"diffuse\", a larger scale with lower surface brightness emission extending by $\\sim \\timeform{5'}$  mostly to the south of the pulsar \\citep{Gaensler2003PSRB1823}. Spectra of both components can be fitted with an absorbed power-law model, and hence are suggested to be a synchrotron origin. The photon index of the core  and diffuse components of  $\\Gamma \\sim 1.6$ and harder $\\Gamma$ of $\\sim 2$, respectively. The absorbing column densities $N_{\\mathrm{H}}$ are both about $1 \\times 10^{22}~\\mathrm{cm^{-2}}$. \\citet{Gaensler2003PSRB1823} interpreted both the components as parts of a single PWN and designated as G18.0$-$0.7. They argued that the non-detection of the PWN in the radio band is due to the limited  sensitivity for a nebula of this large size.  The one-sided morphology of the diffuse component suggests that the PWN is extended exclusively to the south of the pulsar.  The Chandra observation~\\citep{Pavlov2008} confirmed the existence of the two components, and measured the X-ray spectrum of the pulsar itself for the first time. The photon index of the pulsar is $\\Gamma \\sim 2.4$, and the luminosity is $8 \\times 10^{31}~\\mathrm{erg~s^{-1}}$ in the 0.5--8 keV band.  HESS~J1825$-$137 is one of the handful Very High Energy (VHE) $\\gamma$-ray sources on the Galactic plane discovered during the H.E.S.S. Galactic plane survey~\\citep{Aharonian2005a, Aharonian2005b, Aharonian2006Survey}. It is extending to the south of the radio pulsar PSR~J1826$-$1334. Subsequently, deep follow-up observation was made with H.E.S.S. ~\\citep{Aharonian2006HESSJ1825}  The $\\gamma$-ray emission is largely extended to $\\sim \\timeform{1D}$, i.e. $\\sim 70$~pc in size at the distance of 4~kpc. It has a larger scale one-sided morphology than X-rays toward the south of the pulsar. However, the detailed morphology is different from the X-ray; the surface brightness decreases and $\\gamma$-ray  spectrum shows  softening with the distance from the pulsar.  These suggest that the HESS~J1825$-$137 is a PWN powered by PSR~J1826$-$1334.  The most probable scenario for the VHE $\\gamma$-ray emission is inverse Compton (IC) scattering of the cosmic microwave background (CMB) photons with ultra-relativistic electrons accelerated by the pulsar, which is also responsible for the X-ray emission via the synchrotron process. However, the VHE $\\gamma$-rays are more extended than the X-rays. Furthermore, the X-rays are weak at the peak position of the $\\gamma$-rays. If the $\\gamma$-rays have the electron origin (IC), the typical energy of electrons is $\\sim$20~TeV. On the other hand, the typical energy of electrons emitting synchrotron X-rays is $\\sim$100~TeV assuming a magnetic field of $B=10~\\mu$G. These apparent differences between X-rays and $\\gamma$-rays may play a key role for the PWNe physics, such as the energy transport mechanism of and the history of energy injections to electrons.  The extent of the diffuse X-ray emission, however, may be underestimated, due mainly to the limited sensitivity XMM-Newton and Chandra for diffuse X-rays.  We hence observed the PWN with the X-ray Imaging Spectrometers~\\citep{Koyama2007XIS} on board the Suzaku satellite~\\citep{Mitsuda2007Suzaku}.  In this paper, we will introduce the observations in section 2, our analysis and results are explained in section 3, and we will discuss our results in section 4.  Uncertainties are quoted at the 90\\% confidence range unless otherwise stated. ",
        "conclusions": "\\subsection{Radial Profile}  In figure~\\ref{fig:radialprofile}, we see largely extended emission up to \\timeform{15'} for the first time. The radial profile becomes flat beyond \\timeform{10'}. The flux in this flat region is 50\\% larger than the mean flux of the background region. One may argue that the difference could be due to the fluctuations of the NXB,  CXB and GRXE. We therefore discuss  whether any  systematic errors of the background subtraction (NXB, CXB and GRXE) may mimic such excess.  The reproductively of the NXB is extensively examined by \\citet{Tawa2008}. Using the same method, the NXB reproductively in the 1.0--9.0 keV band at the 99\\% confidence level is less than $\\mathrm{2.2 \\times10^{-5}~counts~s^{-1}~arcmin^{-2}}$. This uncertainty is very  small compared with the excess level of the radial profile of the source region (figure~\\ref{fig:radialprofile}).  The fluctuation of the CXB in the XIS FOV is estimated to be $\\sim$ 36\\% at the 99\\% confidence level with the same way as \\citet{Tawa2008}.  Since the CXB flux is only 10\\% of the background spectrum, the overall uncertainty is less than 4\\%.  The largest uncertainty would come from the subtraction of the GRXE. \\citet{Yam93} and \\citet{Kan97} reported that the GRXE has a large scale structure depending on the galactic longitude ($l$) and latitude ($b$). The $l$-dependency is rather null in this region hence can be neglected in our case. The $b$-dependency is expressed by the scale height of the medium and hard components given by 2 and 0.7 degrees, respectively \\citep{Yam93, Kan97}. Since the flux of the medium and hard components in our region is nearly the same in the 1.0-9.0 keV band, we can estimate that GRXE flux at the source position is about 20\\% larger than that in the background position.  The most unknown factor is a possible fluctuation of the GRXE from position to position, which may cause under/over subtraction of the GRXE. If there is under/over subtraction of the GRXE, the source spectra should have emission/absorption line structures at the Si, S and Fe K-shell transition energies, because the GRXE has strong lines of these elements (see figure \\ref{fig:spectrumBG}). No such line/absorption structures are found in the spectra of the regions B, C, and D given in figure 5, which supports that the GRXE is properly subtracted. We further checked possible under/over subtraction of the GRXE by fitting the source spectrum (combined B+C+D) with a model of absorbed power-law + $\\Delta\\times$ GRXE, where the spectral parameters of the GRXE, other than normalization ($\\Delta$) were fixed to the best-fit values given in table~\\ref{tab:BGfitting}. As a result, we can set the 90\\% upper-limit of the possible contamination of the GRXE to be 1.5\\% of the excess flux. Accordingly, the largely extended power-law component from the source region is real, not an artifact of improper GRXE subtraction.   \\subsection{Comparison with earlier X-ray observation} In the XMM-Newton observation, \\citet{Gaensler2003PSRB1823} reported that the photon indices of the core and diffuse components are $\\sim$ 1.6 and $\\sim$ 2, respectively. Since the region A includes both the components, our result of $\\Gamma \\sim 1.7$ is reasonable.  From the region B, we determined more accurate  photon index and interstellar absorption than the diffuse component of XMM-Newton. Moreover, we detected the X-ray emission farther out of the previously reported region. The peak position of the extended emission coincides with the pulsar PSR~J1826$-$1334, and the surface brightness decreases with the distance from the pulsar. The major fraction of the extended emission cannot be a thin thermal plasma of the SNR, but is non-thermal origin. The photon indices are smaller than the canonical value of synchrotron emissions found in non-thermal SNRs (2.5--3.0; Bamba~et al.~2005).  Furthermore, the morphology of the extended emission is different from limb-brightened structures such as SN1006~\\citep{Koyama1995}.  Thus the extended emission cannot be explained as non-thermal emission from shell-like SNR. The photon indices are close to typical values of PWNe~\\citep{Gaensler2006}. Therefore, we conclude that the extended emission is synchrotron radiation from a PWN, as suggested first by \\citet{Gaensler2003PSRB1823}.  Our results indicate that the PWN extends over $\\timeform{15'}$, which corresponds to 17$d_\\mathrm{4kpc}$~pc, where $d_\\mathrm{4kpc}$ is the distance of PSR~J1826$-$1334 normalized by 4~kpc. The interstellar absorption of $N_{\\mathrm{H}} \\sim 10^{22} \\mathrm{cm}^{-2}$ is consistent with the distance of $\\sim 4$~kpc obtained by the radio measurements~\\citep{Taylor1993,Cordes2002}. The extent of $\\sim 17d_{\\rm 4kpc}$~kpc is larger than typical X-ray PWNe; for example, one of the largest PWNe is 3C58, where the X-ray nebula extends up to $\\sim 6$~pc from the pulsar~\\citep{Slane2004}.  The total X-ray luminosity ($L_{\\rm X}$) from the regions A--D are $8.6 d_{\\rm 4kpc}^2 \\times 10^{33}$~erg~s$^{-1}$ in the 2--10 keV band. The ratio of $L_{\\rm X}$ to the spin-down luminosity is not unusual compared with other PWNe~\\citep{Che04,Li07}.   \\subsection{Comparison with $\\gamma$-ray observation}  Since the region C overlaps the bright region of the VHE $\\gamma$-rays, we can make  a reliable spectral energy distribution (SED) between the X-rays and the VHE $\\gamma$-rays for the first time. In figure~\\ref{fig:SED}, we show the X-ray SED of the region C and the $\\gamma$-ray SED of the brightest region (Radius \\timeform{0.15D} in table~2 of \\cite{Aharonian2006HESSJ1825}).  If the origin of the VHE $\\gamma$-rays is the inverse Compton scattering of the CMB photons by high-energy electrons, these electrons should emit synchrotron radiation. We also plot the estimated synchrotron intensity in the magnetic fields of $B= 10,~5,~1~\\mu$G.  In the case of $B \\sim 7~\\mu$G, the estimated synchrotron intensity smoothly connects to the X-ray intensity.  Thus both the X-rays and VHE $\\gamma$-rays can be explained by high-energy electrons of a single population in a magnetic field of $B \\sim 7~\\mu$G, approximately the same as the core region of $B \\sim 10~\\mu$G ~\\citep{Gaensler2003PSRB1823}.  \\begin{figure} \\begin{center} \\FigureFile(80mm,60mm){figure6.eps} \\end{center} \\caption{ Spectral energy distribution of the extended component from the X-ray to TeV $\\gamma$-ray bands. The intensity of the region C in table~\\ref{tab:fittings} and that of the Radius \\timeform{0.15D} region in table~2 of \\citet{Aharonian2006HESSJ1825} are used for X-ray and VHE $\\gamma$-ray. The synchrotron radiation from the electrons responsible for VHE $\\gamma$-ray (0.25$-$10~TeV) is plotted toward the left for three different values of the magnetic field. \\label{fig:SED}} \\end{figure}  According to \\citet{deJager2008}, required electron energy $E_\\mathrm{e}$ in a magnetic field $B$ to radiate synchrotron photons of mean energy $E_\\mathrm{syn}$ is given by $E_\\mathrm{e}$ $=$ $120~\\mathrm{TeV}~(B/7\\mu \\mathrm{G})^{-1/2} (E_\\mathrm{syn} / \\mathrm{2keV})^{1/2}$. On the other hand, the mean electron energy $E_\\mathrm{e}$ required to inverse Compton scatter the CMB photons to energies $E_\\mathrm{IC}$ is typically $E_\\mathrm{e}=18~\\mathrm{TeV}~(E_\\mathrm{IC} /\\mathrm{1TeV})^{1/2}$ \\citep{deJager2008}. Thus electrons responsible for X-rays have larger energies than those for VHE $\\gamma$-rays.  The peak position of the VHE $\\gamma$-rays does not coincide with the pulsar~\\citep{Aharonian2006HESSJ1825}, while the peak position in X-rays are consistent with the pulsar. There is no indication of the $\\gamma$-ray peak in the XIS images (figures~\\ref{fig:image} and \\ref{fig:radialprofile}). The offset is a mystery, and the reason is not yet clarified. However, as \\citet{Pavlov2008} mentioned, the local excess of the seed photons for the inverse Compton scattering created by the near hidden star might cause the offset. In this case, a magnetic field stronger than $\\sim 10~\\mu$G might be required.  \\citet{Har99} reports an EGRET source 3~EG J1826--1302 in the north of PSR~J1826--1334, and the relation between them has been discussed  by some authors (e.g. Zhang, Chen \\& Fang~2008). The EGRET source is , however, out of the FOV of the XIS in our observation. Thus direct comparison between the intensities of the EGRET source and the X-ray we found is impossible, and the relation between them is not clarified.    \\subsection{Widely extended PWN}  Except for the neighborhood of the pulsar, the photon index is constant at $\\Gamma=2.0$ from the pulsar to $\\timeform{15'}$. The smooth decay of the intensity with the distance from the pulsar shown in figure~\\ref{fig:radialprofile} suggests that the accelerator is the pulsar itself.  This means the high-energy electrons should be transported over $17d_\\mathrm{4kpc}$~pc within its synchrotron life time. According to \\citet{deJager2008}, the synchrotron life time of an electron emitting photons of mean energy $E_\\mathrm{syn}$ in an isotropic magnetic field of the strength $B$ is $\\tau_\\mathrm{syn}$ $=$ $1.9~{\\rm kyr}~(B/7\\mu \\mathrm{G})^{-3/2} (E_\\mathrm{syn} / \\mathrm{2keV})^{-1/2}$.  Thus, the velocity for transporting the high-energy electrons should be more than $17 d_\\mathrm{4kpc}$~pc$/ \\tau_\\mathrm{syn}$ $\\sim$ $8.8 \\times 10^3~{\\rm km~s}^{-1} d_\\mathrm{4kpc} (B/7\\mu \\mathrm{G})^{3/2} (E_\\mathrm{syn} / \\mathrm{2keV})^{1/2}$.  A simple transport mechanism is diffusion by a magnetic field or convection. In the case of diffusion, our results indicate that the diffusion coefficient is $D=R^2/(2\\tau) = 2.3 \\times 10^{28} \\mathrm{~cm^2~s^{-1}} (R/\\mathrm{17pc})^2 (\\tau/1.9\\mathrm{kyr})^{-1}$, where $R$ is the size of the extended emission, and $\\tau$ is the transport time for which we assume the synchrotron life time.  We can express the mean free path of the electron in a magnetic field $B$ as $f r_\\mathrm{L}$, where $r_\\mathrm{L} = E_e/(eB)$ is the gyro radius of an electron with energy $E_e$ and charge $e$, and the parameter $f$ characterizes the efficiency of diffusion. Then the diffusion coefficient can also be expressed as $D$ $\\sim$ $f r_\\mathrm{L} c /3$ $=$ $2.3\\times 10^{28} \\mathrm{~cm^2~s^{-1}}~(B/7\\mu \\mathrm{G})^{-1}~(E_e/120\\mathrm{TeV}) (f/40)$. Thus, our results suggest $f \\sim 40$. According to \\citet{deJager2008}, $f$ should be $\\le 1$ in perpendicular diffusion. Therefore, our results suggest that the transport mechanism is not the perpendicular diffusion, but it might be diffusion parallel to the magnetic field or convection.  Collisions between a PWN and a reverse shock from the surrounding SNR were proposed as a scenario to explain the morphology of one-sided PWNe \\citep{Pavlov2008}, such as Vela pulsar~\\citep{Blondin2001, Gaensler2003PSRB1823}. Whether or not this scenario can explain the large extent of $\\sim$17$d_{\\rm 4kc}$~pc and the smooth radial profile shown in figure~\\ref{fig:radialprofile} is an open problem."
    },
    "0808/0808.3957_arXiv.txt": {
        "abstract": "Measuring g-mode pulsations of isolated white dwarfs can reveal their interior properties to high precision. With a spectroscopic mass of $\\approx 0.51 M_{\\odot}$ ($\\log g = 7.82$), the DAV white dwarf HS 1824+6000 is near the transition between carbon/oxygen core and helium core white dwarfs, motivating our photometric search for additional pulsations from the Palomar 60-inch telescope.  We confirmed (with much greater precision) the three frequencies: $2.751190 \\pm 0.000010$ mHz (363.479 sec), $3.116709 \\pm 0.000006$ mHz (320.851 sec), $3.495113 \\pm 0.000009$ mHz (286.114 sec), previously found by B. Voss and collaborators, and found an additional pulsation at $4.443120 \\pm 0.000012$ mHz (225.067 sec). These observed frequencies are similar to those found in other ZZ Ceti white dwarfs of comparable mass (e.g. $\\log g<8$).  We hope that future observations of much lower mass ZZ Ceti stars ($< 0.4 M_{\\odot}$) will reveal pulsational differences attributable to a hydrogen covered helium core. ",
        "introduction": "\\begin{deluxetable}{lcc} \\tablewidth{0pt} \\tablecaption{Properties of HS 1824+6000\\label{tab:prop}} \\tablehead{ \\colhead{ } & \\colhead{\\citealt{bv06}} & \\colhead{\\citealt{ag07}} } \\tablecolumns{3} \\startdata $T_{\\rm eff}$ (K) & $11192\\pm300$\\tablenotemark{a} & $11380\\pm140$\\tablenotemark{b} \\\\ $\\log g$ (dex) & $7.65\\pm0.10$\\tablenotemark{a} & $7.82\\pm0.04$\\tablenotemark{b} \\\\ $m_B$ (mag) & 15.7 & 15.7 \\\\ \\cutinhead{Observed Frequencies (in mHz)} \\citealt{bv06} & \\multicolumn{2}{l}{$2.6\\pm0.4$, $3.0\\pm0.9$, $3.3\\pm0.4$, $3.4\\pm0.8$} \\\\ This Paper & \\multicolumn{2}{l}{$2.751190 \\pm 0.000010$, $3.116709 \\pm 0.000006$} \\\\ & \\multicolumn{2}{l}{$3.495113 \\pm 0.000009$, $4.443120 \\pm 0.000012$} \\enddata \\tablenotetext{a}{Photometrically determined.} \\tablenotetext{b}{Spectroscopically determined.} \\end{deluxetable}  \\begin{figure*}[t] \\centering \\epsscale{1.0} \\plotone{f1.eps} \\caption{The empirical ZZ Ceti instability strip.  The circles represent systems for which temporal observations have been performed.  Filled circles indicate systems not observed to vary while open circles indicate systems with observed periods.  These data are from \\citet{pb04} and \\citet{ag05,ag07}.  The vertical crosses represent low-mass WDs selected from the SDSS with $\\log g$ and $T_{\\rm eff}$ redetermined from MMT spectra from \\citet{mk07a}.  The diagonal crosses are from \\citet{mk07a} except they are reanalysis of SDSS spectra and are merely candidate low-mass WDs until better spectra can be obtained.  The square is LP 400-22 \\citep{ak06,mk07a}.  The star is HS 1824+6000 (see Table \\ref{tab:prop}, \\citealt{ag07}).  The error bar in the instability strip represents the typical error for measurements within the instability strip.  The dashed lines are empirical fits to the instability strip as determined by \\citet{ag07}.  The solid \\citep{jap07} and dash-dotted \\citep{lga01} lines correspond to cooling tracks of He WDs of the labeled mass.  The labeled masses correspond to the following model masses (in $M_{\\odot}$) for \\citet{jap07} and \\citet{lga01} respectively: 0.19 to 0.1869 and 0.196, 0.24 to 0.2495 and 0.242, 0.40 to 0.3986 and 0.406, 0.45 to 0.4481 (\\citealt{jap07} only).} % \\label{fig:instrip} \\end{figure*}  The hydrogen line ZZ Ceti variable (DAV) white dwarfs (WDs) occupy a discrete strip in the $T_{\\rm eff}-\\log g$ plane known as the ZZ Ceti instability strip.  Many groups have assessed the location of this instability strip both empirically \\citep{fw91,asm04b,ag05,bgc07} and theoretically \\citep{pb97,yw99,gf03}, and despite minor discrepancies it spans $11,000 \\lesssim T_{\\rm eff} \\lesssim 12,250$ K for $\\log g \\approx 8.0$.  \\citet{gf82} suggested that the instability strip is pure, meaning that all WDs within the strip are variable.  However, \\citet{asm05} found numerous objects from the SDSS with associated low signal-to-noise spectra that were tenuously identified to be in the strip but did not vary to their observed detection limits.  Recent observations \\citep{bgc07} have found some of these to be low amplitude pulsators.  If the instability strip is pure, then it strongly implies that ZZ Ceti stars are a phase of evolution through which all DA WDs must evolve as they cool.  White dwarfs less massive than $\\approx 0.45-0.47 M_{\\odot}$ ($\\log g \\approx 7.67$ at $T_{\\rm eff} \\approx 11,500$ K) \\citep{nld96,id99,ap04,dav06,jap07} do not undergo a He core flash in the course of their evolution and therefore are left with a He core.  Two modes of evolution can truncate the red giant branch evolution and prevent the He core flash:  mass loss due to winds and mass loss due to binary interaction.  In systems of high metallicity, mass loss due to stellar winds on the red giant branch can be significant enough to lose the H envelope prior to the core flash \\citep{nld96,bmh05,mk07b}.  Binary interaction through a common envelope also leads to significant mass loss \\citep{ii93,trm95}.  Helium is thus the expected core composition for WDs below $\\approx 0.45-0.47 M_{\\odot}$.  However, little direct evidence exists of the He core.  Possible evidence would be the apparent over-brightness of old WDs \\citep{bmh05} in the star cluster NGC 6791 \\citep{lrb05}.  Though uncertainties remain \\citep{cjd02,lrb08a,lrb08b}, recent detection of low $\\log g$ young WDs \\citep{jsk07} makes it plausible for many of the old WDs to be He core.  A detailed asteroseismological study of these low-mass WDs could provide convincing evidence for the core composition.  Observation and analysis of a full spectrum of the pulsation modes in a WD can produce a wealth of information about the interior structure of the WD.  The mean period spacing of the modes, the rate of change in a mode's period over time, and multiplet splitting of individual modes can provide information on the total mass, spin rate, magnetic field strength, mass of H envelope, and core composition of the WD \\citep{ahc02,bgc08}.  This has already been theoretically applied to distinguish between C/O and O/Ne core WDs by \\citet{ahc04}.  The measured change in an observed mode period in G117-B15 has also been used to constrain significantly the C/O core composition of this object \\citep{sok91,sok95,sok00,sok05a}.  We plot a version of the empirical instability strip in Figure \\ref{fig:instrip} .  Included are not observed to vary (NOV) systems and pulsating ZZ Ceti stars from the observations of \\citet{pb04} and \\citet{ag05,ag07}.  Also included are low-mass WDs from \\citet{mk07a}.  There is a notable absence of low-mass ($\\log g \\lesssim 7.67$) WDs within the instability strip.  There are a few possible ZZ Ceti stars of this low mass that are not plotted due to the absence of spectroscopically determined $\\log g$ and $T_{\\rm eff}$ measurements \\citep{bv07}.  Also shown are the He WD cooling tracks of two models for $\\approx 0.19$, 0.24, 0.40, and $0.45 M_{\\odot}$ WDs \\citep{lga01,jap07}.  The difference between these two models at low mass is due to the different H envelope masses.  \\citet{lga01} used a $1 M_{\\odot}$ main sequence star and truncated its evolution up the red giant branch at various stages to produce He WDs of varying masses.  \\citet{jap07} used close binary evolution expectations for main sequence stars of many masses to produce He WDs of varying masses.  These different approaches cause the different remnant H envelope masses that yield a degeneracy in the He WD mass and its position in the $T_{\\rm eff}-\\log g$ plane.  This degeneracy would be broken in the case of a ZZ Ceti He WD where the pulsation mode spectrum would reveal the H envelope mass.  HS 1824+6000 (hereafter HS 1824, see Table \\ref{tab:prop}) was initially observed by \\citet{bv06} to exhibit pulsations.  Their photometrically determined $\\log g$ and $T_{\\rm eff}$ placed its mass at $\\approx 0.40 M_{\\odot}$ using the tables of \\citet{lga97}.  This mass was well within the theoretical expected mass range for He core WDs making it an excellent object to compare and contrast its pulsation frequencies with other C/O core and possible He core DAVs.  However, later spectroscopic measurement by \\citet{ag07} determined its mass to be $\\approx 0.51 M_{\\odot}$, beyond the expected mass range for He core WDs.  In \\S\\ref{sec:obsana} we discuss our own observations and differential photometry of HS 1824.  In \\S\\ref{sec:timing} we apply a Lomb-Scargle Periodogram approach to a non-uniformly sampled time series in order to obtain the pulsation frequencies of HS 1824.  In \\S\\ref{sec:flcpgs} we report the results of our observations and analysis.  In \\S\\ref{sec:conc} we compare all observed ZZ Ceti periods with $\\log g < 8.0$.  It is our hope this will yield a `zeroth-order' approach to He core identification in much the same way $T_{\\rm eff}$ and $\\log g$ measurements of field WDs identify likely ZZ Ceti stars.  No singular distinction is currently present.   \\clearpage ",
        "conclusions": "\\label{sec:conc}  \\begin{figure} \\centering \\epsscale{1.0} \\plotone{f3.eps} \\caption{Top Panel: Lomb-Scargle Periodogram for combined data set (see \\S\\ref{sec:flcpgs}).  Bottom Panel: Lomb-Scargle Periodogram for de-signaled (using only the four significant to 90\\% detected frequencies) data set .  For both plots the dashed lines denote power level required for 90\\% significance.} \\label{fig:cspec} \\end{figure}  \\begin{figure*} \\centering \\epsscale{1.0} \\plotone{f4.eps} \\caption{Spectrum of reported pulsation periods for published ZZ Ceti systems with spectroscopically determined $\\log g < 8$.  HS 1824 is highlighted in gray with its four observed period locations marked with four vertical lines.  Some marks represent more than one (very closely spaced) observed pulsation period, see references for details.  HL Tau 76 lists only the verified independent pulsation modes of \\citet{nd06}.  Several systems listed here are not included in Figure \\ref{fig:instrip} as their $\\log g$ and $T_{\\rm eff}$ measurements are not of sufficient precision. References: 1  - \\citealt{bgc06}, 2 - \\citealt{asm04a}, 3 - \\citealt{gv00}, 4 - \\citealt{rs05}, 5 - \\citealt{ag07}, 6 - \\citealt{ag06}, 7 - \\citealt{fm05}, 8 - \\citealt{pb04}, 9 - \\citealt{nd06}, 10 - \\citealt{pb95}, 11 - \\citealt{sok05b}, 12 - \\citealt{asm02}, 13 - \\citealt{asm03}.} \\label{fig:perspec} \\end{figure*}  We have successfully detected four pulsation frequencies (periods), 2.751190 mHz (363.479 sec), 3.116709 mHz (320.851 sec), 3.495113 mHz (286.114 sec), and 4.443120 mHz (225.067 sec), in multiple observations of HS 1824+6000.  There are also two possible pulsation frequencies (periods) at 4.450643 mHz (224.687 sec) and 5.755451 mHz (173.748 sec).   With these periods of pulsation in HS 1824, the question remains if it, or other low gravity systems, can be empirically distinguished from the normal C/O core ZZ Ceti population.  To answer this we compiled all known ZZ Ceti stars with published pulsation periods and spectroscopically measured gravities of $\\log g < 8.0$.  This search resulted in 30 systems including HS 1824.  In Figure \\ref{fig:perspec} we plot all reported periods for these 30 systems.  Across all of these ZZ Ceti systems there exist many reported pulsation periods ranging from $100-1400$ sec.  However, it is apparent that better than half of the reported periods reside within the range of $150-400$ sec.  The four periods of HS 1824 are indistinguishable from the rest of this set of ZZ Ceti stars.  Further, there does not appear to be any distinction between the two low-mass ($\\log g \\lesssim 7.67$) systems (HE 0031-5525, SDSS J2135-0743, \\citealt{bgc06}) and the rest of the set.  With this current set of data it appears that this empirical analysis of reported pulsation periods is not sufficient to distinguish a suspected He core from a normal C/O core.  However, HE 0031-5525, and SDSS J2135-0743 \\citep{bgc06} are very close to the boundary of He and C/O core WDs and within the errors of their $\\log g$ measurements may be C/O cores.  It remains  uncertain as to what degree this period spectrum comparative analysis can succeed. There are two primary differences between He and C/O core WDs that affect  g-modes: the contrast in mean molecular weights in their cores, and the one fewer stratified layer in a He core object.  G-modes penetrate deeply into the core, so that  differences in the Brunt-V\\\"{a}is\\\"{a}l\\\"{a} profile there (due to the mean molecular weight; see \\citealt{cjd02}) significantly change the resulting mode period spectrum \\citep{pa06}. The stratified layers of material within the WD also affects how different pulsation modes are trapped, driven, and excited \\citep{ahc02,pa06}. He core WDs possess only two zones of He and H, while C/O core WDs possess the additional zone of C/O. Qualitatively, both of these differences would produce differences in the mode period spectra, and are the subject of current theoretical work we are pursuing. Once  these full mode calculations are available, we can answer whether clear differences are observable.  Most reported systems in Figure \\ref{fig:perspec} were found in observational campaigns looking only for pulsations in an effort to constrain the ZZ Ceti instability strip.  In most cases, no attempt was made to distinguish observed pulsation periods as independent modes, as opposed to linear combinations of modes.  This analysis was neglected in large part due to the lack of extensive follow up.  Our observations of HS 1824 showed most single nights of data contain the pulsations of one specific period and it was a rarity to find a night of data with multiple pulsation periods.  Ideally, very long gapless observations on the order of several days would address these problems very well.  These observations could be obtained through the use of telescope networks such as the Whole Earth Telescope\\footnote{http://www.physics.udel.edu/darc/wet} as was done with HL Tau 76 \\citep{nd06} and G117-B15A \\citep{sok91,sok95} and the Las Cumbres Observatory Global Telescope\\footnote{http://www.lcogt.net}.  We look toward future, more detailed observations of many low-mass and normal-mass ZZ Cetis to help provide a measurable distinction between He and C/O core compositions in WDs."
    },
    "0808/0808.2210_arXiv.txt": {
        "abstract": "{We present the first part of a new catalog of variable stars (OIII-CVS) compiled from the data collected in the course of the third phase of the Optical Gravitational Lensing Experiment (OGLE-III). In this paper we describe the catalog of 3361 classical Cepheids detected in the $\\approx40$ square degrees area in the Large Magellanic Cloud. The sample consists of 1848 fundamental-mode (F), 1228 first-overtone (1O), 14 second-overtone (2O), 61 double-mode F/1O, 203 double-mode 1O/2O, 2 double-mode 1O/3O, and 5 triple-mode classical Cepheids. This sample is supplemented by the list of 23 ultra-low amplitude variable stars which may be Cepheids entering or exiting instability strip.  The catalog data include {\\it VI} high-quality photometry collected since 2001, and for some stars supplemented by the OGLE-II photometry obtained between 1997 and 2000. We provide basic parameters of the stars: coordinates, periods, mean magnitudes, amplitudes and parameters of the Fourier light curve decompositions. Our sample of Cepheids is cross-identified with previously published catalogs of these variables in the LMC. Individual objects of particular interest are discussed, including single-mode second-overtone Cepheids, multiperiodic pulsators with unusual period ratios or Cepheids in eclipsing binary systems.  We discuss the variations of the Fourier coefficients with periods and point out on the sharp feature for periods around 0.35~days of first-overtone Cepheids, which can be explained by the occurrence of 2:1 resonance between the first and fifth overtones. Similar behavior at $P\\approx3$~days for 1O Cepheids and $P\\approx10$~days for F Cepheids are also interpreted as an effect of resonances between two radial modes. We fit the period--luminosity relations to our sample of Cepheids and compare these functions with previous determinations.}{Cepheids -- Stars: oscillations -- Magellanic Clouds} ",
        "introduction": "The Optical Gravitational Lensing Experiment (OGLE) is a wide-field sky survey originally motivated by search for microlensing events (Paczy{\\'n}ski 1986). The observing strategy of the project is to regularly monitor brightness of about 200~million stars in the Magellanic Clouds and Galactic bulge in the time-scales of years. A~by-product of these observations is an enormous database of photometric measurements, which can be used for selecting long lists of newly discovered variable stars.  The OGLE project yielded a wealth of information about variable stars. The second phase of the survey, conducted between 1997 and 2000, resulted in catalogs of thousands Cepheids, RR~Lyr stars, eclipsing binaries and long period variables in the Magellanic Clouds. Moreover, the huge catalogs of variable sources found in the OGLE-II fields in the Magellanic Clouds ({\\.Z}ebru{\\'n} \\etal 2001) and in the Galactic bulge (Wo{\\'z}niak \\etal 2002) were released. In this paper we present the first part of the OGLE-III Catalog of Variable Stars (OIII-CVS) -- the catalog containing virtually all variable stars in the fields regularly observed by OGLE since 2001 with the Warsaw telescope at Las Campanas Observatory, Chile.  Classical Cepheids ($\\delta$~Cep stars, type I Cepheids), as the primary distance indicator, are among the most important variable stars. The Large Magellanic Cloud (LMC) is one of the most fundamental extragalactic targets of modern astrophysics, because it is our nearest non-dwarf neighbor galaxy. For this reason we begin the OIII-CVS with the catalog of classical Cepheids in the LMC.  Large number of variable stars in the LMC, including classical Cepheids, were discovered by Leavitt (1908). However, the first period derivation and the plot of the period--luminosity (PL) diagram for 40 LMC Cepheids was made by Shapley (1931). In 1955 the periods of 550 classical Cepheids in the LMC were published (see Shapley and McKibben Nail 1955 for the bibliography). Then, a considerable survey for LMC Cepheids was done by Woolley \\etal (1962). The catalog prepared by Payne-Gaposchkin (1971) on the basis of Harvard photographic plates contained about 1100 Cepheids in the LMC. After hiatus, a number of Cepheids in the LMC were also discovered by Hodge and Lee (1984), Kurochkin \\etal (1989), van Genderen and Hadiyanto Nitihardjo (1989) and Mateo \\etal (1990). In the late 1990's very large catalogs of Cepheids were published as a by-product of gravitational microlensing surveys: EROS (Beaulieu \\etal 1995, Afonso \\etal 1999), MACHO (Welch \\etal 1997, Alcock \\etal 1999b) and OGLE-II (Udalski \\etal 1999d, Soszy{\\'n}ski \\etal 2000).  The catalog described in this work contains the largest sample of classical Cepheids detected to date in the LMC and, likely, in any other environment. Almost 1000 objects are new identifications. Double-mode Cepheid sample presented in this paper is three times more numerous than the largest sample presented so far. We also show individual objects of particular interest, like triple-mode Cepheids, Cepheids with non-radial pulsations, Cepheids in eclipsing binary systems, etc.  The paper is organized as follows. In Section~2 we describe how the observations were obtained and reduced. Section~3 gives the details about the process of Cepheid selection. In Section~4 we describe the catalog itself. We compare our sample with previously published catalogs of the LMC classical Cepheids in Section~5. In Section~6 we discuss the Fourier coefficients as a function of periods. In Section~7 we fit the period--luminosity relations. Finally, Section~8 summarizes the paper. ",
        "conclusions": "In this paper we present the largest catalog of classical Cepheids in the LMC and probably the largest sample of such stars identified to date in any environment. Our list of Cepheids is supplemented by the high quality, long-term standard photometry enabling precise analysis of these stars. These data are ideal for studying many fundamental problems, such as interpretation of the pulsational and evolutionary models of Cepheids, non-radial oscillations in the pulsating stars, possible non-linearity of the PL relation, structure and history of the LMC.  The catalog contains very rare objects, such as Cepheids with three radial modes excited, 1O/3O double-mode Cepheids, single-mode second-overtone pulsators, Blazhko Cepheids, eclipsing binary systems containing Cepheids including system of two Cepheids eclipsing each other. Our data show that first-overtone classical Cepheids and high amplitude $\\delta$~Sct stars follow continuous PL relation. Distribution of the Fourier parameters suggests that the internal resonance between radial modes may occur twice for the first-overtone pulsators: for periods of about 0.35~days and 3~days. The PL relation for first-overtone Cepheids is possibly non-linear, with a~discontinuity in the slope around $P=0.5$~days.  In the next parts of the OIII-CVS we will present other members of the Cepheid family: type II Cepheids, anomalous Cepheids, HADS and RR~Lyr stars. It is possible that the list of classical Cepheids described in this paper will be supplemented with additional objects of this type detected during further analysis.  \\Acknow{The authors wish to thank Prof.~W.A.~Dziembowski, Prof.~M.~Feast, and Prof.~W.~Gieren for many helpful suggestions which improved the paper. We thank Drs. Z.~Ko{\\l}aczkowski, T.~Mizerski, G.~Pojma{\\'n}ski, A.~Schwar\\-zenberg-Czerny and J.~Skowron for providing the software and data which enabled us to prepare this study.  This work has been supported by the Foundation for Polish Science through the Homing (Powroty) Program and by MNiSW grants: NN203293533 to IS and N20303032/4275 to AU.  The massive period searching was performed at the Interdisciplinary Centre for Mathematical and Computational Modeling of Warsaw University (ICM UW). We are grateful to Dr. M.~Cytowski for helping us in this analysis.}"
    },
    "0808/0808.2483_arXiv.txt": {
        "abstract": "{Blazars are thought to emit highly-collimated outflows, so-called jets. By their close alignment to our line of sight, relativistic beaming effects enable us to observe these jets over the whole electromagnetic spectrum up to TeV energies, making them ideal laboratories for studying jet physics. In the last years multiwavelength observations of blazars provided us with detailed data sets which helped to characterize the two main components of the non-relativistic emission, peaking in the optical to X-ray and GeV/TeV energy region, respectively. In leptonic acceleration models, they are explained by synchrotron radiation of electrons and inverse-Compton emission from the same electron population and thus, correlations of both emission regimes are expected. We review recent observational results on the presence and absence of such correlations in blazars, and discuss constraints on emission models by quantitative correlation analyses.}  \\FullConference{Workshop on Blazar Variability across the Electromagnetic Spectrum\\\\ April 22-25 2008\\\\ Palaiseau, France}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "The up to now best-studied objects concerning correlations of the two bumps in the spectral energy distribution of blazars all belong to the class of high-peaked BL~Lac objects. These bright TeV blazars have been studied in great detail for about the last ten years on various timescales and during various emission-level episodes. Correlated variability up to now has been observed on timescales ranging from hours to months, up to years and seems to be an observationally established fact for the well-studied TeV blazars. Most of the results, however, still suffer from various possible experimental caveats, including the need for an as-simultaneous-as-possible time coverage and a good sensitivity of the instruments included in the observational campaigns. Last but not least, one needs some luck to be able to follow a blazar during an outburst and be able to infer the physics of the flare production and decay.  There have been, however, few campaigns, as e.g. the 2001 campaign on Mkn 421, which provided the possibility to study correlations during the rise and decay phase of individual flares; particularly also the observation of the 2006 flares of PKS 2155--304 seem to offer almost a dissecting knife for studying the quantitative imprints of the physics processes that produce flares on the correlations.  At the same time, interesting phenomena like ``orphan flares'' without X-ray counterparts and ``childless X-ray flares'' need to be explained, and more generally, a systematic understanding of lags between the two photon populations is still elusive.  While correlations of the emission on the high-energy tails of the two photon populations have been studied in great detail, lately also potential correlations of VHE $\\gamma$ rays with the optical or radio emission have been investigated. Optical triggers seem to be a successful proxy to finding new TeV emitters, while first studies of radio-TeV correlations seem to help locating the region in which the SSC phenomenon takes place, profiting from the high spatial resolution that VLBI can offer."
    },
    "0808/0808.0165_arXiv.txt": {
        "abstract": "The CoNFIG (Combined NVSS-FIRST Galaxies) sample is a new sample of 274 bright radio sources at 1.4~GHz. It was defined by selecting all sources with \\SoneG$\\ge$1.3~Jy from the NRAO-VLA Sky Survey (NVSS) in the North field of the Faint Images of the Radio Sky at Twenty-cm (FIRST) survey. New radio observations obtained with the VLA for 31 of the sources are presented. The sample has complete FRI/FRII morphology identification; optical identifications and redshifts are available for 80\\% and 89\\% of the sample respectively, yielding a mean redshift of $\\sim$0.71. One of the goals of this survey is to get better definitions of luminosity distributions and source counts of FRI/FRII sources separately, in order to determine the evolution of the luminosity function for each type of source. We present a preliminary analysis, showing that these data are an important step towards examining various evolutionary schemes for these objects and to confirm or correct the dual population unified scheme for radio AGN. Improving our understanding of radio galaxy evolution will give better insight into the role of AGN feedback in galaxy formation. ",
        "introduction": "\\cite{Long66} determined that powerful radio sources undergo strong differential evolution, the first indication of cosmic downsizing. Since then our understanding of the space density of AGN as a function of cosmic epoch has steadily continued to advance.\\\\ \\indent With the development of evolutionary models for radio sources came the idea of a dual-population model. The initial version of this dichotomy \\citep{Long66} was in terms of low and high luminosities. But with high-frequency surveys, and the large number of flat/inverted spectrum sources revealed in them, an alternative classification emerged, based exclusively on the source spectra: sources with a spectral index $\\alpha \\le -0.5$ (where $S^{\\alpha}_{\\nu} \\propto \\nu^{\\alpha}$),  corresponding to optically thin synchrotron radiation, were classified as steep-spectrum, whereas sources with lower spectral indices ($\\alpha \\ge -0.5$) inevitably showed features of synchrotron self-absorption and were classified as flat-spectrum. Initial indications suggested that these two populations underwent different evolution \\citep{Schmidt76,Masson77}. However, \\cite{Dunlop90} studied the radio luminosity function (density of sources with a given luminosity per unit of co-moving volume) of these two classes of radio sources, and came to the conclusion that both populations were undergoing a similar evolution, implying that they might not actually be distinct.\\\\ \\indent The Fanaroff-Riley (FR) classification \\citep{FR74}, originally a sub-classification for steep-spectrum objects, was employed to provide a further categorization of radio sources. This classification divides radio sources into two classes of double-lobed sources based on the appearance of their jets. The FRI objects have the highest brightness along the jets and core, reside in moderately rich cluster environments \\citep{Hill91} and include sources with irregular structure \\citep{Parma92}. In contrast, FRII sources show hot spots in the lobes and more collimated jets, are found in more isolated environments and display stronger emission lines \\citep{Rawlings89,Baum89}. \\cite{FR74} found these two classes to be divided in radio power, with a break luminosity $P_{\\rm 178~MHz} \\sim 10^{25}\\, {\\rm W   Hz^{-1}sr^{-1}}$, with FRII sources lying above this limit. Subsequently \\cite{Owen94} showed that the break was a function of both radio and optical luminosity.\\\\ \\indent During the 1980s the 'unification' hypothesis emerged to describe how viewing aspect could relate RQSOs, (Radio Quasi-Stellar Objects of either flat or steep-spectrum)  to FRII radio galaxies \\citep[e.g.][]{Peacock87,Scheuer87,Barthel89}. However, the scheme did not include lower-luminosity AGNs such as FRI galaxies and BL Lac objects. The unifying connection between these was introduced by \\cite{Marcha95}. The unified model of AGN proposed by \\cite{Wall97} and \\cite{Jack99} assumes that the cosmic evolution of radio loud AGN is based on a division of the radio sources into a low-luminosity ($P_{\\rm 178~MHz}<10^{25}\\, {\\rm W Hz^{-1}sr^{-1}}$) component corresponding to FRIs, and a high-luminosity component corresponding to FRIIs. In this scheme, the various forms of AGN observed (FRI and FRII extended double sources, flat- and steep-spectrum RQSOs and BL~Lac objects) result from the orientation of the extended parent objects with respect to the observer's line-of-sight. Indeed, because the double-sided ejection of synchrotron blobs in AGN is at relativistic speed, the orientation of the ejection axis to the line-of-sight becomes crucial: sources viewed side-on appear as double radio galaxies (FRI or FRII) and sources viewed along the jets appear as RQSOs (beamed counterparts of FRII sources) or BL~Lac objects (beamed counterparts of FRI sources). The relativistically-boosted jet emission in the beamed counterparts of the extended sources dominates the extended emission, making the overall radio emission appear compact down to VLBI scales.\\\\ \\indent Initially, in modelling the space density of radio AGN, it was assumed for simplicity that the low-luminosity radio galaxies including FRI sources showed no cosmic evolution \\citep{WPL80,Jack99}, the strong cosmic evolution confined only to the higher luminosities and the FRII galaxies. With the advent of large-scale redshift surveys for nearby galaxies, many authors, including \\cite{Brown01}, \\cite{Snellen01}, \\cite{Willott01}, \\cite{Sadler07} and \\cite{Rigby07}, found significant evolution for low power sources -- but mild evolution in comparison with that of the high-luminosity sources. \\cite{Rigby07} argued that if both FRIs and FRIIs have similar evolution, the dual-population scheme could be reduced to a single-population model.\\\\  The FRI/FRII dual-population scheme has encountered several problems. One of these concerns the correspondence between FRI galaxies and BL~Lac objects. \\cite{Urry95} noted that some BL~Lac objects have non-FRI-like morphologies and that the density of FRI sources is too low to account for the entire BL~Lac population, a concern also raised by \\cite{Wall97}. Looking at BL Lac objects from another point of view, \\cite{March96} demonstrate that only about one third of low-luminosity core dominated radio sources - which are supposedly the beamed counterpart of FRI sources - are conventional BL~Lac objects. Most of the remaining sources have optical classification such as Seyfert objects or elliptical galaxies.\\\\ \\indent A related issue concerns the existence of FRI RQSOs. Until recently, these objects were thought not to exist, leading to the hypothesis that FRI and FRII central engines were of different nature \\citep{Baum95} and that the torus opening in FRI sources was too small to observe a quasar nucleus \\citep{Falcke95}. However, the discovery of an FRI QSO, E1821+643, by \\cite{Blun01} overthrew those assumptions. More recent VLA observations \\citep{Heywood07} uncovered another 4 sources of this type.\\\\ \\indent Finally, if sources with different FR classes undergo different evolution, this might imply that their fundamental characteristics, such as the black hole spin or jet composition, are different too. However, the existence of hybrid sources, which display both FRI and FRII morphological characteristics \\citep{Capetti95}, then becomes puzzling. Based on observations of hybrid sources, \\cite{Kaiser97} argued that the FR dichotomy is based purely on the interaction between the jets and the environmental medium, and not on intrinsic properties of the central engine.  This view is also supported by \\cite{Gopal00} and \\cite{Gawr06}. However, \\cite{Wang92} suggested that some AGN engines could be capable of ejecting jets of unequal power, resolving the problem of hybrid sources. \\cite{Gopal00} found no evidence for such a process in their sample.\\\\  \\indent Other schemes have been suggested to resolve these difficulties. \\cite{Kaiser07} explained the FR dichotomy by postulating that all sources start as FRII objects and used an analytical model in which the evolution of the radio sources is governed by energy losses from both radiating relativistic electrons in the lobes and turbulence in the jets.\\\\ \\indent \\cite{Willott01} used an approach to a dual-population unified scheme based on the luminosity of sources instead of morphology. Optical spectra of FRII sources are heterogeneous and they can display both strong  and weak low-excitation emission lines \\citep{Laing94}. Therefore, radio sources can also be grouped based on their emission lines, with one population composed of low-luminosity sources having weak emission lines (containing both FRI and FRII objects), and the other composed of high-luminosity sources with strong emission lines (containing only FRII objects). With this model, \\cite{Willott01} concluded that high luminosity objects undergo a stronger evolution with epoch than low luminosity sources, as found by by e.g. \\cite{Long66}, \\cite{WPL80} and \\cite{Urry95}, and that the radio luminosity function has the form of a broken power-law, similar to the conclusions of \\cite{Dunlop90}. We note that, since conventional accretion-disk systems are expected to be strong X-ray sources, luminosity-based dual-population models are also used in modelling the luminosity function of X-ray selected RQSOs \\citep{Hardcastle06}.\\\\  Defining the relation between the different radio sub-populations together with their cosmic evolution is becoming fundamental to our understanding of galaxy formation.\\\\ The current paradigm for galaxy formation follows hierarchical build-up in a Cold Dark Matter (CDM) universe. Nevertheless, serious difficulties arise from this model in its simplest form, as discussed by \\cite{Bower06}. It implies that current epoch galaxies must be the largest and bluest and have the highest star forming rate of all galaxies. Observations show that they are red, old galaxies, whereas the bulk of star formation is observed at earlier epochs. This is known as downsizing, first described by \\cite{Cowie96}. AGN negative feedback could be the key to understanding this phenomenon. The ignition of the nucleus in a star forming galaxy could eject the gas into the inter-galactic medium, thus reducing or even stopping star formation, breaking the hierarchical buildup \\citep{Silk98,Granato01,Quilis01}. Note that there is also some positive AGN feedback \\citep{vanB04,Klamer04} in which  the pressure from the jets compresses the inter-stellar medium and induces star formation. The balance between these processes remains to be understood; establishing the cosmic behaviour of the radio AGN is important in revealing the precise role of the feedback mechanisms.\\\\  Although FRI and FRII sources show different evolution, they also lie in different luminosity ranges. There is therefore a possibility that both types may show similar evolution for overlapping luminosities (i.e. high-luminosity FRIs and low-luminosity FRIIs). In order to sort out the FR dichotomy and its details, accurate models of the evolution of each population are needed. This implies compiling accurate statistics, such as luminosity distributions and source counts, for both types separately. This is the goal of the CoNFIG (Combined NVSS-FIRST Galaxies) sample presented here, a new sample of bright radio sources with complete morphological identification.\\\\ \\indent The CoNFIG sample is defined as all sources with \\SoneG$\\ge$1.3~Jy from the NRAO-VLA Sky Survey \\citep[NVSS,][]{Condon98} catalogue within the north region of the Faint Images of Radio Sources at Twenty-cm survey \\citep[FIRST,][]{White97}, a 1.5~sr region defined roughly by $-8^{\\circ} \\le {\\rm Dec}\\le 64^{\\circ}$ and 7 hr $\\le {\\rm RA}\\le$ 17 hr. The flux density limit of 1.3~Jy was chosen so that the number of sources in the sample was of statistical significance while allowing us to identify the morphology for each source individually. Optical identifications were obtained from the  SuperCOSMOS Sky Survey \\citep[SSS,][]{Hambly01} and redshift information extracted using the SIMBAD  database. With the accompanying VLA observations described here, the sample has complete morphology information, a median flux density of \\SoneG$\\sim$1.96~Jy and optical identifications and redshift information for $\\sim$80\\% and $\\sim$89\\% of the sources respectively.\\\\  The structure of this paper is as follows. The details of the construction of the CoNFIG sample are explained in \\S\\ref{sample} while \\S\\ref{morpho} describes how the morphologies were determined. Optical identifications and redshift information are discussed in \\S\\ref{IDandz} and \\S\\ref{lumdist} outlines the computation of the morphology-dependent luminosity distributions. Finally, \\S\\ref{extSC} describes and discusses the FRI/FRII source counts.\\\\ \\indent Throughout this paper, we assume a standard $\\Lambda$CDM cosmology with \\Hzero=70 km s$^{-1}$ Mpc$^{-1}$, \\OM=0.3 and \\OL=0.7. ",
        "conclusions": "The CoNFIG sample is constructed as a sample of 274 radio sources from NVSS with $S\\ge 1.3$~Jy. Redshift information is available for $\\sim$80\\% of the sample, and morphological classifications were obtained for all of the sources, either from NVSS and FIRST contour plots, or, for 46 sources, from 8~GHz VLA observations. These data allow us to compute morphology-dependent luminosity distributions and source counts.\\\\ \\indent To increase the number of sources with morphology information, three more samples were constructed in sub-areas of the main region with flux density limits of 0.8~Jy, 0.2~Jy and 50~mJy. Morphological identifications were obtained only from NVSS and FIRST contour plots for those sources. Morphological information for the CENSORS and BDFL sample were obtained from the Ledlow-Owen relation between radio power and optical magnitude and from published classification respectively. Combining these six samples allowed us to compile source counts for FRI and FRII sources separately. A simple, single evolution model for space density was then fitted to these data.\\\\ \\indent Our data show mild evolution of the FRI sources at low redshift; however, they do not participate in the ``evolution bump'' around \\SoneG$\\sim$1~Jy. The results also support the observation that a large number of mJy sources are FRIs galaxies and not starburst galaxies as previously assumed.\\\\"
    },
    "0808/0808.0171_arXiv.txt": {
        "abstract": "We present results of VLBI observations of the water masers associated with IRAS 4A and IRAS 4B in the NGC 1333 star-forming region taken in four epochs over a two month period. Both objects have been classified as extremely young sources and each source is known to be a multiple system. Using the Very Long Baseline Array, we detected 35 masers in Epoch I, 40 masers in Epoch II, 35 in Epoch III, and 24 in Epoch IV.  Only one identified source in each system associates with these masers.  These data are used to calculate proper motions for the masers and trace the jet outflows within 100 AU of IRAS 4A2 and IRAS 4BW. In IRAS 4A2 there are two groups of masers, one near the systemic cloud velocity and one red-shifted. They expand linearly away from each other at velocities of 53 \\kms. In IRAS 4BW, masers are observed in two groups that are blue-shifted and red-shifted relative to the cloud velocity.  They form complex linear structures with a thickness of 3 mas (1 AU at a distance of 320 pc) that expand linearly away from each other at velocities of 78 \\kms. Neither of the jet outflows traced by the maser groups align with the larger scale outflows. We suggest the presence of unresolved companions to both IRAS 4A2 and 4BW. ",
        "introduction": "Water masers are excellent probes of astrophysical flows. Their large flux densities and exceedingly compact sizes make them perfect Very Long Baseline Interferometer (VLBI) targets. Many masers observed in star-forming regions are thought to form in shocks either in or along the outflow of material commonly seen associated with young stellar objects (YSOs).  They provide a high resolution probe of the base of the stellar wind that is unaffected by extinction from dust or extensive interaction with ambient material. The milliarcsecond resolution of VLBI is critical to resolve the individual maser components and outflows, particularly in high stellar density environments.  Long term total power monitoring has shown that masers around low mass YSOs are more episodic than those around higher mass YSOs, however, their detectable phases are sufficiently long for proper motion measurements if the observations are spaced by no more than two to three weeks \\citep{Furuya:2003,Brand:2003, CLAUSSEN:1996,WILKING:1994}.  Furthermore, if they arise in the warm (400K) spatially confined post-shock region, as postulated in various models \\citep{MODEL1, MODEL2, MODEL3, MODEL4}, they should have space velocities sufficient to produce measurable proper motions over one to three weeks. However, regular monitoring of low mass YSOs is necessary to ensure that VLBI observations are conducted during periods of high maser activity.  Using the NRAO's \\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} Very Long Baseline Array (VLBA), proper motion studies have been made of water masers associated with high mass and intermediate mass YSOs such as Cepheus A HW2 \\citep{TORRELLES:2001}, S~106 FIR \\citep{FURUYA:2000},  NGC 2071 \\citep{SETH:2002}, and IRAS~20050+2720 MMS1 \\citep{FURUYA:2005} as well as with lower mass YSOs such as Serpens SMM1, RNO 15-FIR, and IRAS~05413-0104 \\citep{Moscadelli:2006,CLAUSSEN:1998}. For example, water masers in IRAS~05413-0104 (associated with HH~212 at a distance of 450 pc), were found to lie in a structure of 10 AU in length and less than 0.5 AU thick. The masers appeared to arise in shock-related structures which showed proper motions along the axis of the outflow of 60 \\kms and displayed coherent structures over time scales of two to three weeks (but which was less discernible over time scales as long as several months).  Comparison with molecular and infrared observations of IRAS~05413-0104 clearly demonstrated that proper motion studies of water masers in the region (within 40 AU) of the central source are among the best tools for studying the kinematics of the jets emanating from embedded YSOs.  IRAS~4, comprised of multiple sources whose submillimeter-dominated 'Class 0' spectral energy distributions suggest extreme youth, is located in the NGC 1333 star-forming region.  As a nearby active and multiple young object,  it presents a prime candidate for VLBI observations of its associated water masers. Located at a distance of 320 pc \\citep{HIPPARCOS}, the NGC 1333 molecular cloud hosts a double infrared cluster of about 200 YSOs identified in near-infrared surveys \\citep{STROM:1976,ASPIN:1994,Lada:1996,Wilking:2004} and x-ray surveys \\citep{PREIBISCH:1997,Getman:2002,PREIBISCH:2003}. Far-infrared observations have not only revealed the higher luminosity sources in NGC 1333, but also YSOs in the earliest phase of evolution such as the Class 0 sources IRAS~4 and IRAS~2 \\citep{Harvey:1984,JENNINGS:1987}. IRAS~4 was found to be a binary at submillimeter wavelengths \\citep{SANDELL:1991}; hence IRAS~4A and 4B.  IRAS~4A is located 31$\\arcsec$ to the NW of 4B at a position angle of -45$^{\\circ}$. Subsequently, both IRAS~4A and 4B were found to be binary systems. \\citet{LAY:1995} performed single baseline interferometry in the submillimeter and found that 4A was a binary with an angular separation of 1.8$\\arcsec$ (or 580 AU at 320 pc).  Interferometric observations in the millimeter and radio continuum confirmed the separation and determined that the position angle of separation was $\\approx -20^{\\circ}$ \\citep{LOONEY:2000,REIPURTH:2002}. While the easternmost source (IRAS~4A1) dominates the millimeter and radio continuum emission, the westernmost source IRAS~4A2) appears to be relatively more evolved with a warm ammonia core \\citep{WOOTTEN:1993,Shah:2000} and a wealth of complex organic molecules (Bottinellii 2007, priv. comm.).  Submillimeter and millimeter wave continuum observations have confirmed the prediction of \\citet{LAY:1995} that IRAS~4B is also a binary with components (4BE and 4BW) separated by about 10$\\arcsec$ (3200 AU in projection) along an E-W direction \\citep{SMITH:2000,LOONEY:2000,Sandell:2001}. \\footnote{We adopt the naming convention introduced by \\citet{Sandell:2001}. The western (eastern) source is referred to as IRAS~4BI (IRAS~4BII) by \\citet{SMITH:2000}, as IRAS~4B (IRAS~4C) by \\citet{LOONEY:2000}, and as IRAS~4B (4B$\\arcmin$) by \\citep{DIFRANCESCO:2001}.}  As in the case of IRAS~4A2, complex organic molecules have been detected toward IRAS~4B \\citep{Bottinelli:2007}, presumably from a hot core associated with 4BW. We note that IRAS~4BE has yet to be detected in the radio continuum \\citep{REIPURTH:2002} or demonstrate any sign of outflow activity.  Infall motions have been detected toward IRAS~4A and 4B \\citep{DIFRANCESCO:2001}, unambiguously toward the former, and strongly indicated toward the latter.  The detection of water masers toward IRAS~4A and 4B suggests that they are YSOs also associated with mass outflow.  Water masers were first detected in single dish observations by \\citet{HASCHICK:1980} and later monitored by \\citet{CLAUSSEN:1996}.  These observations established the maser emission as highly variable on monthly time scales, sometimes being completely absent while at other times reaching peaks of 10 Jy.  Detected emission was found from $-$10 to $+$15 \\kms but the angular resolution was not sufficient to determine which masers were associated with 4A and 4B.  Beginning in 1983, VLA observations have shown water maser activity associated only with IRAS~4A2 and/or 4BW \\citep{RODRIGUEZ:2002,Wootten:1998,Furuya:2003}.  VLBA observations of the water masers in IRAS~4A and 4B, acquired in 2003, have been presented by \\citet{Desmurs:2006}.  While masers were detected in only two of the four epochs observed, they confirmed the association of masers with IRAS~4A2 and expansion motion between the two spots was detected. In IRAS~4BW, the five maser spots detected were red-shifted and formed a chain 80 mas in extent with some masers displaying proper motion toward the north.  Molecular outflows have been mapped toward both IRAS~4A and 4B. IRAS~4A is associated with a highly collimated outflow about 20,000 AU in extent seen in HCN and SiO with a position angle close to 20$^{\\circ}$ \\citep{Girart:1999,Choi:2001,Choi:2005}.  The origin of the outflow is likely IRAS~4A2 \\citep{Choi:2005}.  On a larger scale ($\\approx 1 \\times 10^{5}$ AU), the outflow defined by CO and molecular hydrogen has a position angle of 45$^{\\circ}$ \\citep{BLAKE:1995,Choi:2006}.  The shift in position angle from small to large scale is perhaps due to a combination of a shift in the direction of the magnetic field, an encounter with denser ambient gas, and precession of the outflow axis \\citep{Choi:2006}.  A more compact outflow is associated with IRAS~4BW, with a position angle close to 0$^{\\circ}$ \\citep{BLAKE:1995,Choi:2001} with perhaps a second outflow with a position angle of -35$^{\\circ}$ \\citep{DIFRANCESCO:2001}.  In this paper we present VLBA observations of water masers associated with the IRAS~4 region obtained over four epochs in 1998 and spaced by about one month. We also describe Very Large Array (VLA)  observations, taken one month in advance of the start of the VLBA observations, that showed IRAS~4 to be a very active maser source and yielded absolute positions for the masers and the 1.3 cm continuum sources IRAS~4A1 and IRAS~4BW.  With the VLBA, we detected 35 masers in Epoch I, 40 masers in Epoch II, 35 in Epoch III, and 24 in Epoch IV associated with IRAS~4A2 and IRAS~4BW.  We use these data to reveal the structure of shock fronts associated with stellar winds from these YSOs and to estimate proper motions for masers detected in all four epochs. The origin of the maser emission and its relationship to the stellar winds and larger scale molecular outflows from these YSOs is discussed. ",
        "conclusions": "We have observed the water masers associated with IRAS~4A and 4B at VLBI resolutions in four epochs over three months.  We have determined that the masers are related to the jets emanating from these YSOs due to their spatio-kinematic distribution.  In both sources the masers are found associated with known single components of multiple systems, with total separation velocities between 53 and 78 km s$^{-1}$.  This is further confirmed by the large proper motions measured for both sources, which clearly rules out rotation due to the mass constraints placed on the central objects by other observations.  The water masers of IRAS~4B form arc-like structures roughly 10 AU in length and less than 0.6 AU in thickness.  The structure of these structures changes rapidly with time, with many new maser components appearing and disappearing in just one month.  The orientation of the structures in the plane of the sky does not agree with larger scale outflow angle, which we attribute to a possible unseen very close companion. Future observations will have to sample the source more frequently than once every three weeks, with once every 3-5 days probably providing the best results."
    },
    "0808/0808.2497_arXiv.txt": {
        "abstract": "Many current and future astronomical surveys will rely on samples of strong gravitational lens systems to draw conclusions about galaxy mass distributions.  We use a new strong lensing pipeline (presented in Paper~I of this series) to explore selection biases that may cause the population of strong lensing systems to differ from the general galaxy population.  Our focus is on point-source lensing by early-type galaxies with two mass components (stellar and dark matter) that have a variety of density profiles and shapes motivated by observational and theoretical studies of galaxy properties.  We seek not only to quantify but also to understand the physics behind selection biases related to: galaxy mass, orientation and shape; dark matter profile parameters such as inner slope and concentration; and adiabatic contraction. We study how all of these properties affect the lensing Einstein radius, total cross-section, quad/double ratio, and image separation distribution, with a flexible treatment of magnification bias to mimic different survey strategies.  We present our results for two families of density profiles: cusped and deprojected S\\'ersic models.  While we use fixed lens and source redshifts for most of the analysis, we show that the results are applicable to other redshift combinations, and we also explore the physics of how our results change for very different redshifts.  We find significant (factors of several) selection biases with mass; orientation, for a given galaxy shape at fixed mass; cusped dark matter profile inner slope and concentration; concentration of the stellar and dark matter deprojected S\\'ersic models.  Interestingly, the intrinsic shape of a galaxy does not strongly influence its lensing cross-section when we average over viewing angles. Our results are an important first step towards understanding how strong lens systems relate to the general galaxy population. ",
        "introduction": "\\label{S:motivation}  The field of strong lensing is undergoing a period of growth that is expected to accelerate in the coming decade.  Current samples with tens of lenses, such as from CASTLES\\footnote{CASTLES is a collection of uniform HST observations of mostly point-source lenses from several samples with differing selection criteria, rather than a single, uniformly-selected survey.} \\citep[e.g.,][]{2001ASPC..237...25F}, CLASS \\citep[e.g.,][]{2003MNRAS.341...13B}, SDSS \\citep[e.g.,][]{2008AJ....135..496I}, and SLACS \\citep[e.g.,][]{2008ApJ...682..964B}, will give way to samples with hundreds or even thousands of strong lensing systems discovered by Pan-STARRS, LSST, SNAP, and SKA \\citep[e.g.,][]{ 2004AAS...20510827F, 2004NewAR..48.1085K, 2004ApJ...601..104K, 2005NewAR..49..387M}. As lens samples grow, they will be even more useful for constraining the mass distributions of the strong lensing galaxy population \\citep[for a review of strong lensing astrophysics, see][]{Saas-Fee}.  There are, however, certain complications in using strong lensing studies to constrain the physical properties of galaxies.  The issue we address here is \\emph{selection bias}.  Suppose we treat galaxies as having an intrinsic probability distribution in some parameter $x$ (such as the inner slope of the density profile), and we want to use observed strong lens systems to infer that distribution $p(x)$.  If the lensing probability itself depends on $x$, then in general the distribution $p_\\mathrm{SL}(x)$ for the strong lens galaxies will not reflect the true, underlying $p(x)$ for all galaxies.  The more diverse the galaxy population, and the more the strong lensing cross-section varies across the population, the more important the selection biases can be.  Obviously, one must account for selection biases in order to draw reliable conclusions from current and future strong lens samples.  Our purpose here is to use a complete, end-to-end, and above all, realistic simulation pipeline for strong lensing (presented in Paper~I of this series, van de Ven, Mandelbaum, \\& Keeton 2008) to quantify and understand strong lensing selection biases and assess their impact in typical situations.  We focus on strong lensing of quasars by early-type, central (non-satellite) galaxies at several mass scales, using two-component mass profiles (dark matter plus stars) that are consistent with existing photometry and stacked weak lensing data from SDSS \\citep{2003MNRAS.341...33K,2006MNRAS.368..715M} as well as with $N$-body and hydrodynamic simulations.  The issue of selection biases in strong lensing is not a new one, and we are building on many previous studies of this issue.  The main galaxy properties that we study in an attempt to understand selection bias are \\begin{itemize} \\item Mass \\citep[e.g., ][]{1984ApJ...284....1T, 1991MNRAS.253...99F, 2007MNRAS.379.1195M}; \\item Orientation \\citep[``inclination bias'', e.g., ][]{1997ApJ...486..681M, 1998ApJ...495..157K, 2007arXiv0710.1683R}; \\item Shape \\citep[e.g., ][]{1997ApJ...482..604K, 2005ApJ...624...34H, 2007arXiv0710.1683R}; \\item Dark matter inner slope \\citep[e.g., ][]{2001ApJ...549L..25K, 2001ApJ...555..504W,2002ApJ...566..652L}; \\item Dark matter concentration \\citep[e.g., ][]{2004ApJ...601..104K, 2007A&A...473..715F}. \\end{itemize} In most cases, previous studies addressed the issue of selection bias using models that were simplified in some way. Simplifications often included testing effects of dark matter slope and concentration using pure (generalised) NFW profiles without a baryonic component; or testing effects of density profile or shape using single-component mass models.  A notably different approach was taken by \\citet{2007MNRAS.382..121H}, \\citet{2007MNRAS.379.1195M}, and \\citet{2008MNRAS.386.1845H}, who used the Millennium simulation with a semi-analytic model of galaxy formation to simulate  strong lensing.  That approach naturally yields a realistic distribution of halo and galaxy properties (to the extent that the true distribution of galaxy properties and their relation to halo properties is encoded in the semi-analytic model), at least for a fixed cosmological model, and it will undoubtedly be useful for modeling the strong lensing population in large, future surveys.  We elect, however, to take a more controlled approach here, using realistic but discrete values of the density profile parameters, so that we can disentangle the different types of lensing selection biases, and understand the physics of each one, before recombining them to determine the net effects.  Our basic approach is to generate realistic galaxy models with various values of the relevant parameters, and to investigate how the lensing cross-section $\\sigma(\\vec{x})$ depends on model parameters $\\vec{x}$ (where $\\vec{x}$ might include mass, shape, concentration, and other parameters, some of which may be correlated). The reason for focusing on the cross-section is that it represents the weighting factor that transforms the intrinsic joint parameter distribution $p(\\vec{x})$ into the distribution $p_\\mathrm{SL}(\\vec{x})$ among observed strong lens systems. Note that we focus on selection biases related to \\emph{physical} effects.  There may also be \\emph{observational} selection biases \\citep[e.g., ][]{1991ApJ...379..517K}, but they are specific to a given survey and are not something that we can address in a general way.   This fact is one reason for our focus on point-source lensing: extended source lensing is even more sensitive to observational selection effects than quasar lensing. For example, the effect of a finite aperture is no longer straightforward for an extended source, and detection of an extended strong lens system depends on factors such as the size of the point-spread function which varies spatially and temporally in any given survey.  Fortunately, with large numbers of point-source lens systems anticipated in future surveys [e.g., $\\sim 10^{5}$ with SKA \\citep{2004NewAR..48.1085K}] this investigation of physical selection biases in point-source lens systems should be highly useful. We emphasize that our purpose is only to determine the mapping function $\\sigma(\\vec{x})$ due to physical differences between galaxies, and not to predict the observed distribution of properties $p_\\mathrm{SL}(\\vec{x})$ in some particular lensing survey, for which knowledge of the intrinsic parameter distribution $p(\\vec{x})$ and of any observational selection effects are both necessary.  In this paper, we present a systematic investigation of physical strong lensing selection biases that both unifies and extends previous work.  We begin in Section~\\ref{S:simulations} with a brief review of the simulation pipeline developed in Paper~I, including the density profiles and shapes of our galaxy models, our computational lensing methods, and the basic lensing properties of our galaxies.  In Section~\\ref{S:orientationshape}, we investigate selection biases related to galaxy orientation and shape.  In Section~\\ref{S:innerslope}, we consider selection biases related to the inner slope of the dark matter component of the density profile (for cusped models).  In Section~\\ref{S:concentration}, we allow the dark matter concentration to vary as well.  In Section~\\ref{S:sersic}, we turn to deprojected S\\'ersic density profiles and examine selection biases with both the stellar and dark matter parameters. Section~\\ref{S:robustness} includes an exploration of the ranges of lens and source redshifts for which our conclusions are applicable.  Finally, in Section~\\ref{S:conclusions}, we summarise our main conclusions and discuss their implications for past and future strong lensing analyses. ",
        "conclusions": "\\label{S:robustness}  Finally, to ascertain how much our conclusions about selection biases in the previous sections depend on our chosen fiducial lens and source redshifts, we now try varying the redshifts in several ways.  If we vary both the lens and the source redshifts in a way that preserves $\\Sigma_c$, then the physical scale within the galaxy that is measured by strong lensing is the same, but it will lead to different image separation distributions, scaled by the angular diameter distance $D_L$ of the lens galaxy. This change may lead to {\\em observational} selection effects that depend sensitively on the survey image search strategy (resolution and maximum angular separation), but no physical selection effects, so we do not attempt to model this in any general way.  Now consider fixing the lens redshift (and therefore the conversion from physical to angular scale), but shifting the source redshift to rescale the critical surface mass density $\\Sigma_c$. This shift will change the Einstein radius $\\Rein$ and therefore the physical scale of the density profile that is probed by strong lensing.  Thus, this change may result in different \\emph{physical} selection biases.  If we fix the lens redshift to our fiducial $z_L=0.3$ but move the source redshift considerably lower, from the fiducial $z_S=2.0$ to $z_S'=0.6$, then the critical surface density increases by $\\sim 70$ per cent. This change means that the strong lensing is due to smaller scales, where the average projected surface density is higher, and that the lensing cross-section for this density profile will be lower. Because of the importance of smaller scales, the projected dark matter fraction $\\fpdm$ will decrease. As a test, we recompute \\Rein\\ and the lensing cross-sections for both mass scales as a function of the dark matter halo inner slope $\\gdm$ % to see the degree to which the selection bias we identified with our fiducial lens and source redshifts persist. We also increase the resolution of the surface density maps used by \\gravlens\\ for these calculations by a factor of two, while maintaining the box size, to ensure that the critical curves can be properly resolved.  We also consider a value of $\\Sigma_c$ that is 40 per cent lower than our fiducial value.  While this low $\\Sigma_c$ does not correspond to any sensible value of $z_S$ for our fiducial $z_L$, we assume a higher $z_L\\sim 0.6$ and $z_S\\sim 3$.  Because of the different $z_L$, all angular scales output by our pipeline must be rescaled by $D_A(z=0.3)/D_A(z=0.6)=0.66$, i.e., the Einstein radii must be rescaled by $0.66$ and the cross-sections by $0.66^2$.  However, the relevant quantities for selection bias associated with true variation of galaxy density profiles, e.g. $\\fpdm$, are preserved.  Since we expect increased Einstein radii and cross-sections, we preserve the resolution but increase the box size by a factor of $2$.  Table~\\ref{T:varyz} shows what happens to relevant quantities for lensing when we vary $\\Sigma_c$ to values 70 per cent higher and 40 per cent lower.  We consider only the unbiased cross-sections for the spherical shape model. \\begin{table*} \\caption{\\label{T:varyz}Summary of the lensing results for the cusped, spherical model when varying $\\Sigma_c$ so that different physical scales are probed.  The factor of $0.66$ in \\Rein\\ for the lower $\\Sigma_c$ cases is shown separately since it is not indicative of the physical scale within the galaxy.} \\begin{tabular}{l|r|r|r} \\hline \\hline Quantity & Fiducial $\\Sigma_c$ & Higher $\\Sigma_c$ & Lower $\\Sigma_c$ \\\\ \\hline \\multicolumn{4}{c}{Lower mass scale} \\\\ $\\Rein(\\gdm=0.5)$          & $0.58$ & $0.36$ & $0.79\\times 0.66=0.52$ \\\\ $\\Rein(\\gdm=1.0)$ (arcsec) & $0.68$ & $0.41$ & $0.95\\times 0.66=0.63$ \\\\ $\\Rein(\\gdm=1.5)$          & $1.04$ & $0.61$ & $1.46\\times 0.66=0.97$ \\\\ $\\fpdm(R'<\\Rein,\\gdm=0.5)$ & $0.13$ & $0.08$ & $0.17$ \\\\ $\\fpdm(R'<\\Rein,\\gdm=1.0)$ & $0.28$ & $0.20$ & $0.37$ \\\\ $\\fpdm(R'<\\Rein,\\gdm=1.5)$ & $0.60$ & $0.51$ & $0.67$ \\\\ $\\st(\\gdm=0.5)/\\st(\\gdm=1.0)$ & $0.79$ & $0.79$ & $0.79$ \\\\ $\\st(\\gdm=1.5)/\\st(\\gdm=1.0)$ & $2.65$ & $2.96$ & $2.51$ \\\\ \\hline \\multicolumn{4}{c}{Higher mass scale} \\\\ $\\Rein(\\gdm=0.5)$          & $1.18$ & $0.63$ & $1.73\\times 0.66=1.14$ \\\\ $\\Rein(\\gdm=1.0)$ (arcsec) & $1.75$ & $0.88$ & $2.75\\times 0.66=1.82$ \\\\ $\\Rein(\\gdm=1.5)$          & $4.44$ & $2.33$ & $6.75\\times 0.66=4.46$ \\\\ $\\fpdm(R'<\\Rein,\\gdm=0.5)$ & $0.22$ & $0.13$ & $0.30$ \\\\ $\\fpdm(R'<\\Rein,\\gdm=1.0)$ & $0.51$ & $0.36$ & $0.63$ \\\\ $\\fpdm(R'<\\Rein,\\gdm=1.5)$ & $0.87$ & $0.80$ & $0.91$ \\\\ $\\st(\\gdm=0.5)/\\st(\\gdm=1.0)$ & $0.52$ & $0.51$ & $0.52$ \\\\ $\\st(\\gdm=1.5)/\\st(\\gdm=1.0)$ & $9.19$ & $12.9$ & $8.04$ \\\\ \\hline \\hline \\end{tabular} \\end{table*} First, we can see that when we increase (decrease) $\\Sigma_c$, we decrease (increase) \\Rein\\ and \\fpdm, as expected.  The change in $\\Rein$ can be significant, up to a factor of $\\sim 2$ for the changes in $\\Sigma_c$ that we have tested.  We find that the selection bias with $\\gdm$, represented in the last two rows of this table, is nearly independent of $\\Sigma_c$ for $\\gdm$ below the fiducial value of $1$, and for higher $\\gdm$ it is a slightly increasing function of $\\Sigma_c$.  The sign of this change is surprising, in the sense that if \\fpdm\\ is lower than we might naively expect the parameters of the dark matter distribution to be less significant. However, we can explain this change by the fact that the cross-section \\st\\ is determined by the outer (source-plane) caustic, which relates to the inner (lens-plane) critical curve.  This in turn is determined by the slope of the density profile on small scales, so if the lensing properties are determined by smaller (larger) physical scales, as when we increase (decrease) $\\Sigma_c$, then \\gdm\\ becomes more (less) important and selection bias becomes stronger (weaker).  However, it is encouraging that the selection bias only changes by $+14$ and $-4$ per cent ($+41$ and $-13$ per cent) for changes in $\\Sigma_c$ of $+70$ and $-40$ per cent for the lower (higher) mass model.  This result suggests that within the range of $\\Sigma_c$ that we have considered in this subsection, the conclusions of our paper regarding physical selection bias are not strongly dependent on our choice of fiducial lens and source redshifts.  (The same is not true for observational selection biases, which may be strongly affected by changing $z_L$ and thus the characteristic angular scale of the system.)  On Fig.~\\ref{F:sigmacrit}, we have indicated the lens and source redshift combinations for which our conclusions are robust according to the tests in this subsection.  \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth,angle=0]{scinvcontour_shrink.ps} \\caption{\\label{F:sigmacrit}Critical surface density, plotted as $\\log{[\\Sigma_c/(M_{\\odot}/\\mathpc^2)]}$, as a function of lens and source redshift.  For $z_S<z_L$, where $\\Sigma_c$ is formally infinite, we have set it to the maximum value shown, $10^4 M_{\\odot}/\\mathpc^2$. Asterisks indicate the minimum source redshift for a given lens redshift for which we have confirmed that our conclusions about selection bias are valid.} \\end{center} \\end{figure}  We must also consider the issue of magnification bias. Magnification bias complicates a change in the redshift of the system in two ways.  First, if the source redshift is different, then for a given survey flux limit, the luminosity threshold is different.  For a broken power-law luminosity function that is steeper at the bright end, this means that if the source is more distant, the limiting luminosity becomes brighter, so the actual number of lensing systems is modulated due to both the different number density of sources above that limit and by the change in effective $\\st$ due to the steeper luminosity function at the limit.  Second, since the source luminosity function may be a function of redshift, we anticipate that the biased cross-sections, which are the relevant quantity for any particular survey given its flux limit and search strategy, may also change due to this effect. For our analysis in this section, we have ignored this change in the magnification bias formalism, but to determine the cross-sections for a real survey this evolution would have to be taken into account.   In any case, both of these magnification bias effects qualify as observational rather than physical selection biases that are more properly simulated in the context of a particular strong lensing survey."
    },
    "0808/0808.1341_arXiv.txt": {
        "abstract": "We present diffraction limited (0.6\\arcsec) 24.5\\,\\micron~ Subaru/COMICS images of the red supergiant $\\mu$ Cep. We report the detection of a circumstellar nebula, that was not detected at shorter wavelengths. It extends to a radius of at least 6\\arcsec~in the thermal infrared. On these angular scales, the nebula is roughly spherical, in contrast, it displays a pronounced asymmetric morphology closer in. We simultaneously model the azimuthally averaged intensity profile of the nebula and the observed spectral energy distribution (SED) using spherical dust radiative transfer models. The models indicate a constant mass-loss process over the past 1000 years, for mass-loss rates a few times $10^{-7}$\\,\\Msunyr.  This work supports the idea that at least part of the asymmetries in shells of evolved massive stars and supernovae may be due to the mass-loss process in the red supergiant phase. ",
        "introduction": "\\label{intro} Although the final stages of the post-main sequence evolution of massive stars do not last long, it is here where most of the mass is lost and the shaping of the pre-supernova ejecta takes place.  Key in this respect are the Red Supergiants (RSGs) which represent a phase in the life of stars with initial masses 10--30\\,M$_{\\odot}$. During this phase lasting 10$^4 -10^5$ years, the stars lose prodigious amounts of mass at a rate of order 10$^{-6} - 10^{-4}$ M$_{\\odot}$ yr$^{-1}$ (van Loon et al. 2005\\nocite{2005A&A...438..273V}; Massey et al. 2008\\nocite{2008arXiv0801.1806M}), and the final mass of the objects is mostly set during this phase. RSGs have been observed to be the direct progenitors of Type IIP supernovae (Smartt et al. 2004\\nocite{2004Sci...303..499S}), but can also evolve towards the blue via a Yellow Hypergiant phase (Oudmaijer et al. 2008\\nocite{2008arXiv0801.2315O}; de Wit et al. 2008a\\nocite{2008A&A...480..149D}) to become a Wolf-Rayet star and eventually explode as a supernova (Meynet \\& Maeder 2000\\nocite{2000A&A...361..101M}).  It is not only the study of the stars themselves that helps us understand the final stages of the evolution of massive stars. Investigating RSGs' mass-loss and their circumstellar material helps us to understand the origin of the aspherical structures found around supernovae, such as SN 1987A's rings, or gamma-ray bursts.  It is possible that at least some of these observed asymmetries originate during the RSG phase (e.g. Chita et al. 2008). Indeed, in several instances, the mass-loss during the RSG phase has been found to deviate from spherically symmetric, and can have a complex appearance. Much effort has been directed towards the objects VY CMa and NML Cyg, extremely bright, cool, objects close the empirical Humphreys Davidson limit (Humphreys \\& Davidson 1979\\nocite{1979ApJ...232..409H}). High resolution studies have revealed highly anisotropic structures in the wind of VY CMa (e.g. Smith et al. 2001\\nocite{2001AJ....121.1111S}), while NML Cyg's aspherical appearance is shaped by the strong radiation from the nearby Cyg OB2 association (Morris and Jura 1983; Schuster et al 2006\\nocite{2006AJ....131..603S}).  A third RSG that is found in the same location in the HR diagram as the above two objects is $\\mu$ Cep (Schuster et al 2006\\nocite{2006AJ....131..603S}). Contrary to these two cooler objects, little is known about the material surrounding $\\mu$ Cep. This can be readily explained by its mass-loss, which is orders of magnitude lower than for NML Cyg or VY CMa (cf. Jura \\& Kleinman 1990\\nocite{1990ApJS...73..769J}). With an $M_{\\rm bol}$ of $-$9.08, this M2Ia star is one of the brightest RSGs known, lying on evolutionary tracks corresponding to stars with initial masses greater than 25 M$_{\\odot}$ (Levesque et al. 2005\\nocite{2005ApJ...628..973L}).  Using Levesque et al.'s parameters, we find a distance to the object of 870 pc, and a luminosity of $\\sim$ 3$\\times 10^5$ L$_{\\odot}$. Here, we report on the discovery of an extended dust shell around the object by means of diffraction limited imaging (0.6\\arcsec) of mid-IR thermal dust emission. ",
        "conclusions": "We have obtained the first diffraction-limited images of a Red Supergiant at 24.5\\,$\\mu$m, and resolved the circumstellar material around $\\mu$ Cep out to 6\\arcsec. The intensity profile and the SED were simultaneously fitted with a dust model. The main results can be summarized as follows:   1. The outer parts of the shell are to first order circular with apparent inhomogeneities, the inner parts, tracing the mass lost in the past 1000 years display a flattened structure, which is possibly a slowly expanding, dense torus. The immediate conclusion from this is that any asymmetries in the shells of massive evolved stars and Supernovae ejecta may find their origins in the Red Supergiant phase. This is qualitatively very similar to what is found for the lower-mass AGB stars which evolve into bipolar Planetary Nebulae.  2. The mass-loss rate for $\\mu$ Cep is found to be a few times 10$^{-7}$\\,M$_{\\odot}\\,$yr$^{-1}$. This low mass-loss rate may explain the fact that $\\mu$ Cep is not detected in scattered light, as opposed to other RSGs which have mass-loss rates that are orders of magnitude larger.  3. A general conclusion is that imaging thermal dust emission at 24.5\\,$\\mu$m is a viable manner to spatially resolve the circumstellar material around evolved objects with low mass-loss rates and as a consequence are otherwise not detected with maser observations, CO imaging or scattered light imaging."
    },
    "0808/0808.0617_arXiv.txt": {
        "abstract": "We present an analysis of the gas physics at the base of jets from five T Tauri stars based on high angular resolution optical spectra, using the Hubble Space Telescope Imaging Spectrograph (HST/STIS). The spectra refer to a region within 100 AU of the star, i.e. where the collimation of the jet has just taken place. We form PV images of the line ratios to get a global picture of the flow excitation. We then apply a specialised diagnostic technique to find the electron density, ionisation fraction, electron temperature and total density. Our results are in the form of PV maps of the obtained quantities, in which the gas behaviour is resolved as a function of both radial velocity and distance from the jet axis. They highlight a number of interesting physical features of the jet collimation region, including regions of extremely high density, asymmetries with respect to the axis, and possible shock signatures. Finally, we estimate the jet mass and angular momentum outflow rates, both of which are fundamental parameters in constraining models of accretion/ejection structures, particularily if the parameters can be determined close to the jet footpoint. Comparing mass flow rates for cases where the latter is available in the literature (i.e. DG Tau, RW Aur and CW Tau) reveals a mass ejection-to-accretion ratio of 0.01 - 0.07. Finally, where possible (i.e. DG Tau and CW Tau), both mass and angular momentum outflow rates have been resolved into higher and lower velocity jet material. For the clearer case of DG Tau, this revealed that the more collimated higher velocity component plays a dominant role in mass and angular momentum transport. ",
        "introduction": "\\label{introduction}  In current star formation theory, jets/outflows from a newly forming star are believed to transport significant amounts of energy and momentum away from the region of the central source (\\citealp{Bally07}; \\citealp{Ray07}; \\citealp{Pudritz07}; \\citealp{Shang07}). This can have a big influence on the way in which the stars form, because, for example, jets may drive the injection of turbulence in the parent cloud, thereby regulating the star formation rate. At the same time, they may be able to extract the excess angular momentum  from the accretion disk, thus allowing the matter to drift through it and finally accrete onto the star. Jets, therefore, are considered a fundamental ingredient in the star formation process. To fully understand the  mechanisms underlying the physics of these commonly observed nebulae, however, it is necessary to know the mass outflow rate, which regulates the dynamics of the flow and, therefore, is the most important input parameter for any model of flow generation and propagation. For example, in a series of recent papers by our group, differential Doppler shifts at the borders in a number of T\\,Tauri jets were reported, suggestive of rotation around the symmetry axis (\\citealp{Bacciotti02}; \\citealp{Woitas05}; \\citealp{Coffey04}; 2007). The combination of mass outflow rate and toroidal velocity estimates allows determination of the angular momentum transported by the flow, and thus also a comparison with the angular momentum that the associated disk has to loose in order to accrete at the observed rate. Indeed in one test case, namely the jet from RW\\,Aur, we found that the angular momentum present in the jet is at least 60 - 70\\,\\% of that required to be extracted from the disk (\\citealp{Woitas05}). This is of course a result of primary importance, since it could be the first long-awaited validation of the popular idea that jets exist to remove the excess angular momentum in forming system, thus providing a solution to one of the main theoretical problems in star formation.  In order to confirm this finding, however, we have to determine the mass and angular momentum outflow rate close to the star, and in a wider sample of objects. A good starting point is to use the {\\it HST}/STIS spectra that were recently acquired by our group in a survey to search for rotation signatures in jets from T\\,Tauri stars. The survey included jets from five young stars, namely TH\\,28, DG\\,Tau, HH\\,30, CW\\,Tau and RW\\,Aur. The spectra, taken with the slit placed at a few tens of arcseconds from the star, and oriented transversely to the flow direction, sampled the collimation region of the jets and allowed us to confirm the presence of differential Doppler shifts at the jet borders and derive estimates of toroidal velocities (\\citealp{Coffey04}; 2007). These spectra, taken with the slit perpendicular to the jet axis, included several optical forbidden emission lines that can be used to diagnose the physical conditions of the gas. Thus, we have embarked in a further analysis of these data to understand how the plasma properties behave close to the star where the jet is launched, collimated and accelerated, and finally to find the mass and angular momentum outflow rate in our sample. We point out that estimating the mass outflow rate at the jet base is more likely to exclude spurious effects, such as the addition of mass from possible entrainment of ambient gas and dust or significant disruption of the flow, which are less likely to have a dominant influence in the early stages of jet propagation.  The determination of the mass outflow rate requires a knowledge of the total gas density, which, contrary to the electron density, is not directly measurable from the observed lines. One, therefore, has to use some method to derive the total density from the available observational data, and indeed several methods have been developed to this aim in the recent past. For example, from a measure of the intrinsic emission line luminosity one can derive the number of emitting particles of a given species in the observed volume, which can then lead to the gas density under an assumption of abundances. However, this method relies on an accurate prior knowledge of reddening estimates, excitation temperature, ionisation state of the given species, and filling factor, all of which bring into the calculation substantial uncertainties (see e.g. \\citealp{Nisini05}). An alternative method is via a determination of the hydrogen ionisation fraction of the emitting region. A direct measure of the electron density can be easily obtained from the [S\\,II]$\\lambda\\lambda$6716,6731 doublet (see e.g. \\citealp{Osterbrock89}). Dividing the electron density by the ionisation fraction then closely approximates the total gas density. One approach to finding the ionisation fraction is to model the line ratios under the assumption of a definite mechanism for the gas heating. There is not yet, however, a general consensus on the mechanism causing the jet emission although, without a doubt, the observed forbidden lines are excited collisionally. The most widely accepted explanation is that the gas is being heated by internal shocks, although other possibilities include ambipolar diffusion, turbulent mixing-layers and compression by jet instabilities (see e.g. \\citealp{Lavalley-Fouquet00} and references therein). On the basis that the gas is shock heated, \\citet{Hartigan94} constructed a grid of planar shock models to calculate the line ratios, and compared the results to spectra of a few stellar jets integrated along the beam. In this way, they were able to find the most likely value of the hydrogen ionisation fraction in the flows, and hence the total density, after accounting for the shock compression. This method, however, is model dependant, and implies heavy and lenghty calculations.  Subsequently, a technique was developed whereby the ratios of the optical forbidden emission lines of S, N and O in the red wavelength range are used to infer the hydrogen ionisation and the electron temperature {\\it without} having to assume a specific method of heating, but only assuming that the atomic levels are populated collisionally in the absence of ionising photons (\\citealp{Bacciotti95}; \\citealp{BE99}).  This so-called 'BE technique' is based on the recognition that, in the conditions present in stellar jets, the ionisation states of oxygen and nitrogen are tightly correlated to that of hydrogen via charge-exchange. Consequently, the line ratios can be easily modeled and compared with the observed ratios to find the most suitable values of the gas physical quantities. In addition to the requirement of very little calculations, the technique does not require an accurate knowledge of distance or extinction. Furthermore, the results do not depend on any specific mechanism for jet formation and/or evolution, and thus can be applied to targets in very different conditions.  The {\\it BE technique} has been applied to jet observations of varying spectral and spatial resolution (\\citealp{Bacciotti95}; \\citealp{Bacciotti96}; \\citealp{Bacciotti99}; \\citealp{Nisini05}; \\citealp{Podio06}), which lately included high angular resolution observations of the DG\\,Tau, RW\\,Aur, LkH$\\alpha$233 jets using {\\em HST}/STIS multiple-slit configuration, (\\citealp{Bacciotti02mexico}; \\citealp{Melnikov08}) as well as sub-arcsecond ground-based data of the DG\\,Tau and RW\\,Aur jets obtained with adaptive optics (\\citealp{Lavalley-Fouquet00}; \\citealp{Dougados02}). Very recently, \\cite{Hartigan07} used a slightly different version of the {\\it BE technique} to diagnose the physical condition of the HH\\,30 jet from {\\em HST}/STIS `slitless' spectra. In all cases, the method has brought very valuable information on the excitation of the jet gas allowing, for instance, constraints on models for jet heating, and allowing confirmation that shocks are, as expected, the most likely cause of the thermal behaviour of the plasma. In most cases, it has also been possible to estimate the mass and momentum flux in the flow from the derived total densities. This in turn has facilitated a discourse on the dynamical relationship between optical jets and coaxial molecular flows (see, e.g. \\citealp{BE99}). Furthermore, as mentioned above, in the test case of the RW Aur jet, using the results of the {\\it BE technique} and Doppler gradient measurements, we could give an estimate of the angular momentum transported by the flow.  For all these reasons, we decided to use the {\\it BE technique} to determine the mass and angular momentum flux in our sample of jets observed at high angular resolution. The STIS data included all the optical lines necessary to apply this procedure, which has proven to be very well suited to the diagnostics of large datasets in extended jet surveys, as is the case of this paper. In the following, we describe the application of the method to our spectra (Section\\,\\ref{observationsanddataanalysis}), and the obtained results (Section\\,\\ref{results}). Given the form of our input, we present the results for the physical conditions as position-velocity (PV) maps, in which the electron density, the ionisation fraction, the electron temperature and the total density are resolved both as a function of velocity and distance from the jet axis. From these values, we derived estimates of the mass and angular momentum outflow rates of the jet close to the launch point. Our findings are discussed in Section\\,\\ref{discussion}. ",
        "conclusions": "\\label{conclusions}  We have analysed the gas physics for several T\\,Tauri jets close to the launching point, based on high resolution {\\it HST}/STIS spectra taken with the slit perpendicular to the flow direction. We have applied the {\\it BE diagnostic technique} to the line emission spectra, to obtain maps of the electron density, temperature and ionisation fraction. To this aim, we have also adapted the {\\it BE code}, to allow inclusion of more datapoints in the calculations. The new approach involves easing the algebraic constraints, which cause difficulties where jets exhibit a broad range of velocities and the various the various lines trace different velocity components, i.e. DG\\,Tau. Our results represent the first survey of physical conditions at the base of T Tauri jets presented in the form of position-velocity diagrams for the physical quantities. In fact, with our high spatial and spectral resolution dataset, we resolved the jet physics as a function of velocity and distance from the jet axis, in the region just a few tens of AU above the disk plane where the flow is launched.  The overall survey results indicate that, at the jet base, the plasma has a high electron density ($>$2~10$^{4}$\\,cm$^{-3}$), high electron temperature (2\\,10$^{4}$\\,K), and low ionisation level (0.03 - 0.3) which varies considerably depending on the target. Indeed, in all cases saturation of the [\\ion{S}{2}] doublet is reached in at least some of the datapoints, and so in this region close to the launch point of the jet the electron density reported is often a lower limit. This is different to findings further along the jet which are a factor of ten lower (\\citealp{Bacciotti99}; \\citealp{Podio06}). In the case where previous studies have been carried out, we find good agreement for values reported close to the star. In the case of TH\\,28 (which presents the best dataset), possible shock signatures are present, thus providing a observational indications that shocks can contribute to heating the jet close to the source.  We determine the mass and angular momentum outflow rates for the jets close to their base. Estimates we determined for the mass and angular momentum outflow rates, both of which are fundamental parameters in constraining models of accretion/ejection structures, particularily if the parameters can be determined close to the jet footpoint. Values for a single jet lobe are in the range 4.0~10$^{-9}$ to 6.7~10$^{-8}$\\,M$_\\sun$\\,yr$^{-1}$ and 1.1~10$^{-6}$ to 1.3~10$^{-5}$\\,M$_\\sun$\\,yr$^{-1}$\\,AU\\,km\\,s$^{-1}$. Again we find good agreement with the literature in cases where values were reported close to the star. Mass flow ratios were found to be $\\dot M_{jet}/\\dot M_{acc}$\\,$\\sim$\\,0.01 - 0.07 where accretion rates were available in the literature (i.e. for 3 of 5 targets). This is in the range predicted by accretion-ejection models (\\citealp{Konigl00}; \\citealp{Casse00}; \\citealp{Shu00}).  Although we have examined a region of the jet at about 50 - 80\\,AU from the source corresponding to the collimation zone, we note that the region where the jet is formed and launched is believed to be on scales of less than 1\\,AU, a region which is currently out of the reach of present instrumentation and often obscured by infalling matter. We await near-infrared interferometry as an opportunity to observing this zone.  \\vspace {0.2in} {\\bf Acknowledgements} \\vspace {0.1in} \\newline The present work was supported in part by the European Community's Marie Curie Actions - Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network, under contract MRTN-CT-2004-005592."
    },
    "0808/0808.2612_arXiv.txt": {
        "abstract": "An analytical method is presented for treating the problem of a uniformly rotating, self-gravitating ring without a central body in Newtonian gravity. The method is based on an expansion about the thin ring limit, where the cross-section of the ring tends to a circle. The iterative scheme developed here is applied to homogeneous rings up to the 20th order and to polytropes with the index $n=1$ up to the third order. For other polytropic indices no analytic solutions are obtainable, but one can apply the method numerically. However, it is possible to derive a simple formula relating mass to the integrated pressure to leading order without specifying the equation of state. Our results are compared with those generated by highly accurate numerical methods to test their accuracy. ",
        "introduction": "The problem of the self-gravitating ring captured the interest of such renowned scientists as \\citet{Kowalewsky85}, \\citet{Poincare85b,Poincare85c,Poincare85d} and \\citet{Dyson92, Dyson93}. Each of them tackled the problem of an axially symmetric, homogeneous ring in equilibrium by expanding it about the thin ring limit. In particular, Dyson provided a solution to fourth order in the parameter $\\sigma=a/b$, where $a$ provides a measure for the radius of the cross-section of the ring and $b$ the distance of the cross-section's centre of mass from the axis of rotation. An important step toward understanding rings with other equations of state was taken by \\citet{Ostriker64,Ostriker64b,Ostriker65}, who studied polytropic rings to first order in $\\sigma$ and found a complete solution to this order for an isothermal limit. \\par  First numerical results for homogeneous rings were given by \\citet{Wong74}, who was not able to clarify the transition to spheroidal bodies that \\citet{Bardeen71} had supposed would exist. \\citet{ES81} and \\citet{EH85} developed improved methods with which they were able to study the connection to the Maclaurin spheroids. Returning to the problem significantly later, \\citet*{AKM03} achieved near-machine accuracy, which allowed them to study bifurcation sequences in detail and correct erroneous results. It was also possible to extend the problem to non-homogeneous rings and even to the framework of General Relativity \\citep*{Hachisu86,AKM03c,FHA05}. \\par  Through the use of computer algebra, we extend Dyson's basic idea and determine the solution to the problem of the homogeneous ring up to the order $\\sigma^{20}$. We also present an iterative method for performing a similar expansion about the thin ring limit for arbitrary equations of state and general results are derived, confirming and generalizing work that had already been published by \\citet{Ostriker64b}. The application to polytropes is considered and ordinary differential equations (ODEs) are found that allow for the determination of the mass density. A closed-form solution can only be found if the value of the polytropic index is $n=1$, and such rings are considered to the order $\\sigma^3$. For other polytropic indices, the ODEs are solved numerically so that results from the approximate scheme can be compared to highly accurate numerical results for a variety of equations of state. The numerical solutions considered here are taken from a multi-domain spectral program, much like the one described in \\citet*{AKM03b}, but tailored to Newtonian bodies with toroidal topologies (see \\citealt{AP05} for more information). The solutions obtained by these numerical methods are extremely accurate and thus provide us with a means of testing the accuracy of the approximate method. ",
        "conclusions": "In their work on homogeneous rings, Poincar\\'e and Kowalewsky, whose results disagreed to first order, both had made mistakes as Dyson has shown. His result to fourth order is also erroneous as we point out in the appendix. It thus seems particularly worthwhile to test the correctness of the solutions presented here. For one thing, we ensured that the transition condition \\begin{align} \\nabla U_\\text{in}|_\\text{s}=\\nabla U_\\text{out}|_\\text{s} \\end{align} is fulfilled up to the appropriate order in $\\sigma$. Furthermore we tested that the virial theorem \\eqref{VI} is fulfilled for each order in $\\sigma$. \\par Please note that one has to be careful in interpreting the results for the thin ring limit. For example, one might think that the squared angular velocity vanishes like $\\sigma^2\\ln\\sigma$. This is true for the dimensionless quantity $\\Omega^2/G\\mu_\\text c$, but need not be true for the squared angular velocity itself. If we fix the `size' $b$ and the mass $M$ of the ring in that limit, then the cross-section shrinks to a point ($a=\\sigma b$). With \\eqref{mass_h00} we can conclude that $\\mu_\\text c\\propto\\sigma^{-2}$ and therefore $\\Omega^2\\propto\\ln\\sigma$, which means that $\\Omega^2$ and hence the velocity of a fluid element tend to infinity. \\par Relativistic rings, including the thin ring limit, were studied in \\citet{AKM03c}, \\citet{Ansorgetal04} and \\citet{FHA05}. From the perspective of General Relativity, the Newtonian theory constitutes a good approximation when certain conditions are fulfilled. For one thing, typical velocities must be small compared to the speed of light $c$ and for another $|U|\\ll c^2$ must hold. We just saw, however, that for rings of finite extent and mass, the velocities grow unboundedly in the thin ring limit. The same holds for $U_\\text s$ as well, see \\eqref{Us_general}. This means that the Newtonian theory of gravity is not appropriate to describe this subtle limit itself, since one cannot expect it to be a good approximation to General Relativity. It is remarkable that the approximation about the point $\\sigma=0$ is nevertheless so successful."
    },
    "0808/0808.2338_arXiv.txt": {
        "abstract": "Axisymmetric magnetorotational instability (MRI) in viscous accretion disks is investigated by linear analysis and two-dimensional nonlinear simulations. The linear growth of the viscous MRI is characterized by the Reynolds number defined as $R_{\\rm MRI} \\equiv v_A^2/\\nu\\Omega $, where $v_A$ is the Alfv{\\'e}n velocity, $\\nu$ is the kinematic viscosity, and $\\Omega$ is the angular velocity of the disk. Although the linear growth rate is suppressed considerably as the Reynolds number decreases, the nonlinear behavior is found to be almost independent of $R_{\\rm MRI}$. At the nonlinear evolutionary stage, a two-channel flow continues growing and the Maxwell stress increases until the end of calculations even though the Reynolds number is much smaller than unity. A large portion of the injected energy to the system is converted to the magnetic energy. The gain rate of the thermal energy, on the other hand, is found to be much larger than the viscous heating rate. Nonlinear behavior of the MRI in the viscous regime and its difference from that in the highly resistive regime can be explained schematically by using the characteristics of the linear dispersion relation. Applying our results to the case with both the viscosity and resistivity, it is anticipated that the critical value of the Lundquist number $S_{\\rm MRI} \\equiv v_A^2/\\eta\\Omega$ for active turbulence depends on the magnetic Prandtl number $S_{{\\rm MRI},c} \\propto Pm^{1/2}$ in the regime of $Pm \\gg 1$ and remains constant when $Pm \\ll 1$, where $Pm \\equiv S_{\\rm MRI}/R_{\\rm MRI} = \\nu/\\eta$ and $\\eta$ is the magnetic diffusivity. ",
        "introduction": "Magnetohydrodynamic (MHD) turbulence is the most promising candidate for angular momentum transport in astrophysical disk systems. Nonlinear behaviors of the magnetorotational instability (MRI) are actively investigated as a driving mechanism of MHD turbulence over the last decades (Balbus \\& Hawley 1991, 1998). The central issue in MRI research is its nonlinear properties, in particular, the saturation amplitude of the instability. Although the key processes governing the nonlinear saturation are scoped by global and local numerical studies, it is not fully explained yet (Hawley \\& Balbus 1992; Hawley et al. 1995, 1996; Brandenburg et al. 1995; Matsumoto \\& Tajima 1995; Stone et al. 1996; Hawley 2000; Machida et al. 2000; Arlt \\& R\\\"udiger 2001; Balbus 2003). Recently, Lesur \\& Ogilvie (2008) argue that, in the shearing box context with zero-net vertical flux, MHD turbulence is sustained through nonlinear classical dynamo activity once the MRI is operated. It would be necessary and significative to study the saturation process from the microscopic viewpoint of the physical sustaining mechanism for MHD turbulence as they have done.  Ohmic dissipation is one of the crucial processes that determine the saturation amplitude of the MRI. Linear growth rate of the MRI can be reduced significantly because of the suppression by ohmic dissipation. Two- and three-dimensional local shearing box simulations (Sano et al. 1998, 2004; Sano \\& Stone 2002) have shown that physical properties of the saturated turbulence depend on the Lundquist number $S_{\\rm MRI } \\equiv v_A^2/\\eta\\Omega $, where $v_A$ is the Alfv\\'en velocity, $\\Omega $ is the angular velocity, and $\\eta $ is the magnetic diffusivity (see also Fleming et al. 2000; Ziegler \\& R\\\"udiger 2001; Liu et al. 2006). Particularly, when $S_{\\rm MRI} \\ll 1$, the size of the saturated stress decreases with the decrease of $S_{\\rm MRI} $ (Sano \\& Stone 2002; Pessah et al. 2007). It is also pointed out that magnetic reconnection plays an important role in the energy dissipation of MRI driven turbulence (Sano \\& Inutsuka 2001).  Numerical issues are one of the main reasons why the saturation physics of the MRI is remained to be understood. Fromang \\& Papaloizou (2007) demonstrate the efficiency of angular momentum transport decreases linearly with the grid spacing as the resolution increases. Although it is very difficult to distinguish between the numerical and physical factors working as the saturation mechanism (King et al. 2007; Silvers 2007), current researches of the MRI pay much attention to the numerical factors with the greatest care. Pessah et al. (2007) derive a scaling law of the saturated stress from past wide variety of numerical results by analytically eliminating the numerical factors such as box size and resolutions.  More straightforward way for decontaminating the numerical factors is to bring explicitly the physical diffusivities much larger than the numerical one into the computational study. Lesur \\& Longaretti (2007) have performed first systematic study of the MRI in the presence of both viscous and magnetic dissipations, which are larger than the numerical diffusivities. For the cases with nonzero net flux of the vertical field, the transport property in the saturated state depends on the magnetic Prandtl number $Pm \\equiv \\nu/\\eta $, where $\\nu $ is the kinematic viscosity. Fromang et al. (2007) have reported that in zero magnetic flux cases the turbulent activity is an increasing function of the magnetic Prandtl number $Pm$. Linear behaviors of the MRI in the presence of both the viscosity and resistivity are analytically studied by Pessah \\& Chan (2008) comprehensively.  The magnetic Prandtl number $Pm$ takes a wide range of values in astrophysical disk systems. In protoplanetary disks surrounding young stellar objects, the magnetic Prandtl number is much smaller than unity because of their low ionization degree (Nakano 1984; Umebayashi \\& Nakano 1988; Sano et al. 2000). In accretion disks of compact X-ray sources and active galactic nuclei, the magnetic Prandtl number ranges from $\\simeq 10^{-3}$ to $10^{3}$ depending on the distance from the central object (Balbus \\& Henri 2008). In collapsar disks which is known as the central engine of gamma-ray bursts (Woosley 1993), the physical state with high magnetic Prandtl number of $Pm \\gtrsim 10^{10}$ is expected to be realized in their evolutionary stage as a result of the large neutrino viscosity (Masada et al. 2007). Therefore, more systematic and deeper study on the MRI in the presence of both the viscosity and resistivity is quite important for understanding the accretion process triggered by the MRI in various disk systems.  One important unsettled matter, in these situations, is the role of the kinematic viscosity at the nonlinear stage of the MRI. In general, the viscosity  as well as the magnetic resistivity can suppress the growth of the MRI. However the dependence of nonlinear outcome on the Prandtl number indicates that the role of the viscosity in MRI turbulence could be different from that of the resistivity. Then, the main purpose of this paper is to reveal nonlinear features of the MRI in viscous accretion disks. As is described in what follows, linear growth of the MRI is characterized by the Reynolds number $R_{\\rm MRI} \\equiv v_A^2/\\nu\\Omega $ in the viscous fluid, and by the Lundquist number $S_{\\rm MRI} \\equiv v_A^2/\\eta\\Omega $ in the resistive fluid. Focusing on these two non-dimensional parameters, we clarify the difference in nonlinear behaviors of the MRI between the viscous and resistive systems.  Our paper is organized as follows. In \\S~2, linear features of the MRI in the viscous fluid are presented.  In \\S~3, nonlinear behavior of the MRI is investigated by two-dimensional MHD simulations taking account of the viscous terms. The differences between the effect of the viscosity and resistivity are also clarified in \\S~3. Finally we make an physical explanation for our nonlinear results with the help of the linear dispersion relation. Applying our results to double diffusive systems, we predict a condition for sustaining active MRI turbulence in the presence of both the viscosity and resistivity in \\S~4. ",
        "conclusions": "\\subsection{Nonlinear Behavior in the Single Diffusive Systems} To give a physical explanation for the nonlinear behavior of the axisymmetric MRI, we focus on the critical wavelength obtained from the linear theory in this section. The diagrams of Figures~\\ref{fig8}a and \\ref{fig8}b indicate the critical and the fastest growing wavelengths of the MRI as a function of the Lundquist number $S_{\\rm MRI}$ and the Reynolds number $R_{\\rm MRI}$, respectively (see also Table~\\ref{table1}). Note that the vertical axes are normalized by $2\\pi (\\eta/\\Omega)^{1/2} $ in Figure~\\ref{fig8}a and by $2\\pi (\\nu/\\Omega)^{1/2}$ in Figure~\\ref{fig8}b. Shaded area denotes the linearly unstable regions for the MRI.  Assuming fixed diffusivities, $S_{\\rm MRI}$ and $R_{\\rm MRI}$ increase as the instability grows because they are proportional to the squared Alfv\\'en velocity. Then the horizontal axis in Figure~\\ref{fig8} can be regarded as the time direction in terms of the evolution of the MRI.  First, we consider the resistive case shown by Figure~\\ref{fig8}a. For the case of $S_{\\rm MRI} \\lesssim 1$, the critical wavelength is described as $\\lambda_{\\rm crit} \\simeq \\eta/v_A$. At the nonlinear stage of the two-dimensional MRI, MHD turbulence decays and it saturates only when $S_{\\rm MRI} \\lesssim 1$ (see Figs.~\\ref{fig4}b and \\ref{fig7}). This behavior can be interpreted schematically using the $\\lambda_{\\rm crit}$-$S_{\\rm MRI}$ diagram (Sano \\& Miyama 1999).  As the MRI grows and amplifies the magnetic field, the critical wavelength shifts to the shorter length-scale. Then, many smaller scale fluctuations can become unstable. Those structures enhance the efficiency of ohmic dissipation in the turbulent state. In other words, the system evolves toward a more dissipative state, and could be saturated at a critical point around $S_{\\rm MRI} \\simeq 1$, at which the critical wavelength reaches the shortest value and the ohmic dissipation is the most efficient. In this way, MHD turbulence can decay if $S_{\\rm MRI} \\lesssim 1$.  When $S_{\\rm MRI} \\gtrsim 1$, on the other hand, the critical wavelength is given by $\\lambda_{\\rm crit} \\simeq v_A/\\Omega $. Two-dimensional calculations of the MRI suggests that the unstable growth cannot saturate if $S_{\\rm MRI} \\gtrsim 1$. The nonlinear behavior in these cases can be explained by Figure~\\ref{fig8}a analogously as follows: At the linear evolutionary stage, the critical wavelength in the radial direction becomes longer due to the exponential growth of the radial field component, while the vertical field grows slower than exponentially. The wavevectors of the unstable modes thus become parallel to the vertical axis. This channel flow mode is the exact solution of the nonlinear MHD equations (Goodman \\& Xu 1994). As the field strength becomes larger, the critical wavelength shifts to the longer one. The dissipation can be much less effective, and thus the channel solution continues to grow without saturation.  Next, let us consider the viscous case shown in Figure~\\ref{fig8}b. In the viscous fluid, the critical wavelength is given by $\\lambda_{\\rm crit} \\simeq v_A/\\Omega $  despite the size of the Reynolds number (see Figure~\\ref{fig1}). Even if $R_{\\rm MRI} $ is much smaller than unity, the critical wavelength thus shifts to larger scale as the instability grows. Then the system always evolves toward a less dissipative state and is not saturated. This interpretation is consistent with our numerical results shown in Figure~\\ref{fig7}.  These results indicates that the saturation process of the MRI would be changed dramatically at the critical point at which the critical wavelength switches from the decreasing function of the field strength to the increasing one. The differences in the nonlinear behavior of the MRI between the viscous and resistive systems would be originated from whether the critical point exists or not. \\subsection{MRI in the Doubly Diffusive System} In this subsection, we apply the discussion above to the doubly diffusive system that includes both the viscosity and resistivity. The saturation behavior of the axisymmetric MRI can be anticipated by the dependence of the critical wavelength on the field strength derived from the linear theory. With the same framework used in \\S~2, a local axisymmetric dispersion equation of the MRI in the presence of both viscous and ohmic dissipations is given by \\begin{equation} a_4 \\tilde{\\gamma}^4 + a_3 \\tilde{\\gamma}^3 + a_2 \\tilde{\\gamma}^2 + a_1 \\tilde{\\gamma} + a_0 = 0 \\;, \\label{eq20} \\end{equation} where \\begin{eqnarray} a_4 & = & 1\\;, \\ \\ \\ \\ a_3 = 2 \\left( \\frac{1}{R_{\\rm MRI}} + \\frac{1}{S_{\\rm MRI}} \\right) \\tilde{k}_z^2 \\;, \\nonumber \\\\ a_2 & = & \\left( \\frac{1}{R_{\\rm MRI}^2 } + \\frac{1}{S_{\\rm MRI}^2 } + \\frac{4}{R_{\\rm MRI}S_{\\rm MRI} } \\right) \\tilde{k}_z^4+ 2\\tilde{k}_z^2 + \\tilde{\\kappa}^2 \\;,  \\nonumber \\\\ a_1 & = & \\left( \\frac{1}{R_{\\rm MRI} } + \\frac{1}{S_{\\rm MRI} } \\right)\\frac{2}{R_{\\rm MRI}S_{\\rm MRI} } \\tilde{k}_z^6 \\nonumber \\\\ && + 2\\left( \\frac{1}{R_{\\rm MRI}} + \\frac{1}{S_{\\rm MRI}}\\right) \\tilde{k}_z^4 + \\frac{2}{S_{\\rm MRI}} \\tilde{k}_z^2\\tilde{\\kappa}^2 \\;, \\nonumber \\\\ a_0 & = & \\frac{1}{R_{\\rm MRI}^2 S_{\\rm MRI}^2} \\tilde{k}_z^8 + \\frac{2}{R_{\\rm MRI}S_{\\rm MRI} }\\tilde{k}_z^6 \\nonumber \\\\ && + \\left( \\frac{1}{S_{\\rm MRI}^2 } \\tilde{\\kappa }^2 + 1 \\right) \\tilde{k}_z^4 + ( \\tilde{\\kappa}^2 - 4) \\tilde{k}_z^2 \\;, \\nonumber \\end{eqnarray} (Menou et al. 2006; Masada et al. 2007; Lesur \\& Longaretti 2007; Pessah \\& Chan 2008). This equation is, as expected, characterized by $R_{\\rm MRI }$ and $S_{\\rm MRI}$.  Here we focus on the system with a constant magnetic Prandtl number ($Pm \\equiv S_{\\rm MRI}/R_{\\rm MRI} = \\nu/\\eta $). Figure~\\ref{fig9} demonstrates the critical wavelength of the MRI as a function of $S_{\\rm MRI}$ for various values of $Pm$ obtained by solving the dispersion equation (18). The $S_{\\rm MRI}$-dependence of the critical wavelength varies with the size of $Pm$. When $Pm \\ll 1$, the linear growth of the MRI is independent of the magnetic Prandtl number, and the critical wavelength is almost identical to the pure resistive case ($Pm =0$). However, if the viscosity effects is added sufficiently, then the critical wavelength is enlarged by the suppression due to the viscosity in the middle range of $S_{\\rm MRI}$. For the cases of $Pm \\gg 1$, the critical wavelength around $S_{\\rm MRI} \\simeq 1$ is given by $\\tilde{k}^{-1}_{\\rm crit} \\propto (S_{\\rm MRI}R_{\\rm MRI})^{-1/3} \\propto (S_{\\rm MRI}^2 / Pm)^{-1/3}$ (Pessah \\& Chan 2008).  Although the critical wavelength shifts to longer length-scales as the magnetic Prandtl number increases, the critical wavelength has a minimum value for all the cases. This implies that the MRI turbulence could be suppressed if the Lundquist number is less than a critical value. This diagram suggests that the critical Lundquist number $S_{{\\rm MRI},c}$ depends on $Pm$. The critical Lundquist number $S_{{\\rm MRI},c}$ is plotted as a function of $Pm$ in Figure~\\ref{fig10}a. In the regime of $Pm \\gg 1$, it is proportional to the square root of the magnetic Prandtl number, that is $S_{{\\rm crit},c} \\propto Pm^{1/2}$. In contrast, it remains to be constant in the range $Pm \\ll 1$.  Since the critical Reynolds number is given by $R_{{\\rm MRI},c} = S_{{\\rm MRI},c}/Pm$, we can also obtain the relation between $Pm$ and $R_{{\\rm MRI},c}$, and which is depicted in Figure~\\ref{fig10}b. Nonlinear growth of the MRI can be expected in the parameter region above this critical curve. The magnetic Prandtl number is proportional to $R_{{\\rm MRI},c}^{-2}$ in the regime of $Pm \\gg 1$ and $Pm \\propto R_{{\\rm MRI},c}^{-1}$ when $Pm \\ll 1$. This curve is reminiscent of the critical curve for MHD turbulence sketched from nonlinear simulations of MRI (Fromang et al. 2007). The Reynolds number $R_{\\rm MRI}$ in their models are at most a few tens\\footnote{The definition of the Reynolds number $Re$ in Fromang et al. (2007) is different from $R_{\\rm MRI}$ in this paper. The relation between these two is $R_{\\rm MRI} \\approx \\alpha_{M} Re$, where $\\alpha_M$ is the Maxwell stress normalized by the (initial) pressure.}, and the critical magnetic Prandtl number is around unity. Thus our prediction is roughly consistent with nonlinear results even quantitatively. Note that the critical curve shown by Figure~\\ref{fig10} is obtained by using only the features of the linear dispersion relation of MRI. This implies that the linear growth of the MRI is very important even at the nonlinear saturated phase to sustain MHD turbulence.  The discussion in this paper is based only on two-dimensional simulations of the MRI. For the understanding the saturation mechanism of the MRI, it is quite important to perform the systematic three-dimensional analysis of the MRI in the presence of multiple diffusivities. Furthermore, the assumption of the local shearing box could affect the nonlinear evolution of the viscous MRI. The necessary ingredients of the unstable growth of the MRI are the velocity shear and the magnetic field. In the numerical setting of the local shearing box, the velocity profile of the background shear flow cannot disappear by the role of the kinematic viscosity, but is imposed by the boundary conditions. It would be very interesting to use global disk models to investigate the MRI in highly viscous disks. These are our next tasks."
    },
    "0808/0808.3422_arXiv.txt": {
        "abstract": "We study dynamics of bars in models of disk galaxies embeded in realistic dark matter halos.  We find that disk thickness plays an important, if not dominant, role in the evolution and structure of the bars.  We also make extensive numerical tests of different $N$-body codes used to study bar dynamics. Models with thick disks typically used in this type of modeling (height-to-length ratio $h_z/R_d=0.2$) produce slowly rotating, and very long, bars. In contrast, more realistic thin disks with the same parameters as in our Galaxy ($h_z/R_d\\approx 0.1$) produce bars with normal length $R_{\\rm bar} \\approx R_d$, which rotate quickly with the ratio of the corotation radius to the bar radius ${\\cal R} =1.2-1.4$ compatible with observations. Bars in these models do not show a tendency to slow down, and may lose as little as 2-3 percent of their angular momentum due to dynamical friction with the dark matter over cosmological time. We attribute the differences between the models to a combined effect of high phase-space density and smaller Jeans mass in the thin disk models, which result in the formation of a dense central bulge. Special attention is paid to numerical effects such as the accuracy of orbital integration, force and mass resolution.  Using three $N$-body codes -- Gadget, ART, and Pkdgrav -- we find that numerical effects are very important and, if not carefully treated, may produce incorrect and misleading results. Once the simulations are performed with sufficiently small time-steps and with adequate force and mass resolution, all the codes produce nearly the same results: we do not find any systematic deviations between the results obtained with TREE codes (Gadget and Pkdgrav) and with the Adaptive-Mesh-Refinement (ART) code. ",
        "introduction": "Barred galaxies represent a large fraction ($\\sim$65\\% ) of all spiral galaxies \\citep[e.g.,][]{Eskridge00,Sheth08}. Bars are ubiquitous. They are found in all types of spirals: in large lenticular galaxies \\citep{Aguerri05}, in normal spirals such as our Galaxy \\citep{Freudenreich98} and M31 \\citep{Lia06,Beaton07}, and in dwarf magellanic-type galaxies \\citep{Valenzuela07}.  An isolated stellar disk embeded into a dark matter halo spontaneously forms a stellar bar as a result of the development of  global disk instabilities \\citep[e.g.,][Sec. 6.5]{BT}. Bars continue to be closely scrutinized because of their connection with the dark matter halo \\citep[\\eg][]{od03, hbk05, cvk06, lia07}.  Because the bars rotate inside massive dark matter halos, they lose some fraction of the angular momentum to their halos and tend to slow down with time \\citep{Tremaine84,Weinberg85}.  The formation of bars and associated pseudobulges is often considered as an {\\it alternative} to the hierarchical clustering model \\citep{Kormendy04}. This appears to be incorrect: recent cosmological simulations indicate that the secular bulge formation is a {\\it part (not an alternative) of the hierarchical scenario}. The simulations of the formation of galaxies in the framework of the standard hierarchical cosmological model indicate that bars form routinely in the course of assembly of halos and galaxies inside them \\citep{Mayer08,Ceverino08}.  The simulations have a fine resolution of $\\sim 100$\\,pc and include realistic treatment of gas and stellar feedback, which is important for the survival of a bar. Bars form relatively late: well after the last major merger ($z\\approx 1-2$ for normal spiral such as our Milky Way), when a collision of gas rich galaxies brings lots of gas with substantial angular momentum to the central disk galaxy. As the disk accretes the cold gas from the halo, it forms a new generation of stars and gets more massive. At some stage, the disk becomes massive enough to become unstable to bar formation.  Once the stellar bar forms, it exists for the rest of the age of the Universe.  The cosmological simulations are still in preliminary stages, and it is likely that many results will change as they become more accurate and the treatment of the stellar feedback improves.  However, existing cosmological simulations already show us the place and the role of traditional $N$-body simulations of barred galaxies, which start with an unstable stellar disk. It was not clear whether and how this happens in the real Universe. Now the cosmological hydrodynamic simulations tell us that this is somewhat idealistic, but still a reasonable setup compatible with cosmological models.  Simulations of bars play an important role for understanding the phenomenon of barred galaxies \\citep[e.g.,][]{Miller78,Sellwood80,lia03,vk03}. Numerical models successfully account for many observed features of real barred galaxies \\citep{Bureau05,Bureau06,Beaton07}. So, it is very important to assess the accuracy of those simulations.  Recent disagreements between results of different research groups \\citep{vk03,od03,sd06} prompted us to undertake a careful testing of numerical effects and to compare results obtained with different codes. This type of code testing is routine in cosmological simulations \\citep{Frenk99,Heitmann05,Heitmann07}, but it never has been done before for bar dynamics. Testing and comparison of numerical codes is important for validating different numerical models. It was instrumental for development of precision cosmology. It is our goal to make such tests for $N$-body models of barred galaxies.  We use four different popular $N$-body codes: ART \\citep{KKK97}, Gadget-1 and Gadget-2 \\citep{volker01,volker05}, and Pkdgrav \\citep{wsq04}.  ART is an Adaptive Mesh Refinement code that reaches high resolution by creating small cubic cells in areas of high density.  Gadget and Pkdgrav, on the other hand, are TREE codes that compute forces directly for nearby particles and use a multipole expansion for distant ones. We use the codes to run a series of simulations using the {\\it same initial conditions} for all codes.  We also address another issue: the effects of disk thickness on the structure and evolution of bars. Only recently have the simulations started to have enough mass and force resolution  to resolve the vertical height of stellar disks. We use different codes to show that the disk height plays an important and somewhat  unexpected role.  One of the contentious issues in the simulations of bar dynamics is the angular speed and the structure of bars in massive dark matter halos.  The amount and the rate at which bars slow down is still under debate. \\citet{ds98,ds00} find in their massive halo models, i.e., those for which the contributions of the disk and the halo to the mass in the central region are comparable, that the bar loses about 40\\% of its initial angular momentum, $L_z$, in $\\sim 10\\ \\gyr$. However, in simulations with much better force resolution and a more realistic cosmological halo setup, \\citet{vk03}  and  \\citet{cvk06} find a decrease in $L_z$ of only 4--8\\% in $\\sim 6\\ \\gyr$.  \\citet{ds98,ds00} also find that bars do not significantly slow down for lower density halos.  \\citet{vk03} presented bar models in which stellar disks were embedded in a CDM Milky-Way-type halos with realistic halo concentrations $c \\sim 15$, where $c$ is the ratio of the virial radius to the characteristic radius of the dark matter halo. These simulations were run with the ART code  with high spatial resolution of 20-40 pc. The bars in the models were rotating fast for billions of years. \\citet{vk03} argued that slow bars in previous simulations were an artifact of low resolution. \\citet{sd06} used initial conditions of one of the models of \\citet{vk03} and run a series of simulations using their hybrid, polar-grid code.  They found, in most cases, a different evolution than that reported in Valenzuela \\& Klypin. In particular, contrary to Valenzuela \\& Klypin's results, they did not find that the bar pattern speed is almost constant for a long period of time.  They attribute the differences to the ART refinement scheme.  While \\citet{vk03} mention numerical effects (lack of force and mass resolution) as the cause for excessive slowing down of bars in earlier simulations, there was another effect, which was not noticed by \\citet{vk03}: Their disk was rather thin, with a scale-height $h_z$ of only 0.07 of the disk scale-length $R_d$. This should be compared with $h_z \\sim 0.2\\ R_d$ used in most studies of stellar bars \\citep[e.g.,][]{lia02,lia03,m-v06}.  Models of \\citet{ds98} have $h_z = 0.1 R_d$, but the resolution of their simulations was grossly insufficient to resolve this scale. \\citet{od03} used $h_z = 0.1 R_d$ for a model, which had very little dark matter in the central disk region: $M_{\\rm dm}/M_{\\rm disk} \\approx 1/4$ inside radius $R=3R_d$. Dependence of bar speed on disk thickness was noted by \\citet{lia00}: thicker disks result in slower bars.   The remainder of this paper is organized as follows. In section~\\ref{sec:height} we give a detailed review of available data on disk scale heights and present a simple analytical model for the relation of the disk scale length with the disk scale height. We describe our numerical models in Section~\\ref{sec:Models}. In Section~\\ref{sec:Numerics} we give a brief description of the codes and present analysis of numerical effects.  Main results are presented in Section~\\ref{sec:Results}. We summarize our results in Section~\\ref{sec:Discussion}. ",
        "conclusions": "\\label{sec:Discussion}  \\begin{figure} \\includegraphics[width=0.475\\textwidth]{DensityMapxySK.5.ps} \\includegraphics[width=0.465\\textwidth]{FreudP.ps} \\caption[fig:Surface]{The iso-contours of the surface stellar density for model \\kgthree~ (top panel) and the Milky Way galaxy (bottom panel).  The \\kgthree~ model was re-scaled to have the evolved disk scale length  3~kpc, which is close to the scale length of our Galaxy. The bottom panel shows one of the models from \\citet[][fig. 14, right panel]{Freudenreich98}. The model represents a multi parameter fit to the COBE DIRBE maps of the near-infrared light coming from central regions of our Galaxy. The small circle shows position of the Sun.  The \\kgthree~ model reproduces the length and the flattening of the Milky Way bar.} \\label{fig:SurfaceMW} \\end{figure}  We find that numerical effects -- the mass and force resolution, and the time integration of trajectories -- can significantly alter results of simulations. The time-step of integration must be small enough to allow accurate integration of expected trajectories. Accuracy of energy conservation can give a misleading impression that the simulation is adequate, while it actually makes substantial errors in position of orbits.  Our experiments with realistic orbits in models with flat rotation curves indicate that even the best available integration schemes require $\\sim 2000$ time-steps per orbit. This requirement is valid for variable-step schemes with more steps required for a constant-step schemes. Numerical tests with full-scale dynamical models confirm the condition.  This condition results in very small time-steps of $dt \\sim 10^4$~yrs. If one uses dimensionless units defined by scaling $G=1$, $M_{\\rm disk}=1$, $R_d =1$, then the dimensionless time step should be $\\tilde dt\\approx 5\\times 10^{-4}$. This should be compared with typically used $\\tilde dt \\approx 0.01$ \\citep[e.g.,][]{lia02}. This time-step would be insufficient for integration of models with central mass concentration presented in this paper \\footnote{Models studied by \\citet{lia02} are less concentrated and do not require a small time-step.}.  A constant time-step $dt =5\\times 10^5$~yrs used by \\citet{Widrow08} is too large for models of the Milky Way galaxy studied in their paper.  The mass and related with it the force resolution also play important role. As Figure~\\ref{fig:ResCon} shows, a low force resolution produces a less dense central region, which results in a long and a very massive bar. The same model run with better force resolution produces a shorter, less massive, and faster rotating bar.  Once the necessary numerical conditions are fulfilled, the codes produce practically the same results.  We do not find any systematic deviations between the results obtained with TREE codes (Gadget and Pkdrgav) and with the Adaptive-Mesh-Refinement (ART) code.  Disk height is an important parameter, which is often ignored in models of barred galaxies. In the models, which we consider in this paper, the disk height determines the global properties of the bars. Figure~\\ref{fig:Surface} shows surface density maps of the models with different initial disk thickness. Models with thin disks produce short bars with $R_{\\rm bar}\\approx R_d$, which rotate relatively fast: ${\\cal R} =1.2-1.4$ and which show very little decline of the pattern speed. Models with  thick disk produce long and slow rotating bar.  In order to facilitate the comparison with the our Galaxy, we re-scaled models to have the evolved disk scale length 2.65~kpc and to have the circular velocity at the solar distance 220~km/s. Any $N$-body system has two arbitrary scaling factors, which can be used to scale the system.  Having scaled the models to fit the disk scale length, we can compare other parameters of the models. Because the simulations with thin disks produce reasonable models, we use one of the models (\\kgthree) and compare it with the Milky Way. Table~\\ref{tab:MW} gives a list of some parameters.  Figure~\\ref{fig:SurfaceMW} compares the surface density maps of the Milky Way \\citep{Freudenreich98} and the \\kgthree~ model. These comparisons show that the model fits the Milky Way reasonably well.  We suggest  that the disk height is only an indicator of a more fundamental property -- the phase-space density in the central ($R<R_d$) region. In our models the phase-space density is uniquely related with the disk height.  Initially in our models the disk height is low and the stability parameter $Q=1.3-1.8$ is constant across the disk. Thus, the phase-space density in the central region is high, and subsequent evolution brings that highly compressible stellar fluid close to the center, where it forms a nearly flat circular velocity curve. The later, as we speculate, is responsible for arresting the growth of the bar.  This relation between the disk height and the central phase-space density may not be true in general case. For example, \\citet{od03} and \\citet{Widrow08} consider models with large central $Q$ and, thus, with a low central phase-space density. Indeed, in their models bars slow down substantially.   The nearly constant pattern speed of bars in the thin disk models is somewhat puzzling. A bar is a very massive non-axisymmetric object, which rotates inside a non-rotating dark matter halo. As such, one might expect that it should experience dynamical friction and slow down. This is why results of \\citet{vk03}, which showed models with little slowing down of bars, were met with skepticism. Simulations presented in this paper confirm the findings of \\citet{vk03} and show that they cannot be related with numerical problems.  Orbital resonances may be responsible for the observed slow dynamical friction. The resonances in barred galaxies have been extensively studied in recent years using $N$-body simulations \\citep[e.g.,][]{lia03,cvk06,Ceverino07,Weinberg07}. It is now well established that a large fraction of {\\it stellar particles} are in resonance with the bar. Some fraction of {\\it dark matter particles} are also in resonance \\citep{cvk06,Ceverino07}, and these are very important for the dynamical friction between the stellar bar and the dark matter.  There are two types of (exact) resonances in an autonomous Hamiltonian dynamical system: elliptic and hyperbolic \\citep{Arnold}. Orbits, which are close to an elliptic resonance librate around the exact resonance and have the structure of a simple pendulum \\citep[][Sec.2.4]{Lich83}, \\citep[][Sec. 8]{Murray99}. Hyperbolic resonances are points on intersections of separatrixes, which separate domains of elliptical resonances. Orbits close to a hyperbolic resonance are unstable and tend to migrate away from the resonance, while orbits close to an elliptical resonance are stable.  \\citet{Ceverino07} study in detail the resonant orbits in simulations of barred galaxies. It appears that the orbits belong to elliptical resonances. These orbits track their resonance: the orbits do not evolve if the resonance does not move, and they follow the resonances if it gradually migrates in the phase space. Thus, the elliptical resonances trap the orbits: if an orbit for whatever reason happens to appear in the domain of the resonance, it will have a tendency to stay  with the resonance. This phenomenon is thought to be common in Solar System dynamics \\citep[e.g.,][]{Malhotra93}.  Resonance trapping explains why the simulations show maxima in the distribution of orbital frequencies at the positions of resonances.  \\citet{cvk06} investigate another aspect of the resonant interaction between the stellar bar and the dark matter. They found that the dark matter particles, which are in resonance with the stellar bar, themselves form a bar, which rotates with the same angular speed and has a very small lag angle ($\\sim 10^o$) as compared with the stellar bar. \\citet{cvk06} argue that the interaction between the dark matter and the stellar bars is the main mechanism for the dynamical friction between the disk and the dark matter. The near alignment of the dark matter and stellar bars means that their interaction is minimized by the resonances.  These results indicate that the resonant interaction between the stellar bar and the dark matter is mostly due to elliptical resonances, and, thus, has a tendency to {\\it minimize} the transfer of the angular momentum from the disk to the dark matter.  Following orbits in such resonances emphasizes the need for conservative time-steps in $N$-body simulations.  Also subtle changes in the underlying global potential could change the relative number of orbits in these resonances leading to disparate results for the slowing of the bar."
    },
    "0808/0808.1611_arXiv.txt": {
        "abstract": "The starburst galaxy M82 shows a system of H$\\alpha$-emitting filaments which extend to each side of the galactic disk. We model these filaments as the result of the interaction between the winds from a distribution of Super Stellar Clusters (SSCs). We first derive the condition necessary for producing a radiative interaction between the cluster winds (a condition which is met by the SSC distribution of M82). We then compute 3D simulations for SSC wind distributions which satisfy the condition for a radiative interaction, and also for distributions which do not satisfy this condition. We find that the highly radiative models, that result from the interaction of high metallicity cluster winds, produce a structure of H$\\alpha$ emitting filaments, which qualitatively agrees with the observations of the M82, while the non-radiative SSC wind interaction models do not produce filamentary structures.  Therefore, our criterion for radiative interactions (which depends on the mass loss rate and the terminal velocity of the SSC winds, and the mean separation between SSCs) can be used to predict whether or not an observed galaxy should have associated H$\\alpha$ emitting filaments. ",
        "introduction": "Supergalactic winds are a very important mechanism in the evolution of the Universe. Lynds \\& Sandage (1963) found a large outflow in M82 and gave the first definition of a starburst. They proposed that material was ejected from the nuclear regions as the consequence of a SN explosion. Different analytical  models have been put forward to explain the so-called supergalactic winds (SGWs), see, e. g., the work of Chevalier \\& Clegg (1985), Tomisaka \\& Ikeuchi  (1988), Heckman et al (1990), Leitherer \\& Heckman (1995), Strickland \\& Stevens (2000), Cant\\'o et al (2000), Tenorio-Tagle et al. (2003), and Cooper et al. (2008). In the nearest starbursts, they have shown that the supergalactic winds (SGW) form through the collective effects of many individual stellar cluster winds, which are in turn  formed by the collective effect of many individual stellar winds.  The best-studied starburst galaxy is M82. It has an extended, biconical filamentary structure in the optical (Shopbell \\& Bland-Hawthorn 1998; Ohyama et al. 2002). This optical emission is embedded in a pool of soft  X-ray emission, detected with the Chandra X-Ray Observatory (Griffiths et al. 2000) as well as with XMM-Newton  (Stevens et al. 2003). This X-ray emission extends several kpc away from the nuclear region, into the IGM.  With the Hubble Space Telescope (HST), it has been possible to isolate the building blocks of some starbursts:Super Stellar Clusters (SSCs). This objects are very massive and dense stellar clusters, with masses in excess of $10^4~\\rm{M_{\\odot}}$ enclosed in radii of $3$--$10$ pc. They are young, and typically contain from $\\sim 100$ up to $\\sim 10^4$ OB stars. In M82, the SSCs $H{\\alpha}$ luminosities are in the range of (0.01-23)$\\times 10^{38}~\\rm{erg~s^{-1}}$  (Melo et al 2005). These authors  catalogued 197 SSCs in M82 with masses in the $10^4< \\rm{M/M}_{\\odot}<10^6$ range. This exceptional density of massive clusters ($620~\\rm{kpc}^2$) produces a suitable scenario for a study of the interaction between cluster winds.  The differences in the populations of SSCs of several starburst  galaxies can help us to understand which properties are more important for producing filaments. For instance  NGC253, which does not show a clear filmentary structure (Rice 1993), has a significantly lower SSC mass loss rate than M82. However, other authors claim that the lack filamentary structure in this galactic wind is simply a result of not having deep enough H$\\alpha$ observations of this objects (Hoopers et al. 1995).   In NGC253 Melo (2005) catalogued a total of 48 SSCs with an average distance between them of $31~\\rm{pc}$, in contrast with the mean separation between SSCs of $12~\\rm{pc}$ found in M82. In M82, a rich structure of filaments, extending out from the galactic plane, is present. There are many other examples such as NGC1569 (Anders et al. 2004; Westmoquette, Smith \\& Gallagher III 2008) and M83 (Harris et al. 2001), that have high density of SSCs with similar masses to the ones of M82.  Tenorio-Tagle et al. (2003) showed that the interaction of stellar winds from a collection of energetic, neighbouring SSCs could form a filamentary structure similar to the one observed in M82. These authors propose the formation of such filaments by means of supernovae explosions that would expel a huge amount of heavy elements into the Interstellar Medium (ISM), thus enhancing the radiative cooling of the outflows.  In their model, Tenorio-Tagle et al. (2003) explained the SGW structure as a  consequence of the interaction of winds from very close SSCs, in which stationary oblique shocks are responsible for shaping  the material into dense and cold, kiloparsec-sized  filaments. More recently, Cooper et al. (2008) presented a numerical study of a starburst-driven galactic wind. Their setup consisted of  a series of neutral dense clouds placed in the galaxy disk. The dense clouds are swept up by the main shock wave, leaving behind a filamentary structure.  In the present paper, we discuss models for the formation of filamentary structures as the result of the interaction between winds from a system of many clusters with a disk-like spatial distribution. In these models we consider only the flow resulting from the wind interactions, and do not include the effect of the stratified ISM present in the region of the galactic disk.  The paper is organized as follows. In \\S 2, we study which combinations of parameters (mass loss rate ${\\dot M}_w$, cluster wind velocity $v_w$, and separation between nearby star clusters $D$) yield a highly radiative SGW flow. In \\S 3 we describe a series of 3D simulations of the interaction of cluster winds in this highly radiative regime. We present predictions of the emission in the optical and X-ray wavelength ranges, and compare them with observations of galactic winds (i.e. M82, NGC253, etc.) in \\S 4. Finally, in \\S 5 we present our conclusions. ",
        "conclusions": "We study the formation of filaments in the galactic winds driven by young SSCs in regions of high star formation. We study models in which the filaments are produced solely by the interaction between winds from SSCs, and do not consider the possible effects of the stratified or clumped ISM present in the galactic disk. In our models, the winds propagate into a low density, homogeneous environment, which could correspond to the intergalactic medium.  We have shown how a  $(\\dot{M}_w/D)$ vs. $v_w$ diagram can be used to determine if the interaction between the winds of SSCs are radiative, in which case cold filaments can be formed. With the $(\\dot{M}_w/D)$ vs. $v_w$ diagram (Figure~1) one can diagnose if a galaxy with a high star formation rate will show filaments in optical emission lines, or only a galactic wind with X-ray emission. At the same time, given an observation of a galactic wind with a rich filamentary structure, one could predict the number of SSCs which are contained within the observed galaxy.  We computed three models of dense starbursts (M1-M3) based on the observed parameters of M82, and a fourth model (M4), based on NGC253. The models are based on cluster distributions with mean separations between clusters similar to the observed values, but with spatial extents which are considerably smaller than the total extent of the starburst regions. The distributions used in the model could correspond to individual ``starburst clumps'' such as the ones observed in M82 by Westmoquette et al. (2007).  For the parameters of M82 we computed four models (M1-M3, and M1a) with different metallicities (from solar to 10 times solar). For the parameters of NGC253 we computed two models (M4 and M4a) with a 10 times solar metallicity. Models M1-M4 asssumed a mass loss rate consistent with the average mechanical luminosity of the SSCs. Models M1a and M4a are obtained with rather extreme mass loss rates, only achievable during a short-lived phase in which the input from SN dominates over stellar winds.  For the models based on M82, we see the formation of cold, H$\\alpha$-emitting, filaments that extend out from the plane of the the galactic disk in the models with metallicities $\\gtrsim 5~Z_{\\odot}$ (models M2 and M3). We also see formation of dense filaments for the models with extreme mass loss rate at the supernova phase of the clusters (models Ma1 and M4a), where a more conservative mass loss rate did not yielded filaments (M1 and M4).   For the model based on NGC253 (M4), no filaments are produced, even for a $10~Z_{\\odot}$ metallicity. This is consistent with the present observations of this object, in which no H$\\alpha$-emitting filaments have been detected. The prediction obtained from our models would be that with the distribution of SSCs of NGC253 (in pre or post SN phase), no H$\\alpha$ filaments should exist. However, we have shown that during the SN phase (model M4a and section~3.1)  the formation of cold filaments is possible in such high metallic  winds.  Hoopes et al. (2005) claim that the lack of evidence of filamentary structure in the galactic wind of NGC253 could simply be the result of not having deep enough H$\\alpha$ maps of this object. In the future we will see whether or not this result is confirmed by deeper observations.  We have also computed X-ray maps in the energy range of $0.3$ to $2~\\rm{keV}$. For the highly radiative models (M2 and M3), we find that the X-ray emission is concentrated in a region close to the galactic plane. In order to reconcile this result with the observation of extended X-ray emission in M82, we would have to assume that the emission comes from another component in the flow, which could be the shock of SGW against the intergalactic medium.  Finally, we have run two models (M2b and M3b) with the parameters based on M82 (and abundances of $5~Z_{\\odot}$ and $10~Z_{\\odot}$, respectively) with lower computational resolutions but with larger spatial extents. From these runs, we find that the models produce a filamentary H$\\alpha$ morphology that extends out to $\\sim 0.5\\to 1$~kpc on each side of the galactic plane.  We end by noting that in this paper we prove that the radiative interaction between the winds from a disk-like distribution of identical SSCs (super stellar clusters) does lead to the production of dense filaments extending out to $100\\to 1000$~pc away from the galactic plane. These structures produce H$\\alpha$ emission that might correspond to the filaments observed in starburst galaxies such as M82 and NGC1569. Of course, in our models we do not consider many elements that are likely to be important in starburst galaxies. Important among these might be: \\begin{itemize} \\item the presence of a dense, stratified galactic disk interstellar medium, \\item clumpiness in this medium, \\item the existence of a distribution of sizes and ages for the SSCs, \\item the effect of galactic mass loading.  \\end{itemize} There is therefore a large field for future theoretical studies on the formation of filamentary strutures in starburst galaxies."
    },
    "0808/0808.1757_arXiv.txt": {
        "abstract": "This paper deals with the gereral subject of particle acceleration in solar flares with the aim of describing the long standing puzzle of extreme enrichments of $^3$He ions in solar energetic particles associated with solar flares. ",
        "introduction": "The aim of this chapter is to investigate to what extent theoretical models of acceleration of particles can explain the SEP observations described in previous chapters. In the next section we  present a general description of the possible acceleration mechanisms for SEPs and other space and astrophysical sources and indicate that plasma waves or turbulence play a major role in all these processes. We then describe the model of stochastic acceleration by plasma turbulence in some detail and review its successes in describing observations of solar flares, in particular the radiative signatures associated with accelerated electrons and protons. In \\S 3 we consider the problem of observed enhancements of \\3he and heavy elements seen in many SEPs.  Here we first present a brief review of some of the earlier models proposed for production of these enhancements and then compare predictions of the stochastic acceleration model with the observations. A brief summary and discussion is presented in \\S 4. ",
        "conclusions": ""
    },
    "0808/0808.0961_arXiv.txt": {
        "abstract": "Using direct simulations of hydromagnetic turbulence driven by random polarized waves it is shown that dynamo action is possible over a wide range of magnetic Prandtl numbers from $10^{-3}$ to 1. Triply periodic boundary conditions are being used. In the final saturated state the resulting magnetic field has a large-scale component of Beltrami type. For the kinematic phase, growth rates have been determined for magnetic Prandtl numbers between 0.01 and 1, but only the case with the smallest magnetic Prandtl number shows large-scale magnetic fields. It is less organized than in the nonlinear stage. For small magnetic Prandtl numbers the growth rates are comparable to those calculated from an alpha squared mean-field dynamo. In the linear regime the magnetic helicity spectrum has a short inertial range compatible with a $-5/3$ power law, while in the nonlinear regime it is the current helicity whose spectrum may be compatible with such a law. In the saturated case, the spectral magnetic energy in the inertial range is in slight excess over the spectral kinetic energy, although for small magnetic Prandtl numbers the magnetic energy spectrum reaches its resistive cut off wavenumber more quickly. The viscous energy dissipation declines with the square root of the magnetic Prandtl number, which implies that most of the energy is dissipated via Joule heat. ",
        "introduction": "Many astrophysical plasmas are turbulent and tend to be magnetized. The magnetic fields can have typical length scales that are either larger or smaller than that of the energy-carrying eddies. We speak then correspondingly of large-scale or small-scale dynamos. Small-scale dynamos can already work in statistically mirror-symmetric isotropic homogeneous turbulence, whereas large-scale dynamos require in general a departure from parity-invariant or mirror-symmetric flows. The excitation conditions of small-scale dynamos depend sensitively on the value of the magnetic Prandtl number, i.e.\\ the ratio of kinematic viscosity to magnetic diffusivity, $\\Pm=\\nu/\\eta$, where $\\nu$ is the kinematic viscosity and $\\eta$ the magnetic diffusivity. This sensitivity is related to the fact that in the {\\it kinematic} regime the spectral magnetic energy is peaked at the resistive scale. As was pointed out originally by Rogachevskii \\& Kleeorin (1997), and more recently by Boldyrev \\& Cattaneo (2004), the slope of the kinetic energy spectrum is important for the onset of small-scale dynamo action. It matters therefore whether the resistive scale lies in the viscous range ($\\Pm\\approx1$), within the inertial range ($\\Pm\\la0.1$), or right within the range where the bottleneck occurs ($\\Pm\\approx0.1$). The bottleneck effect refers to the spectral subrange just before the dissipation range where the kinetic energy spectrum is shallower than in the inertial range. Within the bottleneck range the velocity increments diverge even more strongly with decreasing separation than in the inertial range, making dynamo action harder still. Indeed, for small values of $\\Pm$ the critical value of the magnetic Reynolds number above which small-scale dynamo action occurs increases therefore sharply toward $\\Pm=0.1$ (Schekochihin et al.\\ 2005), and then decreases slightly for $\\Pm<0.05$ (Iskakov et al.\\ 2007). However, even with the computing power available today, direct simulations of small-scale dynamo action are still only marginally possible at such small values of $\\Pm$.  For certain types of flows dynamo action is easier to achieve even though $\\Pm$ is small. The Taylor-Green flow is an example where the critical value of the magnetic Reynolds number becomes constant for $\\Pm<0.1$ (Ponty et al.\\ 2004, 2005). In this flow there can be large-scale patches with finite kinetic helicity of opposite sign. A completely different example is fully helical turbulence where the excitation condition for dynamo action is virtually unchanged as $\\Pm$ decreases from 1 to 0.1 (Brandenburg 2001, hereafter B01). For an ABC-flow dynamo, Mininni (2007) found a weak dependence of the threshold value of the magnetic Reynolds number $\\Rm$ on $\\Pm$. For $\\Pm<0.1$ the threshold value seemed to become asymptotically independent of $\\Pm$ and dynamo action was demonstrated for values of $\\Pm$ down to $5\\times10^{-3}$.  In many astrophysical bodies, $\\Pm$ is indeed rather small (around $10^{-5}$). Such systems still possess dynamo action and can have large-scale magnetic fields. It is likely that such systems belong to the second class of systems where the excitation conditions are not drastically altered toward small values of $\\Pm$. Indeed, large-scale magnetic fields are found regardless of whether $\\Pm$ is small (e.g., the Sun and other stars with outer convection zones, as well as planets) or large (e.g., in spiral galaxies, because of their low densities).  The purpose of this paper is to point out that a strong $\\Pm$ dependence does not occur in systems where the magnetic field generation is predominantly due to a large-scale dynamo. Such systems have been studied in idealized settings such as periodic boxes using explicit forcing functions for driving the turbulence. This has significant advantages in that periodic boundary conditions can be used, energy spectra are easily computed and, most importantly, isotropy and homogeneity eases comparison with turbulence theory. A disadvantage is that the magnetic helicity can only change on resistive timescales, which slows down the saturation (B01).  With these provisions in mind, we consider now simulations of maximally helical turbulence in triply periodic boxes where we keep in most cases the fluid Reynolds number, $\\Rey=\\urms/\\nu\\kf$, constant and vary the magnetic Reynolds number, $\\Rm=\\urms/\\eta\\kf$, and thereby $\\Pm$ ($\\equiv\\Rm/\\Rey$). Here, $\\urms$ is the rms velocity of the turbulence and $\\kf$ is the forcing wavenumber. According to B01 the dynamo should be excited whenever the domain is large enough (2--3 times larger than the forcing scale) and the magnetic Reynolds number exceeds unity ($\\Rm\\geq1.1...1.4$ or so). This was confirmed for magnetic Prandtl numbers as low as 0.1. In the present work we consider kinematic dynamo action down to values of $\\Pm=10^{-2}$ and nonlinear saturated dynamos down to $\\Pm=10^{-3}$. ",
        "conclusions": "In many astrophysical bodies the magnetic Prandtl number is small, while in most simulations its value is chosen to be close to unity. As we have shown here, this mismatch is of relatively minor consequence for large-scale dynamos that are driven by helical forcing.  In the nonlinear stage, the velocity and magnetic field patterns are remarkably independent of the value of $\\Pm$. The only thing that changes is the length of the inertial range. A small magnetic Prandtl number simply means that the magnetic energy spectrum turns into the dissipation range more quickly than the kinetic energy spectrum. It also means that essentially all the energy is dissipated via Joule heat. This was recently also demonstrated by Mininni (2007). One reason is that the case of fully helical turbulence studied in the present paper is a particu\\-larly simple one, because it leads to uniform mean-field dynamo action with large-scale pattern formation covering the entire domain. In this paper we have seen that, at least for values of $\\Rey$ up to 4400, the dynamics of this large-scale pattern, i.e., of the large-scale magnetic field, is quite independent of how long the inertial range of the turbulence is. In the absence of helicity, there is only small-scale dynamo action, which is driven by the dynamics at the smallest possible scale, i.e.\\ the resistive scale. In that case it does matter what the dynamics of the turbulence is at that scale. However, in that case it has not yet been possible to find dynamo action for values of $\\Pm$ down to the values considered here. Nevertheless, it is possible that even in that case there is an asymptotic regime for large enough values of $\\Rm$ where the dynamics of the magnetic field is independent of the value of $\\Pm$, even though this regime is not yet accessible with present day computers.  To estimate the value of the magnetic Prandtl number in dense astrophysical bodies, one has to use the {\\it Spitzer} formulae for $\\eta$ and $\\nu$. The resulting magnetic Prandtl number is (e.g., Brandenburg \\& Subramanian 2005b) \\EQ \\Pm=1.1 \\times10^{-4} \\left({T\\over10^6\\K}\\right)^4 \\left({\\rho\\over0.1\\g\\cm^{-3}}\\right)^{-1} \\left({\\ln\\Lambda\\over20}\\right)^{-2}, \\EN so at the bottom of the solar convection zone the magnetic Prandtl number is clearly rather small ($\\sim10^{-4}$). Nevertheless, simulations of solar and stellar dynamos available so far $\\Pm$ are set to values of the order or unity. Although we have shown here that the resulting large-scale fields are similar to the more realistic case of small values of $\\Pm$, an important difference is that the small-scale dynamo may be more pronounced when $\\Pm$ is of order unity. In practice this means that a positive detection of dynamo action in a simulation might not necessarily be relevant for understanding the Sun, unless suitable conditions for the excitation of large-scale dynamo action are also met. On the other hand, once the large-scale dynamo is really excited, and if it is fully saturated, it is then quite feasible to lower the value of $\\Pm$ significantly---without losing the large-scale dynamo. In fact, lowering $\\Pm$ in a saturated large-scale dynamo means that most of the energy will be dissipated via Joule heating, and that the kinetic energy cascade only carries a small fraction of the total energy. This allows us to  increase the value of $\\Rey$, and hence to decrease the viscosity and thereby the value of $\\Pm$ even further. Simulations of large-scale dynamo action in turbulent convection (K\\\"apyl\\\"a et al.\\ 2008) provide one example where it is indeed feasible to lower $\\Pm$, although in that case the system is not uniform and so energy dissipation via Joule heating is only possible in those locations where the dynamo is strong enough (P.\\ J. K\\\"apyl\\\"a 2008, private communication)."
    },
    "0808/0808.0685_arXiv.txt": {
        "abstract": "{We explore the possible observational signatures of different types of kink modes (horizontal and vertical oscillations in their fundamental mode and second harmonic) that may arise in coronal loops, with the aim of determining how well the individual modes can be uniquely identified from time series of images. A simple, purely geometrical model is constructed to describe the different types of kink-mode oscillations. These are then `observed' from a given direction. In particular, we employ the 3D geometrical parameters of 14 TRACE loops of transverse oscillations to try to identify the correct observed wave mode. We find that for many combinations of viewing and loop geometry it is not straightforward to distinguish between at least two types of kink modes just using time series of images. We also considered Doppler signatures and find that these can help obtain unique identifications of the oscillation modes when employed in combination with imaging. We then compare the modeled spatial signatures with the observations of 14 TRACE loops.  We find that out of three oscillations previously identified as fundamental horizontal mode oscillations, two cases appear to be fundamental vertical mode oscillations (but possibly combined with the fundamental horizontal mode), and one case appears to be a combination of the fundamental vertical and horizontal modes, while in three cases it is not possible to clearly distinguish between the fundamental mode and the second-harmonic of the horizontal oscillation. In five other cases it is not possible to clearly distinguish between a fundamental horizontal mode and the second-harmonic of a vertical mode. ",
        "introduction": "The solar corona is characterized by highly dynamic loop-like structures believed to outline magnetic field lines. Various types of coronal loop oscillations have been observed for decades in radio, visible, EUV, and X-rays \\citep[see, e.g., reviews by][]{nak03, nak05, asc04, wan04a, wan05a}. Recently, temporally and spatially resolved transverse and longitudinal oscillations have been detected in coronal loops in the high-resolution observations by the Transition Region And Coronal Explorer (TRACE) and the Solar and Heliospheric Observatory (SOHO). Compressible longitudinal waves in cool ($\\sim1$ MK) loops were first detected by the EUV Imaging Telescope (EIT) on board SOHO \\citep{ber99}, which were confirmed by TRACE observations \\citep[e.g.][]{dem00,dem02a,dem02b}, and identified as propagating slow magnetoacoustic waves \\citep{nak00}. Global kink-mode oscillations were first found in cool ($\\sim$ 1~MK) coronal loops from TRACE EUV imaging observations \\citep{asc99,asc02,nak99}. They were seen as spatial displacements with periods of 3$-$5 minutes and were apparently excited by flares or erupting filaments \\citep{asc02,sch02}. Later, strongly damped standing slow-mode waves were discovered with the Solar Ultraviolet Measurement of Emitted Radiation (SUMER) spectrometer on-board SOHO by virtue of their Doppler signatures in the flare lines Fe~{\\small XIX} and Fe~{\\small XXI} \\citep[e.g.][]{wan02, wan03a, wan03b, ofm02}. These oscillations are usually set up immediately following a pulse of hot plasma flowing from one of the loop footpoints, associated with small flares or even micro-flares \\citep{wan05b, wan06}. Extensive studies on coronal loop oscillations in recent years have made {\\it MHD coronal seismology} a practical, new tool for the determination of previously poorly known parameters of the coronal structure \\citep[see, e.g.,][for reviews]{rob84, rob03, nak05}. For example, the measured oscillation period has been applied to determine the mean magnetic field strength of coronal loops in case of either kink mode \\citep{nak01, ver04} or slow mode oscillations \\citep{wan07}. \\begin{figure*} \\centering \\includegraphics[width=16cm]{10230f01.eps} \\caption{ A sketch of the horizontal and vertical oscillations and their second harmonics. (a) The fundamental horizontal mode, (b) the second harmonic horizontal mode, (c) the fundamental vertical mode, and (d) the second harmonic vertical mode. The top, middle and bottom rows correspond to views from the side, top and top-left, respectively.  } \\label{sketch} \\end{figure*}  \\citet{wan04b} first suggested that the longitudinal curvature of loops may lead to two types of kink modes (i.e., horizontal and vertical modes) in a coronal loop based on observations. The motion is termed a {\\it horizontal} mode if the loop rocks perpendicularly to the loop plane. If the oscillation is polarized in the loop plane and manifests expanding and shrinking motions, we refer to it as a {\\it vertical} mode. The horizontal kink mode oscillation for an individual loop has been numerically studied by \\citet{miy04} with a 3D MHD model. They set an initial horizontal velocity field near the loop top as a perturbation to induce the oscillation. The vertical mode oscillation can be excited either by a pressure pulse below the loop apex\\citep[e.g.][] {sel05, sel06, sel07} or by a velocity driver in one of the footpoints \\citep[e.g.][] {bra05, bra06, ver06a, ver06b} as simulated in a 2D arcade model. Very recently, \\citet{mcl08} performed a 3D nonlinear MHD simulation of wave activity in active regions in which individual loop density structure is included. They found that the impact of the fast wave impulsively excites both horizontal and vertical loop oscillations. \\citet{van04} first studied the eigenmodes of a curved loop and showed that curvature only slightly affects the damping times of the observed kink oscillations in comparison with a straight tube. By numerically solving the eigenvalue problem of a similar model, \\citet{ter06} found that curvature and density structuring introduce preferential directions of oscillation showing the existence of two kink eigenmodes with horizontal and vertical polarizations, but their frequencies and damping rates are very similar. Their calculation supports the initial notion of \\citet{wan04b}.  In several examples, the second harmonic was found together with the fundamental mode via time series analysis \\citep{ver04, van07}. The first spatially resolved example of the second harmonic of transverse loop oscillations was recently identified by \\citet{dem07}. \\citet{dia06} and \\citet{dia06b} analytically studied the vertical kink mode and its second harmonic (they termed it the {\\it swaying} kink mode). Observations of multiple periodicities in kink loop oscillations \\citep[e.g.,][]{ver04} have been used to extract information on the internal longitudinal structuring of coronal loops by \\citet{and05, mce06} and has been demonstrated to have high potential value for coronal seismology. In addition, \\citet{ter07} theoretically studied the excitation of the eigenmodes of coronal loops due to an initial disturbance. Their results suggest that longitudinal harmonics of kink modes are in principle more easily excited than azimuthal harmonics (fluting modes).  \\begin{figure*} \\centering \\includegraphics[width=12cm]{10230f02.eps} \\caption{Positions of the 14 studied oscillating TRACE loops on the solar disk.} \\label{loops} \\end{figure*}  Therefore, differentiating between various types of kink modes and their higher harmonics in observations will improve the applications of coronal seismology in precisely determining the unknown physical parameters of the coronal plasma. \\citet{asc02} identified oscillations in 26 coronal loops as kink modes. They did not distinguish between the different types of modes. \\citet{wan04b} have for the first time identified the vertical kink mode oscillation of a TRACE loop above the limb based on its spatial signature, which is distinct from the horizontal mode.  Our aim here is to explore observational signatures of different types of kink-mode oscillations based on the loop geometries measured by \\citet{asc02} and check if each observed oscillation can be identified with a given type of kink mode based on these signatures.  In Section 2, we describe the methods which are used to model the horizontal and vertical kink modes. In Section 3, we discuss their similarities and differences in signatures of the spatial pattern and Doppler shift. In Section 4, we discuss the simultaneous presence of multiple (harmonics) modes. Conclusions are summarized in Section 5. ",
        "conclusions": "We have modeled the geometric distortion to a simple circular loop produced by four types of kink modes (fundamental and second harmonic modes of horizontal and vertical oscillations) and considered if it is possible to distinguish between them for a given observational geometry. We analyzed in detail 14 selected TRACE loops observed to harbor transverse oscillations and found that they can be divided into three categories based on the location of the observer relative to the loop plane (viewing geometry). The apparent similarity and difference between the four types of kink modes in spatial signatures is found to depend strongly on the category. We also examined if the kink modes displaying a similar spatial signature can be distinguished based on their Doppler shift signatures.  Our main conclusions regarding the possibility of distinguishing between modes are outlined below (see also Table \\ref{cmptab}).  (1) For those loops whose loop plane is almost in the line-of-sight direction (i.e. loops in Category III), the fundamental and second-harmonic horizontal kink modes can be easily and uniquely identified based on the signature of evident lateral displacements of the loop. The fundamental and second-harmonic vertical kink modes show smaller displacements with similar signature in this viewing geometry and are difficult to detect and distinguish. It is also impossible to determine if a fundamental horizontal mode is combined together with a fundamental vertical mode in this case. For those loops located on or near the limb (e.g. Category II loops), the fundamental vertical kink mode as well as its combination with the fundamental horizontal mode are most easy to distinguish from the other types of kink modes. The fundamental horizontal and the second-harmonic vertical modes look similar, displaying a cross-over signature near the loop top. For those loops on the solar disk whose loop plane is inclined at an angle sufficiently different from $0^{\\circ}$ and $90^{\\circ}$ to the line-of-sight direction (e.g., the loops in Category I), the second-harmonic vertical mode is most easy to be uniquely identified based on the cross-over signature near the loop top. The fundamental horizontal and vertical oscillations may be distinguished from each other by considering more subtle effects, e.g., determining on which side of the loop relative to the loop apex the largest amplitude displacement is located and comparing it with a simple model of the type considered here.  However, it may be difficult to determine if the identified fundamental vertical mode is present together with the fundamental horizontal mode.  (2) The Doppler shift signatures provide additional constraints that can help distinguish between two types of kink modes which exhibit similar spatial oscillation signatures if the used spectrometers have a high resolution in Doppler velocity and if the spectrometer slit crosses the loop at two or more locations (or if the loop is scanned sufficiently rapidly) and if the studied loop is well isolated from the ambient loops.  (3) The re-examination of the swaying-like motions of the loop top observed in a 2D arcade loop simulated by \\citet{sel06} shows several features such as the period ratio and weak damping relative to the fundamental mode that are consistent with the theoretically expected ones for the vertical second harmonic oscillation. They could be evidence for the swaying-like motions resulting from the second-harmonic vertical mode in the loop. However, there are differences that suggest that the fundamental vertical mode was also excited. Simultaneous excitation of multiple modes complicates the identification further. The simple analysis presented here may not be adequate in such cases.  In Sect.~\\ref{dplsct} we only considered the shape of velocity profile along a loop, but the absolute amplitude of Doppler velocity may also be used to distinguish between different kink modes. The amplitude of Doppler velocity can be inferred from the apparent displacement amplitude divided by the measured oscillation period for a given geometry and oscillation mode. The amplitude is expected to be a strong function of the oscillation mode and could be used to identify the mode by comparing with direct measurements of Doppler velocity obtained with a slit spectrograph at only a single location.  We have determined which type of kink modes each of the 14 transverse loop oscillations observed by TRACE belongs to based on a comparison between the observed and modeled signatures (see the last second column in Table~\\ref{partab}). For 2 loops out of the 6 loops in Category I the oscillation corresponds to the fundamental vertical mode  (but may be combined with the fundamental horizontal mode), and for another it is identified as the fundamental horizontal kink mode. The remaining 3 could not be uniquely identified.  Out of the 6 loops in Category II one loop shows the interesting feature which first indicates the dominant vertical motion and then agrees with the combined fundamental vertical and horizontal modes. The oscillations of the other loops could not be uniquely identified. The two loops in Category III are clearly identified as displaying the fundamental horizontal kink mode. The results suggest that horizontal oscillations seem to be more easily excited under solar condition. This may be because the disturbing sources (e.g., flares) are commonly located near to but rarely immediately in the plane of the oscillating loops. Note, however, that any vertical oscillations in loops of Category III would not easily have been detected in TRACE data.  In summary, our analysis has shown that for a given viewing geometry it may be difficult to distinguish between some of the different kink modes and in future when identifying kink modes the results of this work should be taken into account, and comparisons should be made to simple geometrical models, like those considered here. Ideally, the application of this method requires the complete loop to be well visible. This is not always the case with observed oscillating loops, which are sometimes only partially visible.  The current analysis is not complete. Thus, it is restricted to loops lying in a plane and does not consider the case of sheared loops such as that reported by \\citet{dem07}. We have also not taken into account the information in the oscillation period, although it is affected by different atmospheric parameters and may not always help distinguish between modes. Another way of distinguishing between vertical and horizontal modes may be through brightness fluctuations \\citep{wan04b}, but before doing that improved statistics of vertical oscillations are needed. Finally, stereoscopy, when corresponding data are available, should be very helpful for obtaining the true geometry of oscillating loops and of the oscillations they harbor \\citep[e.g.][]{fen07}."
    },
    "0808/0808.0016_arXiv.txt": {
        "abstract": "The Efstathiou, Lake and Negroponte (1982) criterion can not distinguish bar stable from bar unstable discs and thus should not be used in semi-analytic galaxy formation simulations. I discuss the reasons for this, illustrate it with examples and point out shortcomings in the recipes used for spheroid formation. I propose an alternative, although much less straightforward, possibility. ",
        "introduction": "\\indent  Technically, it is not yet possible to make full cosmological simulations which include both the dark matter and the baryons, down to the scale of individual galaxies. Yet the dark matter only simulations have reached very high resolutions (Springel et al. 2008, priv. comm.) and it is crucial to find ways of exploiting them fully. A scheme involving resampling has been devised, in which a given object and its surrounding region are singled out and resimulated at high resolution (e.g. Katz \\& White 1993). Nevertheless, this technique can only give information on one, or at best a few objects. For this reason, semi-analytical models were introduced, which can give information on properties of galaxy populations. These models are tagged on to the dark matter only simulations and use `recipes' to describe the evolution of the baryons. It is clear that the results will be useful only if the recipes in question are correct and adequately chosen. This is not an easy task, since these recipes need to include in a relatively simple way a fair fraction of the available information on many key astrophysical processes.  In their quest to have more spheroids, either bulges or ellipticals, some semi-analytic modellers are now considering  disc instabilities. They use the Efstathiou, Lake and Negroponte (1982, hereafter ELN)) criterion to distinguish between bar stable and bar unstable discs. Using this criterion, this information can be obtained very simply from the maximum rotational velocity and the disc mass and scale-length. Once a disc is found to be bar unstable, it is turned instantaneously in the model into an elliptical (e.g. Bower et al. 2006), or, alternatively, half of its mass is turned into a bulge and this repeatedly until disc stability according to the ELN criterion is achieved (e.g. De Lucia \\& Helmi 2008). In this way a considerable amount of mass is turned into a spheroid and several problems concerning the K-band luminosity function are alleviated (see Parry et al. 2008, for a discussion of the relative merit of the two approaches). There are, however, a number of shortcomings in this modelling approach. ",
        "conclusions": "\\label{sec:summary} \\indent  In the previous section I showed that the ELN criterion is too simplistic to be able to describe a complex phenomenon such as bar formation and to distinguish bar stable from bar unstable discs. It thus cannot be used in semi-analytic calculations.  Further problems concern the subsequent evolution, after the bar has formed. No simulation has ever shown an unstable disc turn into an elliptical, or acquire a very massive classical bulge instantaneously. The bar can form a boxy/peanut bulge (see Figs~\\ref{fig:unstable} and \\ref{fig:stable}) or a discy bulge, and the formation of a classical bulge from either of those is not excluded\\footnote{See Athanassoula (2005) for the definition and properties of the different types of bulges.}. This classical bulge, however, would be much smaller than required by the semi-analytic models, and, furthermore, it would not form instantaneously since the bar needs some time to form, more time is necessary for the boxy/peanut, or discy bulge formation and yet more time is necessary for the putative conversion into a classical bulge. Thus the second part of the De Lucia \\& Helmi (2008) model has shortcomings, but these may turn out to be quantitative rather than qualitative. On the contrary, the Bower et al. (2006) model has shortcomings at the qualitative level, since the disc can not disappear, i.e. an elliptical could not form from a bar unstable disc without a merger. Both models of course have already problems in the first part, i.e. deciding whether a disc is bar unstable or not, since they both use the ELN criterion.  Can we replace the simple ELN recipe for handling disc instabilities by a better one? This task is not easy, since we need to include in this recipe information on disc stability, bar formation and evolution and the subsequent formation of the different types of bulges, all of which are complex, nonlinear processes. I firmly believe that it is not possible to find one, or a few simple formulas, like the ELN criterion, that can describe all the necessary ingredients of these processes. It might thus be preferable to seek a solution intermediate between the full cosmological simulation including both the dark matter and the baryons at sufficient resolution, which will not be available in the foreseeable future, and the equally difficult task of finding appropriate recipes.  Computer simulations allow us today to routinely make $N$-body simulations which can describe the evolution of a single galaxy. In fact a very large number are already available. Using their results instead of the recipes is a possibility well worth exploring. For example, for the problem at hand one would need a small library of $N$-body simulations with and without gas, describing bar formation and evolution. These simulations should cover, albeit very crudely, the necessary parameter space describing the halo, stellar disc and gas components. This would include not only the mass and scale-length ratios of the various components, but also the different amounts of random motion in the stellar disc and in the spheroids. From these simulations one should extract all the necessary parameters, such as the times necessary for bar and peanut formation, the amount of mass in the boxy/peanut bulge or the discy bulge, etc. This information would be assembled in some sort of table. Then at each time step of the semi-analytic simulation, instead of checking the criterion or applying the recipe, one would extract from the above table the relevant information as a function of the properties of the galaxy at the time under consideration.  The above scheme is of course not a substitute for full scale cosmological simulations, starting {\\it ab initio}. It also has a number of shortcomings, the most important of which is that our knowledge about the velocity dispersions in disc galaxies is very limited. Nevertheless, it is a very important improvement with respect to the presently used scheme."
    },
    "0808/0808.2539_arXiv.txt": {
        "abstract": "More reliable constraints on the microlensing optical depth comes from a better understanding of the Galactic model. Based on well-constrained Galactic bulge and disk models constructed from survey observations, such as, $\\it{HST}$, 2MASS, and SDSS, we calculate the microlensing optical depths toward the Galactic bulge fields, and compare them with recent results of microlensing surveys. We test $\\chi^2$ statistics of microlensing optical depths expected from those models, as well as previously proposed models, using two types of data: optical depth map in $(l, b)$ and averaged optical depth over the Galactic longitude $l$ as a function of the latitude $b$. From this analysis, we find that the Galactic bulge models of 2MASS, Han $\\&$ Gould (2003), and G2 of Stanek et al. (1997) show a good agreement with the microlensing optical depth profiles for all the microlensing observations, compared with E2 of Stanek et al. (1997). We find, on the other hand, that models involving an SDSS disk model produce relatively higher $\\chi^2$ values. It should be noted that modeled microlensing optical depths diverge in the low Galactic latitude, $|b| \\la 2 \\arcdeg$. Therefore, we suggest the microlensing observation toward much closer to central regions of the Galaxy to further test the proposed Galactic models,  if it is more technically feasible than waiting for large data set of microlensing events. ",
        "introduction": "Microlensing survey was originally proposed  as a tool for detecting massive astronomical compact halo objects (MACHOs) in the Galactic halo \\citep{pac86}. Searches for the microlensing events towards the Magellanic Clouds by the MACHO \\citep{alc93} and EROS groups \\citep{aub93} have placed constraints on the  fraction of the Galactic MACHO populations. Searches for dark objects in the halo of M31 have also been performed (AGAPE, Ansari et al. 1999; WeCAPP, Riffeser et al. 2003; MEGA, de Jong et al. 2004; VATT-Columbia, Uglesich et al. 2004; POINT-AGAPE, Calchi Novati et al. 2005; Angstrom, Kerins et al. 2006). In addition, microlensing has proven itself as a powerful tool in constraining the distribution of faint stellar objects in the Milky Way \\citep{kir94,han95,han03,woo05,woo07,cal08}. For instance, \\citet{han95} examined theoretical Galactic models by constructing a microlensing map. They explored a triaxial bar-shaped bulge model, which was based on by $\\it{COBE}$-DIRBE multiwavelength observations of the Galactic bulge \\citep{wei94}, and demonstrated that observationally determining the microlensing optical depth may provide supplementary information on the Galactic mass distribution.  Unfortunately, however, reality was not so simple. Early measurements of the microlensing optical depth towards the Galactic bulge were significantly higher than predictions, suggesting both observational and theoretical sides make efforts to reconcile. Measurements at Baade's window by OGLE from 9 microlensing events \\citep{uda94} and at $(l,b)=(2\\fdg55, -3\\fdg64)$ by MACHO from 13 events of clump giant sources \\citep{alc97a} have yielded $\\tau=3.3\\pm1.2\\times10^{-6}$ and $\\tau=3.9^{+1.8}_{-1.2}\\times10^{-6}$, respectively. \\citet{alc97a} pointed out a possibility of a systematic bias such as blending effect in the microlensing optical depth measurement to explain those high optical depth values in the sense that bulge fields are high dense regions. \\citet{pop01} suggested that the bias due to the blending can be avoided by using the events only with bright source stars such as red clump giant stars. Such stars are identified by their well-defined position in the color-magnitude diagram. This position ensures that they are most likely in the Galactic Bulge. Their large flux also makes blending problems relatively unimportant. Measurements of $\\tau$ using red clump giants have tended to reduce discrepancies with models. Based on 16 events, the EROS-2 collaboration \\citep{afo03} gave the microlensing optical depth of $0.94\\pm0.29\\times 10^{-6}$ at $ (l,b) =(2\\fdg5 ,-4\\fdg0)$. The MACHO group calculated the microlensing optical depth toward the Galactic bulge using their seven-year survey data \\citep{pop05}. The measurement of $\\tau=2.17^{+0.47}_{-0.38}\\times10^{-6}$ at $(l,b)=(1\\fdg50,-2\\fdg68)$ was obtained from 42 microlensing events detected during the monitoring of about $7.4\\times10^5$ clump giant sources covering 4.5 deg$^2$. More recently, OGLE-II presented the measurement of the microlensing optical depth toward the Galactic bulge based on recent four-year survey data \\citep{sum06}. Using the sample of 32 microlensing events in 20 bulge fields covering $\\sim5$ deg$^2$, they found $\\tau=2.55^{+0.57}_{-0.46}\\times10^{-6}$ at $(l,b)=(1\\fdg16,-2\\fdg75)$. Efforts in a theoretical side also refined the microlensing optical depth estimate by using more sophisticated models based on observational results from wide field surveys. Theoretical estimates involving a bar in the bulge oriented along the line of sight approach to a value in the range of $0.8\\times10^{-6}$ to $2.0\\times10^{-6}$ \\citep[see e.g.][]{pac94,zha95,zha96a,bin00,han03}.  The optical depth in microlensing is defined as the probability that a source star is being placed within the Einstein radius of a foreground lens star. If bulge stars are distributed over a distance $d$ from the Earth, the microlensing optical depth is given by \\begin{equation} \\tau=\\frac{4\\pi G}{c^2}\\frac{\\int^{d}_0 \\left[\\int^{D_s}_0dD_l\\;\\rho(D_l)\\;D\\right]dD_s\\;D_s^2\\;n(D_s)} {\\int^{d}_0dD_s\\;D_s^2\\;n(D_s)}, \\end{equation} where $n(D_s)$ is the number density of source stars, $\\rho(D_l)$ is the mass density of lenses along the line of sight, $D_l$, $D_s$ and $D_{ls}$ are the distances from the observer to the lens and source, and the distance from the lens to the source, respectively, and $D$ is $D_lD_{ls}/D_s$.  Given that the microlensing optical depth is subject to an underlying mass distribution model, it is crucial to employ a mass model, which can be constructed from independent observations, e.g., star counts. In fact, the stellar contents of the disk and the bulge have been measured by survey projects and modeled to estimate the optical depth. For instance, star counts from {\\it Hubble Space Telescope} ({\\it HST}) provided constraints on faint stars down to the hydrogen burning limit in the disk and 0.15 $M_{\\odot}$ in the bulge \\citep{hol98,zoc00,zhe01}. Using those observational results, one may put relatively stronger constraints on the mass distribution in terms of the microlensing optical depth \\citep[e.g.][]{han03}. Results from large-scale sky surveys such as 2MASS \\citep{skr97} and SDSS \\citep{yor00} are also available in the near-infrared and optical wavelength bands. On the basis of those results, new models of the Galactic bulge and disk have been suggested \\citep{lop05,jur08}. In the present paper, we calculate the microlensing optical depth with Galactic bulge and disk models recently by various models based on recent survey observations as well as with those studied earlier. We then compare the optical depths based on models with the results of MACHO, OGLE-II, and EROS-2 surveys.  The paper is organized as follows. We begin with a brief description of models of the Galactic bulge and the disk in $\\S$ 2. We present expected microlensing optical depths from various models and compare those with observed optical depths in $\\S$ 3. Finally, we conclude with a summary of our results and discussion. ",
        "conclusions": "By comparing the microlensing optical depth map constructed by a mass model with the observed microlensing optical depth one may put a tight constraint on the Galactic model. We have estimated the microlensing optical depths toward the Galactic bulge using the Galactic bulge and disk models constructed from survey observations, and compared them with the sample of microlensing events observed towards the Galactic bulge with red clump giant sources reported by the MACHO, OGLE-II and EROS-2 collaborations. From the analysis we carried out, we show that model estimates are in general compatible with microlensing observations. According to calculated $\\chi^2$ we find that 2MASS, HG03, and G2 models reproduce the microlensing optical depth distribution from observations well. We also find that for those well-fit models contribution to the microlensing optical depth of the disk components is not crucial. For example, for 2MASS cases two disk models do not make any difference in $\\chi^2$. It is, however, interesting to note that the analyses of SDSS disk models give us somewhat different results. That is, models including SDSS disk model produce relatively higher $\\chi^2$ values than others do.  It should be pointed out that current observational data of both the microlensing optical depth in two dimensional fields and the averaged optical depth over $l$ are yet statistically insufficient to rule out conclusively any specific models we have considered since the number of observed microlensing events from monitoring RCGs in recent microlensing surveys is still small. On the other hands, we note that the difference of the average optical depth among models increases as closer to the Galactic center. Therefore, we suggest the microlensing observation toward closer to the Galactic center regions if possible may be more effective in constraining the Galactic model than extending microlensing searches to detect more microlensing events."
    },
    "0808/0808.4129_arXiv.txt": {
        "abstract": "An accurate knowledge of the mineralogy (chemical composition and crystal structure) of the silicate dust in the interstellar medium (ISM) is crucial for understanding its origin in evolved stars, the physical and chemical processing in the ISM, and its subsequent incorporation into protostellar nebulae, protoplanetary disks and cometary nuclei where it is subjected to further processing. While an appreciable fraction of silicate dust in evolved stars, in protoplanetary disks around pre-main sequence stars, in debris disks around main sequence stars, and in cometary nuclei is found to be in crystalline form, very recent infrared spectroscopic studies of the dust along the sightline toward the Galactic center source $\\sgrA$ placed an upper limit of $\\simali$1.1\\% on the silicate crystalline fraction, well below the previous estimates of $\\simali$5\\% or $\\simali$60\\% derived from the observed 10$\\mum$ absorption profile for the local ISM toward Cyg OB2 No.12. Since the sightline toward $\\sgrA$ contains molecular cloud materials as revealed by the detection of the 3.1 and 6.0$\\mum$ water ice absorption features, we argue that by taking into account the presence of ice mantles on silicate cores, the upper limit on the degree of silicate crystallinity in the ISM is increased to $\\simali$3--5\\%. ",
        "introduction": "The mineralogical composition of dust contains important information about its origin and evolution, and may reveal the physical, chemical and evolutionary properties of the astrophysical regions where the dust is found. Silicate, one of the major components of cosmic dust species, is ubiquitously seen in various astrophysical environments, ranging from the Galactic diffuse interstellar medium (ISM), HII regions, and the dust torus around active galactic nuclei to dust envelopes around evolved stars, protoplanetary disks around pre-main sequence stars, debris disks around main sequence stars, cometary comae and interplanetary space. Their chemical composition and crystal structure vary with local environments.  Interstellar spectroscopy provides the most diagnostic information on dust composition. In the infrared (IR), the strongest interstellar spectral features are the 9.7 and 18$\\mum$ absorption (or emission) bands which are generally attributed to the Si--O stretching and O--Si--O bending modes, respectively. While the absorption profiles measured in laboratory for crystalline olivine and pyroxene show many sharp features, in the diffuse ISM, the observed 9.7 and 18$\\mum$ silicate features are broad and relatively featureless, suggesting that interstellar silicates are mainly amorphous rather than crystalline. Since the crystallinity of silicate dust is intimately connected to the energetic processing occurred in the evolutionary life cycle of dust and depends on the environmental properties, of particular interest to silicate astromineralogy is the abundance of crystalline silicates in different astrophysical environments, especially in the diffuse ISM.  Recently, a number of studies have been made to determine the crystallinity degree of silicates in the diffuse ISM. From the observed 9.7$\\mum$ absorption profile for dust in the local ISM toward Cyg OB2 No.12, Li \\& Draine (2001) estimated the fraction of Si in crystalline silicates to be $\\simlt$5\\%. Demyk et al.\\ (1999) derived an upper limit of $\\simali$1--2\\% of crystalline silicates in mass toward two massive protostars. However, Bowey \\& Adamson (2002) argued that a complex mixture of crystalline silicates (60\\% by mass) and amorphous silicates (40\\% by mass) could explain the observed smooth silicate absorption profile at 9.7$\\mum$ toward Cyg OB2 No.12. More recently, Kemper, Vriend \\& Tielens (2004) placed a more strict constraint of $\\simlt 1.1\\%$ on the crystallinity of interstellar silicates, based on a detailed analysis of the 9.7$\\mum$ feature obtained with the Short Wavelength Spectrometer (SWS) on board the Infrared Space Observatory (ISO). This upper limit $\\simali$1.1\\% of crystalline silicates was derived from a direct comparison of the $\\sgrA$ spectrum with theoretical spectra for pure silicates.  However, it is evident from the detection of the 3.1 and 6.0$\\mum$ water ice features respectively attributed to the O--H stretching and bending modes that there are molecular cloud materials along the line of sight toward $\\sgrA$ (McFadzean et al.\\ 1989; Tielens et al.\\ 1996; Chiar et al.\\ 2000). This is further confirmed by the detection of solid CO$_2$ absorption toward $\\sgrA$ (Lutz et al.\\ 1996; de Graauw et al.\\ 1996). In order to infer the precise mineralogical composition of dust, we have to take the grain ice mantles into account when interpreting the observed silicate features. Indeed, in this paper we show that ignoring the ice mantles coated on the silicate cores can result in an underestimation of the crystalline degree of silicate dust, as much as $\\simali$3--5\\%.  The purpose of this work is to address the question how the ice mantles affect the determination of dust mineralogical composition (i.e. whether and to what degree the inclusion of ice mantles will hide the sharp features of crystalline silicates and lead to an underestimation of the silicate crystallinity), not to model any specific astronomical objects. Therefore, neither the specific choice of silicate dielectric functions nor the precise constituents of the ice mantles would affect our conclusion. ",
        "conclusions": "We have investigated the effects of ice mantles coated on silicate cores of various shapes on the determination of the crystallinity degree of silicates. It is found that for the dust in the line of sight toward the Galactic center source $\\sgrA$, $\\simali$3\\% crystalline silicates could be hidden by ice mantles, well exceeding the upper limit $\\simali$1.1\\% of Kemper et al.\\ (2004) derived from the assumption of the absence of dense molecular materials in this line of sight. We take the standard approach by assuming that the silicates are amorphous and then adding in crystalline components. But as shown in Kemper et al.\\ (2004), the smooth, featureless 10$\\mum$ amorphous silicate spectrum allows as much as $\\simali$2.2\\% crystalline silicates (or even higher; see Bowey \\& Adamson 2002). Therefore, the total allowable degree of crystallinity would be $\\simali$5\\%, consistent with the earlier estimates of Li \\& Draine (2001)."
    },
    "0808/0808.3153_arXiv.txt": {
        "abstract": "Using a large sample of optical spectra of late-type dwarfs, we identify a subset of late-M through L field dwarfs that, because of the presence of low-gravity features in their spectra, are believed to be unusually young. From a combined sample of 303 field L dwarfs, we find observationally that 7.6$\\pm$1.6\\% are younger than 100 Myr. This percentage is in agreement with theoretical predictions once observing biases are taken into account. We find that these young L dwarfs tend to fall in the southern hemisphere (Dec$<$0$^\\circ$) and may be previously unrecognized, low-mass members of nearby, young associations like Tucana-Horologium, TW Hydrae, $\\beta$ Pictoris, and AB Doradus. We use a homogeneously observed sample of roughly one hundred and fifty 6300-10000 \\AA\\ spectra of L and T dwarfs taken with the Low-Resolution Imaging Spectrometer at the W.\\ M.\\ Keck Observatory to examine the strength of the 6708-\\AA\\ \\ion{Li}{1} line as a function of spectral type and further corroborate the trends noted by \\cite{kirkpatrick2000}. We use our low-gravity spectra to investigate the strength of the \\ion{Li}{1} line as a function of age. The data weakly suggest that for early- to mid-L dwarfs the line strength reaches a maximum for a few$\\times$100 Myr, whereas for much older (few Gyr) and much younger ($<$100 Myr) L dwarfs the line is weaker or undetectable. We show that a weakening of lithium at lower gravities is predicted by model atmosphere calculations, an effect partially corroborated by existing observational data. Larger samples containing L dwarfs of well determined ages are needed to further test this empirically. If verified, this result would reinforce the caveat first cited in \\cite{kirkpatrick2006} that the lithium test should be used with caution when attempting to confirm the substellar nature of the youngest brown dwarfs. ",
        "introduction": "Determining the difference between an object stably burning hydrogen (a main sequence star) and an object of lower mass incapable of such stable fusion (a brown dwarf) is not always straightforward. Although objects with main sequence spectral types of O, B, A, F, G, or K are always stars and objects with a spectral type of T are always brown dwarfs, dwarfs of type M or L comprise a mixture of stellar and substellar objects.  By definition, distinguishing a brown dwarf from a star requires knowledge that the object is not burning hydrogen in its core. There is no direct way to probe the interiors of these objects, so other techniques are required. The most common indirect method used is the ``lithium test'' (\\citealt{rebolo1992}). This test recognizes the fact that the temperature ($T > 3{\\times}10^6$K; \\citealt{burrows1997, nelson1993}) needed to sustain core hydrogen fusion in the lowest mass stars is only slightly higher than that needed to burn lithium ($T \\approx 2{\\times}10^6$K; \\citealt{pozio1991}). If lithium does not burn, that means hydrogen is also not being fused. Fortuitously lithium, once destroyed, is not easily manufactured again in stellar interiors, so stars and brown dwarfs will never have more than their natal abundance of this element. In principle, the presence or absence of lithium in the spectrum can thus provide knowledge of inner temperatures.  In practice there are several drawbacks when applying the lithium test to M and L dwarfs:  (1) The ground-state resonance line of \\ion{Li}{1} is located at 6708 \\AA, a portion of the spectrum where the flux is quite low for late-M and L dwarfs. The low flux levels means that large telescopes are needed for observation of the line. Even then, very few late-type dwarfs are bright enough that they can be acquired at high resolution so to obtain a sizable sample of spectra, moderate resolution ($\\Delta{\\lambda}\\approx 5-10$ \\AA) is required. Fortunately, this is sufficient for measuring \\ion{Li}{1} equivalent widths of a few \\AA\\ or greater provided that the signal-to-noise ratio is high.  (2) As stated above, there is a ${\\sim}1{\\times}10^6K$ difference between the temperatures for lithium ignition and normal hydrogen ignition. In other words, some brown dwarfs are capable of burning lithium and for those the lithium test will fail. These higher mass brown dwarfs are unrecognizable as substellar based on this test alone.  (3) A sufficient amount of time needs to elapse for the natal supply of lithium to be extinguished. The \\ion{Li}{1} line may therefore still be present not because the object is substellar, but because it is a star too young to have fused its lithium store. In a very young, low-mass M star, lithium may still be present in the spectrum either because the core has not yet reached lithium-burning temperatures or because there has been insufficient time for convective recirculation to have destroyed all the lithium in the star and in particular all of the lithium in the photosphere (but this recirculation is expected to be rapid compared to the other two timescales considered here; \\citealt{basri1998}). Despite these potential pitfalls, \\cite{basri1998} has demonstrated that any object with lithium absorption and Teff $<$ 2670K is unambiguously a brown dwarf. Thus if we confine our focus to objects with temperatures cooler than this, the lithium test can be applied regardless of these considerations. Using the relation of optical dwarf spectral type to temperature from \\cite{kirkpatrick2008} -- $T_{eff}(K) = 3759 - 135x$, where, e.g, x=0 for M0, 9 for M9, 10 for L0, and 18 for L8 -- we find that an effective temperature of 2670K corresponds to a dwarf of optical type M8. For the remainder of this paper, we will therefore restrict our discussion of the lithium test to dwarfs with optical types of M8 and later.  (4) It is expected that the 6708-\\AA\\ line will disappear at very cool temperatures as lithium forms molecular species. \\cite{lodders1999} has shown that below $T_{eff} \\approx 1500$K Li-bearing molecules like LiCl and LiOH begin to form. For brown dwarfs of solar age, this temperature corresponds to optical spectral types of late-L through mid-T (see Figure 7 of \\citealt{kirkpatrick2005}). Stars of solar abundance do not exist at temperatures this low (see Figure 10 of \\citealt{kirkpatrick2005}), so we can safely assume that objects in this range of spectral types are all brown dwarfs. For these the lithium test is not needed. This theoretical prediction of lithium disappearance at these types should nonetheless be tested with empirical data.  (5) Young ($<$100 Myr) brown dwarfs have radii larger than older stars and brown dwarfs of the same spectral type because they are still contracting to their final structural configuration (\\citealt{burrows1997}). These young brown dwarfs are also lower in mass than stars and older brown dwarfs of the same spectral type. Both the larger radius and the lower mass mean that these objects will have lower surface gravities, which should be detectable even with low-resolution spectroscopy. (See \\S4.1.) Lower gravity will weaken the \\ion{Li}{1} line due to the lower atmospheric pressure, making it much harder to detect. This has been observed by \\cite{kirkpatrick2006} in the low-gravity, early-L dwarf 2MASS J01415823$-$4633574, which shows no \\ion{Li}{1} line down to an equivalent width of 1 \\AA. Weak \\ion{Li}{1} is seen despite the fact that this object is believed to be a young (1-50 Myr), low-mass (6-25 M$_{Jupiter}$) brown dwarf and despite the fact that other brown dwarfs of similar spectral type are known that show prominent \\ion{Li}{1} lines -- such as the L0 dwarf 2MASS J11544223$-$3400390 with a \\ion{Li}{1} equivalent width of 3 \\AA.  We can test points (4) and (5) using a unique set of spectroscopic data. Over the past several years we have amassed a library of optical spectra of L and T dwarfs primarily using the W.\\ M.\\ Keck Observatory. Data from \\cite{kirkpatrick1999, kirkpatrick2000, burgasser2000, kirkpatrick2001, wilson2001, burgasser2003, burgasser2003b, thorstensen2003, mcgovern2004, kirkpatrick2006} combined with new data presented in this paper now gives us a library of $\\sim$150 spectra spanning the wavelength regime 6300-10000 \\AA\\ and covering spectral types L0 through T8. With such a large set of homogeneous spectra, we can better characterize the behavior of the \\ion{Li}{1} line as a function of spectral type. Moreover, this large set of spectra enables us to identify peculiar objects, particularly those with unusual spectra indicative of low gravity. As stated above, low gravity in L and T dwarfs is an indicator of youth, so these spectra also enable us to characterize the \\ion{Li}{1} strength as a function of age.  In this paper we present the newest spectra in our optical spectral library (\\S2), analyze objects with detected H$\\alpha$ (\\S3), and identify objects with low-gravity (young) signatures (\\S4). We perform further analysis of the young objects (\\S5) and build statistics on the behavior of \\ion{Li}{1} line strengths as a function of both spectral type and age (\\S6). Conclusions are given in \\S7. ",
        "conclusions": "Using a large set of optical specta, we have shown that a subset of objects with peculiar spectra can be identified as low-gravity brown dwarfs. Low gravity is a hallmark of youth, as these objects have lower masses and more extended atmospheres than higher mass, older field dwarfs. We find that the observed percentage of young ($<$100 Myr old) L dwarfs in our sample -- 7.6$\\pm$1.6\\% -- is consistent with theoretical predictions once selection biases are considered. Even though low gravity enables these young objects to be identified readily via spectroscopy, it complicates the brown dwarf ``lithium test'' because it weakens the observable \\ion{Li}{1} line for a given lithium abundance and spectral type. Our sample of young brown dwarfs is found to lie primarily in the southern hemisphere, which (probably not coincidentally) is the location of known, nearby young associations. Further work on these young objects is needed to establish or refute membership in these groups."
    },
    "0808/0808.0546_arXiv.txt": {
        "abstract": "We present results of multi-epoch VLBI observations with VERA (VLBI Exploration of Radio Astrometry) of the 22~GHz H$_{2}$O masers associated with a young stellar object (YSO) IRAS~22198+6336 in a dark cloud L1204G. Based on the phase-referencing VLBI astrometry, we derive an annual parallax of IRAS~22198+6336 to be 1.309$\\pm$0.047~mas, corresponding to the distance of 764$\\pm$27~pc from the Sun. Although the most principal error source of our astrometry is attributed to the internal structure of the maser spots, we successfully reduce the errors in the derived annual parallax by employing the position measurements for all of the 26 detected maser spots. Based on this result, we reanalyze the spectral energy distribution (SED) of IRAS~22198+6336 and find that the bolometric luminosity and total mass of IRAS~22198+6336 are 450$L_{\\odot}$ and 7$M_{\\odot}$, respectively. These values are consistent with an intermediate-mass YSO deeply embedded in the dense dust core, which has been proposed to be an intermediate-mass counterpart of a low-mass Class~0 source. In addition, we obtain absolute proper motions of the H$_{2}$O masers for the most blue-shifted components. We propose that the collimated jets aligned along the east-west direction are the most plausible explanation for the origin of the detected maser features. ",
        "introduction": "One of the crucial issues in astronomy and astrophysics is to understand formation processes of stars, planetary systems, and their clusters. They are the most basic constituents in the Galaxy and hence, play essential roles in its formation and evolution. In order to understand physics and dynamics in the star-formation processes, it is necessary to obtain accurate physical properties of target sources such as size, mass, and luminosity. Because all of these values strongly depend on the adopted distance, the distance is the most fundamental parameter for quantitative discussion.  In spite of its importance, however, it has long been difficult to measure accurate distances to star-forming regions without assumptions. The most reliable and direct way to determine the distance is an annual trigonometric parallax method, based on precise measurements of the position and motion of the object. In 1990's, the Hipparcos satellite measured annual parallaxes for more than 100~000 stars with a typical precision of 1~mas level \\citep{perryman1997}. Nevertheless, accurate distance measurements with Hipparcos for star-forming regions were limited to nearby sources within a few hundred pc from the Sun (e.g. Lombardi et al. 2008). Although distances to several OB clusters associated with nearby star-forming regions up to $\\sim$600~pc have been reported \\citep{dezeeuw1999}, Hipparcos could not directly determine the distances to newly born stars deeply embedded in dusty molecular cloud cores, because Hipparcos was capable of observing only in the optical wavelengths.  In the last few years, phase-referencing VLBI astrometry has been developed drastically, yielding the annual parallaxes of nearby star-forming regions. For instance, radio continuum sources in the Taurus \\citep{loinard2005, loinard2007, torres2007}, Ophiuchus \\citep{loinard2008}, and Orion \\citep{sandstrom2007, menten2007} regions were observed with the NRAO Very Long Baseline Array (VLBA) to derive their annual parallaxes with typical uncertainty of better than a few percent.  In addition, we have started the initial scientific project \"Measurements of annual parallaxes of nearby molecular clouds\" with VERA (VLBI Exploration of Radio Astrometry). VERA is a Japanese VLBI network dedicated to astrometric observations aimed at revealing 3-dimensional structure of the Galaxy \\citep{honma2007, sato2007, sato2008}. In our project for nearby molecular clouds, we specially focus on the accurate distance measurements of star-forming regions within 1~kpc from the Sun \\citep{dame1987}. Our results will contribute to refine all of the observational studies on nearby star-forming regions by providing their most accurate and directly measured distances. We have already conducted astrometric observations of the H$_{2}$O maser sources in the Orion \\citep{hirota2007}, Ophiuchus \\citep{imai2007}, and Perseus \\citep{hirota2008} regions. All of the results provide direct measurements of their distances with much higher accuracy than the previous photometric methods and the parallax measurements with Hipparcos.  As a part of this project, we present results of astrometric observations of the H$_{2}$O maser source IRAS~22198+6336 in a nearby dark cloud core L1204G\\footnotemark. \\footnotetext{In some literature, IRAS~22198+6336 is called as L1204A \\citep{fukui1989, palla1991, migenes1999}. However, it is quite confusing because L1204A is identified to be another core associated with the well known massive young stellar object (YSO) S140~IRS1 \\citep{tafalla1993, tofani1995}. Thus, we refer to the host of IRAS~22198+6336 as L1204G in this paper according to \\citet{tafalla1993}.} L1204 is located in the Cepheus-Cassiopeia molecular cloud complex \\citep{yonekura1997} and is associated with a well known high-mass star-forming region S140. It is also identified to be a molecular cloud WB176 \\citep{wouterloot1989}. According to the molecular line observations of L1204, it consists of several high density molecular cores labeled by A-G \\citep{tafalla1993}. A deeply embedded YSO IRAS~22198+6336 is associated with one of the dense cores, L1204G, and it is identified to be a powering source of molecular outflow \\citep{fukui1989, wilking1989}. The centimeter and millimeter counterpart of IRAS~22198+6336 is detected with the VLA \\citep{sanchez2008}. The H$_{2}$O masers were detected \\citep{palla1991} and their distributions were investigated by the VLA \\citep{tofani1995} and VLBA \\citep{migenes1999}. However, they did not measure the proper motions of the H$_{2}$O masers because observations by \\citet{tofani1995} could not achieve sufficient high resolution and \\citet{migenes1999} observed only at 1 epoch. Therefore, the kinematics of the circumstellar materials around IRAS~22198+6336 remains still unclear.  The distance to L1204 is estimated by a photometric method to be 910~pc \\citep{crampton1974} assuming that a B0V star HD211880 is the ionization source of the H{\\sc{ii}} region S140, which is also associated with L1204. On the other hand, \\citet{dezeeuw1999} identified HD211880 (HIP110125) as a member of the Cepheus~OB2 cluster. Using the Hipparcos data, \\citet{dezeeuw1999} determined the distance to the Cepheus~OB2 cluster to be closer value of 615$\\pm$35~pc from the Sun. The kinematic distance of 1.3~kpc is sometimes adopted in several literatures \\citep{wouterloot1989, molinari1996} according to the systemic velocity of $-11$~km~s$^{-1}$ \\citep{tafalla1993}. Depending on the adopted distance, the luminosity of IRAS~22198+6336 could differ by a factor of 4, and hence, its physical properties sometimes contain large uncertainties.  Based on our accurate astrometric observations with VERA, we will provide the parallactic distance to IRAS~22198+6336 in L1204G, together with the kinematics of the circumstellar gas traced by the H$_{2}$O masers. We will also discuss the basic properties of the powering source of the H$_{2}$O masers by adopting our new distance to L1204G. ",
        "conclusions": "In this paper, we presented the results of multi-epoch VLBI astrometry with VERA of the 22~GHz H$_{2}$O masers associated with IRAS~22198+6336 in L1204G. Using the absolute positions of total 26~maser spots that were detected in at least three observing sessions, we derived the annual parallax of IRAS~22198+6336 to be 1.309$\\pm$0.047~mas, corresponding to the distance of 764$\\pm$27~pc from the Sun. This is the most accurate distance to L1204G with the uncertainty of only 4\\%.  Because the newly derived distance is closer than those well adopted in the previous literatures (910~pc or 1.3~kpc), the bolometric luminosity of IRAS~22198+6336 should be significantly reduced. According to the SED of IRAS~22198+6336, the bolometric luminosity and total mass of IRAS~22198+6336 were refined to be 450$L_{\\odot}$ and 7$M_{\\odot}$, respectively. These values are not consistent with a massive YSO but are in agreement with an intermediate-mass YSO. We confirmed that IRAS~22198+6336 would be a deeply embedded intermediate-mass YSO analogous to a low-mass Class~0 source \\citep{sanchez2008}.  Together with the annual parallax, we measured the absolute proper motions of the H$_{2}$O masers. We discussed several kinematical models of the circumstellar gas such as a small scale expanding shell, collimated jets, or rotating disk. We proposed that the distributions, proper motions, and the previously reported velocity drifts \\citep{valdettaro2002, brand2003} of the maser features associated with IRAS~22198+6336 are consistent with the collimated bipolar jets and possibly rotating disk perpendicular to them.  IRAS~22198+6336 is one of the rare sources in the very early stage of intermediate-mass YSOs. Therefore, further high-resolution and high-sensitivity observations of the continuum emission from centimeter to submillimeter wavelength, in particular with the future interferometers such as EVLA and ALMA, and the proper motion measurements of the H$_{2}$O masers and continuum sources will be essential to reveal the kinematics and more detailed properties of IRAS~22198+6336, which will contribute to the understanding of very early stage of intermediate-mass star-formation processes.  \\vspace{12pt} We thank the anonymous referee for valuable comments and suggestions. We also thank Guillem Anglada and Carlos Carrasco-Gonzalez for providing us with information on the VLA observations of IRAS~22198+6336. We are grateful to the staff of all the VERA stations for their assistance in observations. TH is financially supported by Grant-in-Aids for Scientific Research from The Ministry of Education, Culture, Sports, Science and Technology (13640242, 16540224, and 20740112)."
    },
    "0808/0808.1630_arXiv.txt": {
        "abstract": "A new idea of deriving a cosmological term from an underlying theory has been proposed in order to explain the expansion history of the universe. We obtain the scale factor with this derived cosmological term and demonstrate that it reflects all the characteristics of the expanding universe in different era so as to result in a transition from inflation to late acceleration through intermediate decelerating phases by this single entity. We further discuss certain observational aspects of this paradigm. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.4015_arXiv.txt": {
        "abstract": "We present an updated spectroscopic orbit and a new visual orbit for the double-lined spectroscopic binary $\\sigma^2$ Coronae Borealis based on radial velocity measurements at the Oak Ridge Observatory in Harvard, Massachusetts and interferometric visibility measurements at the CHARA Array on Mount Wilson.  $\\sigma^2$~CrB is composed of two Sun-like stars of roughly equal mass in a circularized orbit with a period of 1.14 days.  The long baselines of the CHARA Array have allowed us to resolve the visual orbit for this pair, the shortest period binary yet resolved interferometrically, enabling us to determine component masses of 1.137 $\\pm$ 0.037 \\msun~and 1.090 $\\pm$ 0.036 \\msun.  We have also estimated absolute $V$-band magnitudes of $M_{\\rm V}{\\rm (primary)} = 4.35 \\pm 0.02$ and $M_{\\rm V}{\\rm (secondary)} = 4.74 \\pm 0.02$.  A comparison with stellar evolution models indicates a relatively young age of 1--3 Gyr, consistent with the high Li abundance measured previously.  This pair is the central component of a quintuple system, along with another similar-mass star, $\\sigma^1$~CrB, in a $\\sim$ 730-year visual orbit, and a distant M-dwarf binary, $\\sigma$~CrB~C, at a projected separation of $\\sim$ 10$\\arcmin$.  We also present differential proper motion evidence to show that components C~\\&~D (ADS 9979C~\\&~D) listed for this system in the Washington Double Star Catalog are optical alignments that are not gravitationally bound to the $\\sigma$~CrB system. ",
        "introduction": "\\label{sec:introduction}  $\\sigma$~CrB is a hierarchical multiple system 22 pc away.  Its primary components, $\\sigma^1$~CrB (HR 6064; HD 146362) and $\\sigma^2$~CrB (HR 6063; HD 146361), are in a visual orbit with a preliminary period of $\\sim$ 900 years \\citep{Sca1979}, of which the latter is an RS CVn binary with a circularized and synchronized orbit of 1.139-day period \\citep[][ SR03 hereafter]{Str2003}.  In addition to these three solar-type stars, the Washington Double Star Catalog\\footnote{$http://ad.usno.navy.mil/wds/$} (WDS) lists three additional components for this system.  WDS components C and D were resolved 18$\\arcsec$ away at 103\\sdeg~in 1984 \\citep{Pop1986} and 88$\\arcsec$ away at 82\\sdeg~in 1996 \\citep{Cou1996}, respectively.  We will show in \\S\\,\\ref{sec:Opt} that both these components are optical alignments that are not gravitationally bound to the $\\sigma$~CrB system.  Finally, WDS component E ($\\sigma$~CrB~C, HIP 79551) which was resolved 635$\\arcsec$ away at 241\\sdeg~in 1991 by \\textit{Hipparcos} \\citep{HIP1997}, was identified as a photocentric-motion binary by \\citet{Hei1990}.  The parallax and proper motion listed for this star in \\citet{van2007}, the improved \\textit{Hipparcos} results based on a new reduction of the raw data, match the corresponding measures for $\\sigma^2$~CrB within the errors, confirming a physical association.  SR03 presented photometric evidence in support of a rotation period of 1.157 $\\pm$ 0.002 days for both components of $\\sigma^2$~CrB, the central pair of this system.  They explained the 0.017-day difference between the rotation and orbital periods as differential surface rotation. \\citet{Bak1984} estimated an orbital inclination of 28\\sdeg, assuming component masses of 1.2 \\msun~based on spectral types.  SR03 subsequently adopted this inclination to obtain component masses of 1.108 $\\pm$ 0.004 \\msun~and 1.080 $\\pm$ 0.004 \\msun, but these masses are based on circular reasoning, and the errors are underestimated as they ignore the uncertainty in inclination.  Several spectroscopic orbits have been published for this pair (\\citealt{Har1925}; \\citealt{Bak1984}; \\citealt{Duq1991}; SR03), enabling the spectroscopic orbital elements to be well-constrained.  We present an updated spectroscopic solution based on these prior data and our own radial velocity measurements (\\S\\,\\ref{sec:historical}, \\S\\,\\ref{sec:SpecOrb}).  Our visual orbit leverages these spectroscopic solutions and derives all orbital elements for this binary (\\S\\,\\ref{sec:visualorbit}), leading to accurate component masses (\\S\\,\\ref{sec:massest}).  This work utilizes a very precise parallax measure for this radio-emitting binary obtained by \\citet{Les1999} using Very Long Baseline Interferometry (VLBI).  Their parallax of 43.93 $\\pm$ 0.10 mas is about 10 times more precise than the \\textit{Hipparcos} catalog value of 46.11 $\\pm$ 0.98 mas and 12 times more precise than the \\citet{van2007} measure of 47.35 $\\pm$ 1.20 mas.  The \\citeauthor{Les1999} value is 2.2-$\\sigma$ and 2.9-$\\sigma$ lower than the \\textit{Hipparcos} and \\citeauthor{van2007} measures, respectively.  To check for systematic offsets, we compared the parallaxes for all overlapping stars in these three sources.  While the difference in parallax is most significant for $\\sigma^2$~CrB, we found no systematic differences.  Moreover, \\citeauthor{Les1999} performed statistical checks to verify the accuracy of their measure, so we adopt their parallax to derive the physical parameters of the component stars (\\S\\,\\ref{sec:PhyPar}).  The Center for High Angular Resolution Astronomy (CHARA) Array's unique capabilities, facilitated by the world's longest optical interferometric baselines, have enabled a variety of astrophysical studies \\citep[e.g.,][]{McA2005, Bai2007, Mon2007}.  This work utilizes the Array's longest baselines to resolve the 1.14-day spectroscopic binary, the shortest period system yet resolved.  While this is the first visual orbit determined using interferometric visibilities measured with the CHARA Array, the technique described here has regularly been employed for longer-period binaries using other long-baseline interferometers \\citep[e.g.,][]{Hum1993, Bod1999}. The $\\sigma^2$~CrB binary has a projected angular separation of about 1.1 mas in the sky, making it easily resolvable for the CHARA Array, which has angular resolution capabilities in the $K'$ band down to about 0.4 mas for binaries. ",
        "conclusions": ""
    },
    "0808/0808.3918_arXiv.txt": {
        "abstract": "I discuss current theoretical expectations of how primordial, Pop III.1 stars form. Lack of direct observational constraints makes this a challenging task. In particular predicting the mass of these stars requires solving a series of problems, which all affect, perhaps drastically, the final outcome. While there is general agreement on the initial conditions, $\\rm H_2$-cooled gas at the center of dark matter minihalos, the subsequent evolution is more uncertain. In particular, I describe the potential effects of dark matter annihilation heating, fragmentation within the minihalo, magnetic field amplification, and protostellar ionizing feedback. After these considerations, one expects that the first stars are massive $\\gtrsim 100\\sm$, with dark matter annihilation heating having the potential to raise this scale by large factors. Higher accretion rates in later-forming minihalos may cause the Pop III.1 initial mass function to evolve to higher masses. ",
        "introduction": "The first, essentially metal-free (i.e. Population III), stars are expected to have played a crucial role in bringing the universe out of the dark ages: initiating the reionization process, including the local effects of their \\ion{H}{2} regions in generating shocks and promoting formation of molecular coolants in the relic phase; photodissociating molecules; amplifying magnetic fields to possibly dynamically important strengths; and generating the mechanical feedback, heavy elements and possible neutron star or black hole remnants associated with supernovae. In these ways Pop III stars laid the foundations for galaxy formation, including supermassive black holes and globular clusters. Many of these processes are theorized to depend sensitively on the initial mass function (IMF) of Pop III stars, thus motivating its study. The formation of the first Pop III stars in a given region of the universe is expected to have been unaffected by other astrophysical sources and these have been termed Pop III.1, in contrast to Pop III.2 (McKee \\& Tan 2008, hereafter MT08). Pop III.1 are important for influencing the initial conditions for future structure formation and for having their properties determined solely by cosmology. There is also the possibility, described in Sect.~\\ref{S:DM}, that Pop III.1 star formation may be sensitive to the properties of weakly interacting massive particle (WIMP) dark matter.  Unfortunately, at the present time and in the near future we expect only indirect observational constraints on the Pop III IMF. The epoch of reionization can be constrained by CMB polarization (Page et al. 2007) and future high redshift 21~cm HI observations (e.g. Morales \\& Hewitt 2004). Metals from individual Pop III supernovae may have imprinted their abundance patterns in very low metallicity Galactic halo stars (Beers \\& Christlieb 2005) or in the Ly-$\\alpha$ forest (Schaye et al. 2003; Norman, O'Shea, \\& Pascos 2004). Light from the first stars may contribute to the observed NIR background intensity, (e.g. Santos, Bromm, \\& Kamionkowski 2002), and its fluctuations (Kashlinsky et al. 2004; c.f. Thompson et al. 2007). If massive, supernovae marking the deaths of the first stars may be observable by JWST (Weinmann \\& Lilly 2005). If these supernovae produce gamma-ray bursts then these may already be making a contribution to the population observed by SWIFT (Bromm \\& Loeb 2002).    The lack of direct observations of Pop III star formation means theoretical models lack constraints, which is a major problem for treating such a complicated, nonlinear process. Numerical simulations have been able to start with cosmological initial conditions and advance to the point of protostar formation (see Yoshida et al., these proceedings), but progressing further through the protostellar accretion phase requires additional modeling of complicated processes, including a possible need to include extra physics such as WIMP annihilation and magnetic fields. Building up a prediction of the final mass achieved by the protostar, i.e. the initial mass (function) of the star (population), is akin to building a house of cards: the reliability of the structure becomes more and more precarious.  In this article we summarize theoretical attempts to understand the formation process and resulting IMF of Pop III.1 stars. We have reviewed much of these topics previously (Tan \\& McKee 2008), so here we concentrate on a discussion of some of the more uncertain aspects in these models, including the potential effects of WIMP annihilation on Pop III.1 star formation, fragmentation during Pop III.1 star formation, the generation of magnetic fields, the uncertainties in predicting the IMF from feedback models, and the evolution of the Pop III.1 IMF. Note, when discussing possible fragmentation during the formation of a Pop III.1 star, we will consider all stars that result from the same minihalo to be Pop III.1, i.e. they are unaffected by astrophysical sources external to their own minihalo. ",
        "conclusions": ""
    },
    "0808/0808.1295_arXiv.txt": {
        "abstract": "We show that the coupled two-fluid gravitating system (e.g. stiff matter and 'vacuum energy') could trap nonlinear gravitational waves (e.g. Einstein-Rosen waves). The gravitational wave amplitude varies harmonically in time transferring the energy coherently to the stiff matter wave, and then the process goes to the backward direction. This process mimics the behaviour of trapped electromagnetic waves in two-level media. We have defined the limits for the frequency of this energy transfer oscillations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.1540_arXiv.txt": {
        "abstract": "{Cygnus~X-1 and LS~5039 are two X-ray binaries observed at TeV energies. Both sources are compact systems, contain jet-like (radio) structures, and harbor very luminous O stars. A TeV signal has been found around the superior conjunction of the compact object in both objects, when the highest gamma-ray opacities are expected.} {We investigate the implications of finding TeV emission from Cygnus~X-1 and LS~5039 around the superior conjunction, since this can give information on the system magnetic field and the location of the TeV emitter.} {Using the very high-energy spectra and fluxes observed around the superior conjunction in Cygnus~X-1 and LS~5039, we compute the absorbed luminosity that is caused by pair creation in the stellar photon field for different emitter positions with respect to the star and the observer line of sight. The role of the magnetic field and electromagnetic cascading are discussed. For the case of inefficient electromagnetic cascading, the expected secondary synchrotron fluxes are compared with the observed ones at X-ray energies.} {We find that, in Cygnus~X-1 and LS~5039, either the magnetic field in the star surroundings is much smaller than the one expected for O stars or the TeV emitter is located at a distance $> 10^{12}$~cm from the compact object.} {Our results strongly suggest that the TeV emitters in Cygnus~X-1 and LS~5039 are located at the borders of the binary system and well above the orbital plane. This would not agree with those models for which the emitter is well inside the system, like the innermost-jet region (Cygnus~X-1 and LS~5039; microquasar scenario), or the region between the pulsar and the primary star (LS~5039; {\\it standard}{} pulsar scenario).} ",
        "introduction": "Cygnus~X-1 and LS~5039 are high-mass X-ray binaries presenting radiation in the very high-energy (VHE) range (Aharonian et al. \\cite{aharonian05}; Albert et al. \\cite{albert07}). Both sources show extended outflows that generate non-thermal radio emission (Stirling et al. \\cite{stirling01}; Paredes et al. \\cite{paredes00}) and are located at distances of $\\approx 2.1$~kpc (Ziolkowski \\cite{ziolkowski05}) and  $\\approx 2.5$~kpc (Casares et al. \\cite{casares05}; C05 hereafter), respectively. The primary object in these two systems is an O star  (Gies \\& Bolton \\cite{gies86}; C05), and the compact object is a $\\sim 20$~M$_{\\odot}$ black hole in the case of Cygnus~X-1 (Ziolkowski \\cite{ziolkowski05}) and is still of unknown nature in the case of LS~5039 (C05). The detection of jet-like radio emitting structures has  led to the classification of these sources as microquasars (Stirling et al. \\cite{stirling01};  Paredes et al. \\cite{paredes00}). Later on, however, not detecting accretion features in the X-ray spectrum of LS~5039 has been interpreted as a hint of a non-accreting pulsar (e.g. Martocchia et al. \\cite{martocchia05}). On the other hand, Cygnus~X-1 is a firmly established accreting source (e.g. Pringle \\& Rees \\cite{pringle72}). It is worth noting that LS~5039 has been associated to an EGRET source (Paredes et~al. \\cite{paredes00}), whereas Cygnus~X-1 has no GeV association.  Cygnus~X-1 and LS~5039 are relatively similar to each other, as shown in Table~1. Besides showing extended radio emission, the two systems harbor massive and hot stars, have quite compact size, $\\sim 0.1$--0.2~AU, and have shown a TeV signal around the superior conjunction of the compact object (SUPC). The TeV emission properties of both sources around SUPC are presented as well in Table~1 (bottom). In the case of Cygnus~X-1, evidence of detection above 100~GeV of 4.1~$\\sigma$ significance has been reported once right before SUPC, at phase $\\approx 0.9$ (Albert et al. \\cite{albert07}). In LS~5039, the VHE emission has been clearly detected (40~$\\sigma$) all along the orbit, being in fact periodic (Aharonian et al. \\cite{aharonian06}) with the orbital period (C05). At the phase $0.00\\pm 0.05$ ($\\sim$ SUPC), the detection significance is 6.1~$\\sigma$. There is however a big difference between the two sources. Whereas Cygnus~X-1 shows a thermal (comptonized) hard X-ray spectrum (e.g. Sunyaev \\& Trumper \\cite{sunyaev79}), with a luminosity $L_{\\rm X}\\sim 10^{37}$~erg~s$^{-1}$ in the epoch when the TeV emission was detected (Albert et al. \\cite{albert07}), the X-ray radiation from LS~5039 is three orders of magnitude less luminous, i.e. $L_{\\rm X}\\sim 10^{34}$~erg~s$^{-1}$, and could be dominated by a non-thermal component (e.g. Bosch-Ramon et~al. \\cite{bosch07}).  In Cygnus~X-1 and LS~5039, if the TeV emitter is deep inside the system and behind the primary star around SUPC, the photon-photon absorption opacity ($\\tau$) for TeV gamma rays in the stellar photon field will be $\\gg 1$ in the direction  of the observer. In this situation, either the magnetic field is low enough to allow efficient electromagnetic (EM) cascading to develop (e.g. Aharonian et al. \\cite{aharonian06b}; Bednarek \\& Giovanelli \\cite{bednarek07}; Orellana et al. \\cite{orellana07}), effectively reducing $\\tau$ around SUPC, or the absorbed energy will be reradiated mainly by secondary pairs via synchrotron emission (e.g. Bosch-Ramon, Khangulyan, \\& Aharonian \\cite{bosch08}). In the latter case, the high (effective) opacities will require an injected gamma-ray luminosity much higher than the observed one, and the synchrotron luminosity of the secondary pairs will be similar to that of the absorbed gamma rays. Because the secondary synchrotron luminosity should not exceed the observational constraints, we can put strong restrictions on either the magnetic field strength or on the emitter location with respect to the star and the line of sight of the observer. In this work, we discuss the implications of two possible situations, i.e. a {\\it low} versus a {\\it high} magnetic field in the stellar surroundings for the TeV emitters in Cygnus~X-1 and LS~5039. ",
        "conclusions": "The detection of Cygnus~X-1 and LS~5039 around SUPC, plus X-ray data, give strong constraints on the magnetic field and the location of the TeV emitter in these systems. In summary, if the TeV emitter is close to the compact object, the magnetic field in the stellar surroundings must be smaller than a few Gauss. These values will allow IC cascades to develop, effectively reducing the photon-photon absorption opacities. However, this magnetic field seems rather low when compared with the values expected in O stars. For a more realistic $B_*>B_{\\rm c}$, EM cascading becomes inefficient, and the large amount of energy absorbed in the stellar photon field will be reemitted by the secondary pairs mainly via the synchrotron process. For $B_*>B_{\\rm c}$, this radiation will peak roughly in the X-ray band. In such a case, to avoid the violation of the observational constraints, the TeV emitter must be located at distances $>10^{12}$~cm from the compact object in both Cygnus~X-1 and LS~5039. It is consistent with how, as noted by K08, extreme acceleration rates are required to explain the highest energy photons from LS~5039 if the accelerator is deep inside the system. Remarkably, the low hydrogen column densities inferred using X-ray data would also hint at an X-ray emitter far from the compact object (Bosch-Ramon et al. \\cite{bosch07}).  A TeV emitter far from the compact object requires a physical mechanism to transport energy to the borders of the system to power the VHE emission. The source of this energy could be accretion in Cygnus~X-1, or either accretion (microquasar scenario) or a powerful pulsar wind (non-accreting pulsar scenario) in LS~5039. The energy carrier could be the jet seen in radio in Cygnus~X-1. In the case of LS~5039, energy could be transported by a jet (the jet-like structure seen in radio) powered by radiatively inefficient accretion (e.g. Bogovalov \\& Kelner \\cite{bogovalov05}), since no accretion X-ray features have been found. In the {\\it standard} pulsar scenario for LS~5039 (e.g. Dubus \\cite{dubus08}; Sierpowska-Bartosik \\& Torres \\cite{sierpowska08}), the TeV emitter would be located close to the pulsar or around the line joining the pulsar and the star. This would imply that, for reasonable energy budgets, very little emission should be expected due to absorption around SUPC (see Fig.~2 in Dubus \\cite{dubus06} -lower panel-, Fig.~4 in Dubus et al. \\cite{dubus08}, and Fig.~2 in Sierpowska-Bartosik \\& Torres \\cite{sierpowska08}), unlike how it is observed. This does not agree with the {\\it standard} pulsar scenario for LS~5039, although in a more general case, a supersonic outflow produced in the star/pulsar wind colliding region (see Bogovalov et al. \\cite{bogovalov08}) may transport energy efficiently to the system borders and farther out.  To conclude, our results strongly suggest that the TeV emitters in Cygnus~X-1 and LS~5039 are located at the borders of the binary system, well above the orbital plane. This would not be compatible with those models for which the emitter is well inside the system, like the innermost-jet region (Cygnus~X-1 and LS~5039; microquasar scenario) or the region between the pulsar and the primary star (LS~5039; {\\it standard} pulsar scenario)."
    },
    "0808/0808.3897_arXiv.txt": {
        "abstract": "We study an accretion disk in which three different regions may coexist: MHD turbulent regions, dead zones and gravitationally unstable regions. Although the dead zones are stable, there is some transport due to the Reynolds stress associated with waves emitted from the turbulent layers. We model the transport in each of the different regions by its own $\\alpha$ parameter, this being 10 to $10^{3}$ times smaller in dead zones than in active layers.  In gravitationally unstable regions, $\\alpha$ is determined by the fact that the disk self--adjusts to a state of marginal stability.  We construct steady--state models of such disks.  We find that for uniform mass flow, the disk has to be more massive, hotter and thicker at the radii where there is a dead zone.  In disks in which the dead zone is very massive, gravitational instabilities are present.  Whether such models are realistic or not depends on whether hydrodynamical fluctuations driven by the turbulent layers can penetrate all the way inside the dead zone.  This may be more easily achieved when the ratio of the mass of the active layer to that of the dead zone is relatively large, which in our models corresponds to $\\alpha$ in the dead zone being about 10\\% of $\\alpha$ in the active layers.  If the disk is at some stage of its evolution not in steady--state, then the surface density will evolve toward the steady--state solution.  However, if $\\alpha$ in the dead zone is much smaller than in the active zone, the timescale for the parts of the disk beyond a few AU to reach steady--state may become longer than the disk lifetime.  Steady--state disks with dead zones are a more favorable environment for planet formation than standard disks, since the dead zone is typically 10 times more massive than a corresponding turbulent zone at the same location. ",
        "introduction": "Protoplanetary disks are believed to have a complex structure.  In the parts which are ionized enough ({\\em active} zones), the magnetic field couples to the gas and the magnetorotational instability (MRI; Balbus~\\& Hawley~1991) develops, leading to turbulence and angular momentum transport.  Some parts however are too cold and dense for the ionization to reach the level required for MHD turbulence to be sustained. However, in these so--called {\\em dead} zones, there may still be some low level of transport due to Reynolds or non turbulent Maxwell stresses (Fleming \\& Stone 2003, Turner \\& Sano 2008).  Finally, parts of the disk may be dense enough that gravitational instabilities develop, redistributing mass and angular momentum in such a way that the disk settles into a state of marginal stability.  To date, there is no numerical simulations of such global disks.  Global three dimensional MHD simulations so far have modeled stratified and fully turbulent disks, but with no dead zones (Fromang~\\& Nelson~2006).  It is of interest to study whether non--uniform disks composed of different regions as described above can be in steady--state or not. Previous studies (Gammie 1999, Armitage et al. 2001) assumed that no transport at all took place in the dead zone.  Therefore mass could only pile--up there until gravitational instabilities would develop.  If the level of transport is non zero in the dead zone however, the disk may be able to adjust to reach a steady--state.  It is this possibility that we investigate in this paper.  In determining the characteristics of the disk vertical strcuture, we ignore reprocessing of the stellar radiation.  Only heating associated with the different (magnetic and hydrodynamic) stresses is considered.  Reprocessing has been shown to be important when the disk surface is flared (Kenyon ~\\& Hartmann 1987, Chiang \\& Goldreich 1997, D'Alessio et al. 1998, Dullemond et al. 2001).  It provides an extra heating source that dominates the disk beyond a few AU and has important consequences on the structure of the dead zone (Matsumura \\& Pudritz 2006).  However, in the models we present below, the parts of the disk beyond 1~AU are shielded from the stellar radiation by the inner parts. These models, which do not include reprocessing, are therefore self--consistent.  The plan of the paper is as follows.  In section~\\ref{sec:modeling} we describe the way we model the transport in a disk with turbulent regions, dead zones and gravitationally unstable regions.  In section~\\ref{sec:vertical} we present the equations governing the disk vertical structure and describe how they are solved.  In section~\\ref{sec:results} we present the results of the calculations, and discuss them in section~\\ref{sec:discussion}.  We find that steady--state solutions exist and correspond to disks which are thicker, hotter and more massive than disks without dead zones. ",
        "conclusions": "\\label{sec:discussion}  In this paper, we have considered a disk in which transport is produced by either: i) Maxwell stress resulting from MHD turbulence in the (gravitationally stable) regions where the gas couples well to the magnetic field; ii) gravitational stress in the parts which are gravitationally unstable; iii) Reynolds stress associated with hydromagnetic waves driven by the turbulence in adjacent layers and Maxwell stress associated with large scale field siphoned off from these turbulent layers in the parts which are not turbulent and are gravitationally stable (dead zones).  We have modeled the transport using the Shakura \\& Sunyaev prescription, with $\\alpha_D$ being 10 to $10^3$ times smaller than $\\alpha_T$, and $\\alpha_G$ being adjusted such as to give a Toomre parameter $Q \\sim 1.5$ in the gravitationally unstable regions.  Although such a modeling is in principle not valid when the transport is non--local, we believe it does not invalidate the gross features of the models we have presented.  We have found that steady--state models of such a disk exist, and that they are physically reasonable. Since $\\dot{M} = 3 \\pi \\langle \\nu \\rangle \\Sigma \\propto \\alpha c_s H \\Sigma$, mass flow through the disk can be uniform only if the disk is more massive, hotter and thicker at the radii where the vertically--averaged value of $\\alpha$ is smaller, i.e. where there is a dead zone.  Note that even though the disk is thicker at the locations where there is a dead zone, it remains thin for the parameters investigated here.  In disks in which the dead zone tends to be massive, which is the case for the lowest values of $\\alpha_D$ and/or the lowest values of $\\Sigma_T$ investigated here, gravitational instabilities control the dynamics of part of the disk.  Whether these models are realistic or not depends on whether hydrodynamical fluctuations driven by the turbulent layers can penetrate all the way inside the dead zone.  This may be more easily achieved when the ratio of the mass of the dead zone to that of the active layer is the smallest, which in our models corresponds to $\\alpha_D/\\alpha_T=0.1$.  If the disk is at some stage of its evolution out of steady--state, then the surface density will change in such a way as for the disk to evolve toward steady--state.  To see this, let us consider a disk in which $\\dot{M}$ decreases inward at some location. Then mass will accumulate there as accretion through this region is slower, until $\\Sigma$ is large enough for the flow to be steady.  In contrast, if $\\dot{M}$ increases inward at some location, accretion there will be faster and mass will be depleted until $\\Sigma$ is reduced to the level where a steady--state is reached.  The timescale $t_{\\nu}$ for establishing steady--state at the radii where there is a dead zone may be long though, as it is given by:  \\begin{displaymath} t_{\\nu} = \\frac{r^2}{3 \\langle \\nu \\rangle} \\sim \\frac{1}{3 \\langle \\alpha \\rangle} \\left( \\frac{r}{H}  \\right)^2 \\Omega^{-1} , \\end{displaymath}  \\noindent where $\\langle \\alpha \\rangle = \\int_{-H}^H \\rho \\alpha dz/ \\Sigma$.  With $\\langle \\alpha \\rangle = 10^{-2}$ and $H/r=0.1$, we get $t_{\\nu} \\sim 5 \\times 10^2$ and $6 \\times 10^3$~years at 1 and 5~AU, respectively, which is much smaller than the disk lifetime.  However, if the dead zone at these radii is vertically very extended, we have $\\langle \\alpha \\rangle \\sim \\alpha_D$, so that the viscous timescale can get 100 times longer if $\\alpha_D=10^{-4}$.  At 1~AU we still get a timescale significantly smaller than the disk lifetime, but this is not the case at 5~AU, where the disk may not be able to reach steady--state. If these parts of the disk build--up from mass inflowing from further out, then the mass in the dead zone will slowly increase.  This process will never produce outbursts as envisioned by Gammie~(1999) though. Indeed, before enough mass could pile--up for an outburst to be produced, a steady--state would be achieved in which mass would be transported either by density waves or Reynolds stress driven by the turbulent layers.  Note that the spectral energy distribution of a steady disk such as those considered in this paper is not different from that of a standard disk with no dead zone, as the total flux emitted at the surface depends only on $\\dot{M}$.  The surface temperature also does not depend on $\\Sigma_T$, as can be seen from equation~(\\ref{Ts}).  However, the temperature just below the surface is significantly larger in disks with dead zones than in standard disks.  The chemistry there would therefore be different.  Grain properties may also differ.  To study this, irradiation of the disk by the central star would have to be taken into account.  Although it will not be important beyond the regions where there is a dead zone, as these parts are shielded from the stellar radiation due to the puffing up of the disk, it will be important in the parts where $H/r$ increases.  Finally, we comment that a steady--state disk with a dead zone is a more favorable environment for planet formation than a standard disk, as the dead zone, which in general encompasses the region of planet formation, is significantly more massive.  Planet formation would therefore be faster there (Lissauer 1993).  Note however that there is an issue as how to stop type~I migration of the cores in these dead zones. Turbulence, which has been shown to alter this type of migration (Nelson \\& Papaloizou 2004), cannot be invoked.  But if large scale fields are siphoned off in these regions from the turbulent layers, they may prevent the cores from migrating too fast (Terquem 2003, Fromang et al. 2005, Muto et al. 2008).  Planetary cores may also be stopped at the interface between the dead zone and the inner turbulent region, where the surface mass density varies sharply, as suggested by Masset et al. (2006).  This process however relies on the (positive) corotation torque becoming dominant over the Lindblad torque due to the density gradient and was studied in the context of a laminar viscous disk.  It is not clear how the corotation torque is affected by the turbulence at the transition between the active and dead zones, and whether it would still be as efficient in this context."
    },
    "0808/0808.3883_arXiv.txt": {
        "abstract": "{ The Medusa (NGC~4194) is a well-studied nearby galaxy with the disturbed appearance of a merger and evidence for ongoing star formation. In order to test whether it could be the result of an interaction between a gas-rich disk-like galaxy and a larger elliptical, we have carried out optical and radio observations of the stars and the gas in the Medusa, and performed $N$-body numerical simulations of the evolution of such a system. We used the Nordic Optical Telescope to obtain a deep V-band image and the Westerbork Radio Synthesis Telescope to map the large-scale distribution and kinematics of atomic hydrogen. A single \\hi tail was found to the South of the Medusa with a projected length of $\\sim 56$~kpc ($\\sim 5'$) and a gas mass of $7\\cdot 10^8\\,M_{\\odot}$, thus harbouring about one third of the total \\hi mass of the system. \\hi was also detected in absorption toward the continuum in the center. \\hi was detected in a small nearby galaxy to the North-West of the Medusa at a projected distance of 91\\,kpc. It is, however, unlikely that this galaxy has had a significant influence on the evolution of the Medusa. The simulations of the slightly prograde infall of a gas-rich disk galaxy on an larger, four time more massive elliptical (spherical) galaxy reproduce most of the observed features of the Medusa.Thus, the Medusa is an ideal object to study the merger-induced star formation contribution from the small galaxy of a minor merger. ",
        "introduction": "The Medusa galaxy (NCG\\,4194, Arp~160, IZw33) is a nearby galaxy ($D=39$~Mpc) with the disturbed appearance typical of a merger (see Fig.~\\ref{n4194chap_contimage}). It has been interpreted as a candidate for a minor merger between a large elliptical and a small spiral galaxy, a so-called S+E merger \\citep{2000A&A...362...42A}. The galaxy earned its nickname from the fuzzy and knotty optical tail to the North, reminiscent of hair or tentacles. A sharp shell-like structure is visible on the opposite side of the galaxy, confining the optical main body. A second, fainter shell can be seen further out on deep images (e.g., \\citealt{optsample}).  Major mergers between two gas-rich disk galaxies can lead to ultraluminous infrared galaxies (ULIRGs, $L_{FIR} > 10^{12}\\,L_{\\odot}$) due to an extreme increase of star formation \\citep{2003AJ....126.1607S}. Strong gas inflows are triggered in these mergers, increasing the gas density in the centre and thus leading to the observed starburst \\citep[e.g.,][]{1991ApJ...370L..65B,1996ApJ...471..115B,1996ApJ...464..641M}. In these simulations the merger remnant evolves into a regular elliptical galaxy. In contrast, in minor, i.e., unequal-mass mergers only a moderate, if at all, starburst is induced implying a lower far-infrared (FIR) luminosity (less than $10^{11}\\,L_{\\odot}$) \\citep{2008MNRAS.384..386C}. The triggered star formation and the remnant galaxy type strongly depend on the mass ratio of the merging pair, but not so much on the total gas mass \\citep{2005A&A...437...69B,2008MNRAS.384..386C}.  In contrast to most nearby minor mergers, in which the major galaxy is a spiral and thus dominating the star formation (e.g., M\\,51), the Medusa gives us an ideal opportunity to investigate the influence of the merger on the smaller companion because it is the only gas-rich partner here. Since the merger-induced star formation does not depend on the gas content of the companion \\citep{2008MNRAS.385L..38M}, the results of the Medusa might be applicable to other minor mergers, independent on the type of the large partner. \\\\ Another significant difference between ULIRGs and galaxies like the Medusa is the extent of the starburst induced by the merger. While the burst region in ULIRGs is compact and typically concentrated to the inner kiloparsec, (e.g., \\citealt{1996ARA&A..34..749S}), the Medusa exhibits an extended region of ongoing star formation which is, however, not as intense as in ULIRGs (\\citealt{1990ApJ...364..471A}, \\citealt{1994ApJ...422...73P}). \\cite{2004AJ....127.1360W} analyzed optical and ultraviolet Hubble Space Telescope (HST) images of the nucleus and found star-forming knots younger than 20\\,Myr in the center. They derived a star formation rate (SFR) of $\\rm \\sim 6\\,M_{\\odot}\\ yr^{-1}$, but argued that the overall  SFR in the Medusa might be as high as $\\rm 30\\,M_{\\odot}\\ yr^{-1}$ since the knots produce only 20\\% of the UV flux in the center. Based on recent HST spectroscopic observations, \\cite{2006AJ....131..282H} estimate an even higher total SFR of $\\rm \\sim 46\\,M_{\\odot}\\ yr^{-1}$, and ages of individual star forming regions of 5.5--10.5\\,Myr.  The Medusa possesses a reservoir of gas, both molecular and atomic. High resolution aperture synthesis maps obtained with the Owens Valley Radio Observatory revealed extended CO emission on a total scale of 25$''$ (4.7~kpc) \\citep{2000A&A...362...42A}. The CO morphology is complex, occupying mainly the center and the north-eastern part of the main optical body. The extended CO gas follows two prominent dust lanes: one which is crossing the central region at right angle with respect to the optical major axis, and a second which curves to the north-east and into the beginning of the northern optical tail. \\cite{2000A&A...362...42A} suggested that the central starburst is being fueled by gas flows along the central dust lane. The presence of multiple spectral peaks ($\\Delta$V $\\approx$ 200 \\kms) suggests that at least two velocity systems coexist in the CO emitting gas. Either this is an effect of substantial warping of the main body, or simply the unrelaxed, overlapping systems of the progenitor galaxies. The multiple features seem to be related to the extended dust lanes, but several velocity components co-exist also at the base of the optical tail. Single-dish measurements revealed the presence CO emission even inside  the optical tail out to 4.7\\,kpc away from the center \\citep{2001A&A...372L..29A} where  there is no sign of ongoing star formation \\citep{1990ApJ...364..471A}. The total molecular gas mass in the inner 2 kpc was estimated to $2\\cdot 10^9\\,M_{\\odot}$. Thus, the star formation efficiency is almost as high as in ULIRGs although the molecular gas density is significantly lower. \\citep{2000A&A...362...42A}. Besides uncertainties in the molecular gas conversion factor, this gives a hint that the ISM in these mergers is different compared to the molecular cloud complexes in normal galaxies and also in extreme starbursts. \\cite{2000A&A...362...42A} argue that the observed solid-body rotation may offer a shear-free environment. This may support cloud collapse and thus star formation. The star formation rate might have increased because a large amount of molecular gas was transported to the centre on short time scales by the interaction. \\begin{figure*} \\includegraphics[angle=270,width=8.5cm]{0408fig01.pdf} \\includegraphics[angle=270,width=8.5cm]{0408fig02.pdf} \\caption[Optical image and 20\\,cm continuum contour map of NGC\\,4194] {Left: Nordic Optical Telescope V-band image of the Medusa. Right: Contour map of the 20\\,cm continuum emission superimposed on the V-band image. The synthesised beam is shown in the lower left corner. The contour levels are 0.2, 0.5, 2, 5, 20, 50, 70\\,mJy\\ beam$^{-1}$. } \\label{n4194chap_contimage} \\end{figure*}  The galaxy contains a similar amount of atomic gas as molecular gas (\\mhi $\\sim 2.2 \\cdot 10^9\\,M_{\\odot}$), inferred from the single-dish \\hi emission spectrum obtained by \\cite{1981ApJ...247..823T} using the Greenbank 91\\,m telescope. Atomic gas has recently been detected in absorption toward the center of the Medusa, through sub-arcsecond MERLIN observations \\citep{2005A&A...444..791B}. The kinematics of the \\hi gas in the central region of the galaxy is similar to that of the molecular gas traced by the CO emission \\citep{2000A&A...362...42A}, which suggests that the atomic and molecular component of the cold interstellar medium are probably physically related. Two compact radio sources were found in the central parsec, embedded in a region of diffuse radio continuum emission. Even though the presence of a weak AGN cannot be ruled out completely, most (if not all) of the continuum emission is related to the starburst \\citep{2005A&A...444..791B}.  NGC~4194 displays conflicting evidence about its merger age. The sub-arcsecond resolution MERLIN data revealed two compact radio components in the center, separated by $\\rm \\sim 0^{''}.35$, corresponding to 65\\,pc only \\citep{2005A&A...444..791B}. The optical tail is fairly diffuse, which is suggestive of an advanced merger where the progenitors have coalesced. In contrast, the fairly widespread CO distribution, irregular optical isophotes, multiple spectral CO features, and the distorted central body hint toward a younger system.  In this paper, we present observations and $N$-body numerical simulations of the stars and the gas in the Medusa merger. We mapped the large-scale distribution and kinematics of the \\hi gas with the Westerbork Radio Synthesis Telescope (WSRT) and used the Nordic Optical Telescope (NOT) to obtain a deep V-band image of the distribution of the stars in the Medusa. Then we carried out numerical simulations which we compared to the observations. \\hi 21 cm line observations are a powerful tool to trace the dynamical history of interacting and merging galaxies and provide constraints to dynamical models of the history of the interaction. They complement CO millimeter-wave line observations that trace mainly the inner part of galaxies where the molecular gas resides. Furthermore, they offer is the only way to trace the distribution and kinematics of gas in large scale tidal features. The observations are presented in section~2 and 3, and the simulations in section~4. Follows a discussion.  \\begin{table} \\caption[Basic properties of NGC\\,4194]{Basic properties of NGC\\,4194. The distance is based on ${\\rm H_0 = 75\\,km\\ s^{-1}\\ Mpc^{-1}}$. } \\label{n4194chap_intro} \\centering \\begin{tabular}{lc} \\hline\\hline RA (2000) & 12 14 09.5\\\\ DEC (2000) & +54 31 37 \\\\ $\\rm v_{hel,opt}$ (km\\ s$^{-1}$) \t\t& 2442 \\\\ D (Mpc) \t\t\t\t& 39\\\\ $\\rm L_B$ ($10^9\\,L_{\\odot}$) \t\t& 15.7\\\\ $\\rm L_{FIR}$ ($10^{10}\\,L_{\\odot}$) \t& 8.5\\\\ FIR flux at 60\\,$\\rm \\mu$m (Jy) \t& 25.66\\\\ FIR flux at 100\\,$\\rm \\mu$m (Jy) \t& 26.21\\\\ \\hline  \\hline \\end{tabular} \\end{table} ",
        "conclusions": "We have presented new optical, radio continuum and \\hi line observations of the Medusa galaxy and numerical N-body simulations of the merger of a small gas-rich disk galaxy with a larger elliptical galaxy. The optical observations consist of a deep V-band image of the distribution of the stars in the Medusa, and the radio observations of maps of the 20 cm radio continuum emission and of the large-scale distribution and kinematics of the atomic hydrogen gas. The simulations are completely self-consistent with a total of 120000 particles. The gas in the infalling disk galaxy is represented by an ensemble of clouds, all of the same mass, that can dissipate energy in inelastic collisions with each other.  \\begin{enumerate} \\item The \\hi observations revealed the presence of a single \\hi tail to the South, extending out to a projected distance of 4.95\\,', or 56 kpc from the center of the Medusa.  The tail contains $\\rm \\sim 7\\cdot 10^8\\, M_{\\odot}$ of \\hi and has no detected optical counterpart. No significant velocity gradient is observed along the tail, indicating that it is seen almost face-on.  \\item \\hi was detected in a dwarf galaxy identified at a projected distance of 91\\,kpc to the North-West of NGC\\,4194. Its systemic velocity of 2380\\,km\\ s$^{-1}$ is only 120\\,km\\ s$^{-1}$ lower than that of the Medusa. This dwarf galaxy harbours $\\rm \\sim 7 \\cdot 10^7\\,M_{\\odot}$ of \\hi. No connection, neither in the optical nor the \\hi could be found between the dwarf galaxy and NGC\\,4194 and it is unlikely that it has had an influence on the dynamical evolution of the Medusa.  \\item We found a large misalignment between optical and gaseous tidal features. The diffuse optical tail is situated on the opposite side of the  \\hi tail, the optical shells do not seem to be connected to the \\hi tail. Although the shells are within the region of the \\hi  tail, no connection is apparent neither in morphology nor in velocity distribution.  \\item We estimated a SFR of 23\\,$\\rm M_{\\odot}\\ yr^{-1}$ based on the 20\\,cm continuum flux density, and of 10\\,$\\rm M_{\\odot}\\ yr^{-1}$ from the FIR fluxes.  This is consistent with a rather intense ongoing burst of star formation in the Medusa, in agreement with the optical and UV observations of \\cite{2004AJ....127.1360W} .  \\item \\hi was detected in absorption toward the radio continuum in the center.  \\item The simulations of a slightly prograde merger between a gas-rich disk galaxy with a total mass four times less than that of the elliptical give the best match to the observations. A extended single tidal tail develops in the outer part while the small galaxy falls in and oscillates around the center of the large elliptical, creating shell-like structures around the elliptical, best visible in the stars. The overall distribution of the gas is similar to that of the stars, although more gas falls toward the center because of the dissipative nature and more gas is found of the extended tail because of the initially more extended  distribution of the gas. The tail is mostly formed by material of the outer \\hi disk of the progenitor spiral galaxy, and the surface brightness of the stellar material in the tail is low.  \\end{enumerate}"
    },
    "0808/0808.3767_arXiv.txt": {
        "abstract": "We investigate the angular two-point correlation function of temperature in the WMAP maps. Updating and extending earlier results, we confirm the lack of correlations outside the Galaxy on angular scales greater than about 60 degrees at a level that would occur in 0.025 per cent of realizations of the concordance model.  This represents a dramatic increase in significance from the original observations by the COBE-DMR and a marked increase in significance from the first-year WMAP maps.  Given the rest of the reported angular power spectrum $C_\\ell$, the lack of large-angle correlations that one infers outside the plane of the Galaxy requires covariance among the $C_\\ell$ up to $\\ell=5$.  Alternately, it requires both the unusually small (5 per cent of realizations) full-sky large-angle correlations, \\textit{and} an unusual coincidence of alignment of the Galaxy with the pattern of cosmological fluctuations (less than 2 per cent of those 5 per cent).  We argue that unless there is some undiscovered systematic error in their collection or reduction, the data point towards a violation of statistical isotropy.  The near-vanishing of the large-angle correlations in the cut-sky maps, together with their disagreement with results inferred from full-sky maps, remain open problems, and are very difficult to understand within the concordance model. ",
        "introduction": "Over a decade ago, the Cosmic Background Explorer Differential Microwave Radiometer (COBE-DMR) first reported a lack of large-angle correlations in the two-point angular-correlation function, $\\mathcal{C}(\\theta)$, of the cosmic microwave background (CMB) \\citep{DMR4}.  This was confirmed by the Wilkinson Microwave Anisotropy Probe (WMAP) team in their analysis of their first year of data~\\citep{Spergel2003}, and by us in the WMAP three-year data~\\citep{wmap123}. Those findings have since been confirmed by \\cite{Hajian:2007pi} and \\cite{Bunn:2008zd}. Here, we present a more detailed analysis of the three-year and (for the first time) of the five-year WMAP data, confirming and strengthening our previous results.  There is a common misconception that this lack of angular correlations is equivalent to the low quadrupole in the two-point angular power spectrum, which on its own does not have sufficient statistical significance to challenge the canonical paradigm.  It is typically assumed both that the angular power spectrum, $C_\\ell$, and the two point angular correlation function, $\\mathcal{C} (\\theta)$, contain the same information; and that consequently studying one is as good as studying both.  Actually, the exact informational equivalence between $C_\\ell$ and $\\mathcal{C} (\\theta)$ holds only when the full sky is observed. Statistically they are equivalent only when the sky is statistically isotropic.  But even if $C_\\ell$ and $\\mathcal{C} (\\theta)$ did contain the same information, we are well aware that transforming between different representations of the same information --- a time series and its Fourier transform for example --- may make a real signal in the data easier or harder to detect.  The Doppler peaks of the CMB, so clearly visible in the $C_\\ell$ representation are quite invisible in the two-point correlation function.  In this work we demonstrate that outside the region of the sky dominated by our Galaxy, both of the CMB-dominated microwave bands --- V and W --- as well as the Independent Linear Combinations (ILC) map synthesized from them as the best map of the CMB, possess above $60\\degr$ a level of two point angular correlation so low that no more than 0.025 per cent of random realizations of the best-fit $\\Lambda$CDM model would have less.  Indeed, above $60\\degr$, $\\mathcal{C} (\\theta)$ is almost entirely due to correlations involving points inside the Galaxy.  This level of statistical unlikelihood ($\\mathcal{O}(10^{-4})$) should be contrasted with what could be inferred from COBE ($\\mathcal{O}(10^{-2})$). This is a strong argument against the criticism that its identification as an anomaly is \\textit{a posteriori}. It may have been \\textit{a posteriori} for COBE, but its reidentification in WMAP at dramatically increased statistical significance is precisely how one goes about confirming that anomalies are actually present rather than statistical accidents of an observation.  While the \\textit{full-sky} $\\mathcal{C} (\\theta)$ itself has unexpectedly low large-angle correlations (occurring only in 5 per cent of random realizations of the concordance model), we find that what little correlation it does have is effectively ``hidden'' behind the Galaxy. In fact, we find that a random rotation of the Galactic cut is as successful in masking the power only 2 per cent of the time, in agreement with the previous claim that the little correlation above $60\\degr$ stems solely from two specific regions within the Galactic cut covering just 9 per cent of the sky \\citep{Hajian:2007pi}.  This further underlies the striking lack of power outside the Galactic cut, and calls into question cosmological uses of full-sky maps.  Finally, we demonstrate that the absence of large-angle correlations is emphatically not a matter just of the low quadrupole.  Rather, given the other measured multipoles, obtaining this little large-angle correlation on the cut sky (i.e. the part outside the Galaxy) requires carefully tuning $C_2$, $C_3$, $C_4$, and $C_5$.  There is also a strong indication that it is not enough to find a model in which the theoretical $C_\\ell$ yield a very small correlation function on large angular scales.  This is because, even if the theoretical $C_\\ell$ were to be set equal to those that are inferred from the cut-sky $\\mathcal{C}(\\theta)$ --- so that the expected $\\mathcal{C}(\\theta)$ nearly vanished above $60\\degr$ --- an actual realization of Gaussian-random statistically independent $a_{\\ell m}$ with these ${C}_{\\ell})$ would yield different observed $\\mathcal{C}_\\ell$ because of cosmic variance.  $C(\\theta>60\\degr)$ would then not be nearly so close to zero.  Thus getting $C(\\theta>60\\degr)$ to vanish as it does seems to require covariance among the low-$\\ell$ $C_\\ell$, and thus among $a_{\\ell m}$ of different $\\ell$.  This is in contradiction to the predictions of standard inflationary cosmological theory.  One is therefore placed between a rock and a hard place.  If the WMAP ILC is a reliable reconstruction of the full-sky CMB, then there is overwhelming evidence (\\citet{deOliveira2004,Eriksen_asym,Copi2004,Schwarz2004, lowl2,wmap123,Land2005a,Land2004a, Land2004b,Land2005b}; for a review see \\citet{Huterer_NewAst_review}) of extremely unlikely phase alignments between (at least) the quadrupole and octopole and between these multipoles and the geometry of the Solar System --- a violation of statistical isotropy that happens by random chance in far less than 0.025 per cent of random realizations of the standard cosmology.  If, on the other hand, the part of the ILC (and band maps) inside the Galaxy are unreliable as measurements of the true CMB, then the alignment of low-$\\ell$ multipoles cannot be readily tested, but the magnitude of the two-point angular correlation function on large angular scales outside the Galaxy is smaller than would be seen in all but a few of every 10,000 realizations.  We can only conclude that (i) we don't live in a standard $\\Lambda$CDM  Universe with a standard inflationary early history; (ii) we live in an extremely anomalous realization of that cosmology; (iii) there is a major error in the observations of both COBE and WMAP; or (iv) there is a major error in the reduction to maps performed by both COBE and WMAP. Whichever of these is correct, inferences from the large-angle data about precise values of the parameters of the standard cosmological model should be regarded with particular skepticism.  In the remainder of this paper we provide the detailed evidence and reasoning behind the above statements. ",
        "conclusions": "\\label{sec:conclude}  In this paper we have studied the angular correlation function in WMAP three- and five-year maps. We have clarified the relation between various definitions of the angular correlation function, and revisited our previous calculation from \\citet{wmap123} in more detail. We confirmed that power on large angular scales --- greater than about 60 degrees --- is anomalously low, at 99.975 per cent CL (see Table \\ref{tab:Cl}).  The measured angular correlation function thus disagrees with the $\\Lambda$CDM theory, but, more significantly, it is consistent with a simple phenomenological ``theory'' --- $\\mathcal{C} (\\theta\\ge60\\degr)\\equiv 0$.  The significance of this disagreement (as measured by the probability of the value of $S_{1/2}$) has now increased by a factor of over 100 since it was first observed in the COBE-DMR four-year analysis.  We have shown that the cut-sky and full-sky large-scale angular correlations differ (see Table~\\ref{tab:Cl} and Figure~\\ref{fig:ctheta}), though the source of these discrepancies remains unknown.  This shows that either the Universe is \\textit{not} statistically isotropic on large angular scales or that correlations are introduced in reconstructing the full sky from the observations.  We have shown that even given the unusually small full-sky angular correlations (95 per cent unlikely) an unusual alignment of the Galaxy with the CMB (2 per cent of realizations) is required to explain the lack of correlations outside the Galactic region.  We have further shown that simply adjusting the theoretical values of the $C_\\ell$ does not solve the problem if the sky is representative of a Gaussian random statistically isotropic process -- the cosmic variance in the $C_\\ell$ is such that less than 3 per cent of all realizations would preserve a low value of \\Shalf\\ .  From these results we argue that $\\mathcal{C}(\\theta)$ is an important quantity to study along with the usual angular power spectrum, $C_\\ell$.  The typical ``rule-of-thumb'' that low-$\\ell$ describes large angular scales is not accurate.  Any theoretical explanations for the ``missing large-scale power'' should concentrate on explaining the low $\\mathcal{C} (\\theta)$, rather than the smallness of the quadrupole and octopole, which are not nearly as significant \\citep{O'Dwyer2004}. As has been pointed out by \\cite{Rakic:2007ve} and \\cite{Bunn:2008zd}, any possible explanation of the multipole alignments that relies on an additive, statistically independent contribution to the microwave sky on top of the primordial one, increases the significance of the lack of angular correlation.  The CMB, as measured by WMAP in particular, provides much support for our current model of the Universe.  It also points the way toward new puzzles that may affect fundamental physics.  On the largest angular scales the microwave sky is inconsistent with theoretical expectations.  These discrepancies between observations and theory remains an open problem.  In the future, combining the current data with new information, such as new data from WMAP, observations from the Planck experiment, and polarization information \\citep{Dvorkin} may be key to determining the nature of the large-scale anomalies."
    },
    "0808/0808.3284_arXiv.txt": {
        "abstract": "We present images and initial results from our extensive {\\it Spitzer} Space Telescope imaging survey of the W5 H\\,{\\sc ii} region with the Infrared Array Camera (IRAC) and Multiband Imaging Photometer for {\\it Spitzer} (MIPS). We detect dense clusters of stars, centered on the O stars: HD 18326, BD +60 586, HD 17505 and HD 17520. At 24 $\\micron$, substantial extended emission is visible, presumably from heated dust grains that survive in the strongly ionizing environment of the H\\,{\\sc ii} region. With photometry of more than 18000 point sources, we analyze the clustering properties of objects classified as young stars by their IR spectral energy distributions (a total of 2064 sources) across the region using a minimal-spanning-tree algorithm. We find $\\sim$40--70\\% of infrared excess sources belong to clusters with $\\geq$10 members. We find that within the evacuated cavities of the H\\,{\\sc ii} regions that make up W5, the ratio of Class II to Class I sources is $\\sim$7 times higher than for objects coincident with molecular gas as traced by $^{12}$CO emission and near-IR extinction maps. We attribute this contrast to an age difference between the two locations, and postulate that at least two distinct generations of star formation are visible across W5. Our preliminary analysis shows that triggering is a plausible mechanism to explain the multiple generations of star formation in W5, and merits further investigation. ",
        "introduction": "Star formation is a self-regulating process---once massive stars form they immediately begin to disrupt their natal environment with their stellar winds and the emission of ionizing radiation. Eventually their parental molecular clouds are destroyed, halting further star formation. However, it has also been argued that the energy input by these massive stars can promote and induce subsequent star formation in the surrounding molecular gas before it disperses, above that which would be produced without external forcing. This process is given the name `triggering' \\citep[see][for reviews of this subject]{elmegreen98,zinnecker07}. It is vital to understand the balance of cloud destruction and triggered star formation if we are to develop a theory that explains the morphology and evolution of star forming regions, star formation efficiencies and the initial mass function of stars in clusters. It also has relevance for star formation on galactic scales in understanding how star formation progresses with time, and indeed how spiral structure of galaxies evolves with time \\citep{seiden90,jungwiert94}. Feedback in star formation can be clearly observed in bright-rimmed clouds, where an edge-on molecular cloud is externally illuminated by nearby young massive stars, creating a cross-section of the photo-evaporation process as the ionization fronts they produce advance into the molecular cloud.  Two triggering mechanisms are of interest in regard to W5. The first is the creation of an ionized H\\,{\\sc ii} region bubble by an initial generation of massive stars within their parental molecular cloud, and its subsequent expansion. In this `collect and collapse' mechanism, investigated analytically by \\citet{whitworth94} and numerically by \\citet{dale07}, the expansion creates a shock front that sweeps up neutral material ahead of it as it progresses outward. The gas accumulated eventually exceeds a critical threshold for collapse and gives rise to a second generation of star formation. Secondly, inhomogeneity in the ISM often leads to small clumps of material remaining exposed inside an H\\,{\\sc ii} region. In this environment, the high pressure of the surrounding ionized gas has been suggested as a mechanism to compress the clumps and form stars \\citep{stutzki88}.  The star forming region W5 is a part of the chain of molecular clouds W3/4/5 \\citep{westerhout58}. Its distance has not been definitively established---\\citet{becker71} found a photometric distance of 2.2 kpc, but more recently, \\citet{hillwig06} found that a distance of 1.9 kpc gave the most consistent results of evolutionary model fits to stellar radii in W5. \\citet{xu06} found a distance to the neighboring region W3OH of 1.95$\\pm$0.04 kpc using maser parallaxes---we adopt a distance to W5 of 2 kpc as a conservative intermediate between these different estimates. W5 is a relatively isolated star forming region, with an apparently simple morphology. Optical imaging shows it is made up of two roughly circular adjoining H\\,{\\sc ii} regions W5 East and W5 West \\citep{karr03}, containing one and four O stars respectively. At least two of the four O stars in W5 West are multiple systems \\citep{hillwig06}. Both $^{12}$CO and 21 cm radio emission \\citep[see][]{normandeau96} demonstrate the same overall shape---$^{12}$CO emission traces the molecular hydrogen gas in W5, which appears as two broken rings of emission. Figure 1 shows a map of $^{12}$CO ($\\lambda$ = 2.6 mm) emission integrated over the range $-$28 to $-$47 km s$^{-1}$, made using observations with the 14m FCRAO telescope. The bulk of the emission is found between a velocity of $-$31 and $-$47 km s$^{-1}$, and thus is likely roughly in the same plane. In this image, following the naming scheme of \\citet{wilking84}, to the northwest of HD 17505 is the molecular cloud W5NW, in between HD 18326 and BD +60 586 is W5NE, and to the east of HD 18326 is W5A. Due to this relatively simple morphology, W5 presents a useful test case for investigating models of triggered star formation and the influence of massive star formation on its surroundings.  \\begin{figure} \\begin{center} \\includegraphics[width=4.5in]{f1.eps} \\caption{$^{12}$CO (1-0) greyscale image of W5 from FCRAO telescope (integrated T$_A^*$ corrected for main beam efficiency $\\eta_{MB}=0.45$, ranging from 0--110 K km s$^{-1}$), oriented North-up and East-left. We label molecular clouds in black following the naming scheme of \\citet{wilking84}. We label the two H\\,{\\sc ii} regions in white. Prominent O stars are marked with white asterisks: HD 18326 in W5 East, and in W5 West from North-East to South-West are BD +60 586, HD 237019 and HD 17505 (upper of pair), HD 17520 (lower of pair).} \\end{center} \\end{figure}  In this paper we present initial results from our mid-IR {\\it Spitzer} survey of W5 with the IRAC and MIPS instruments. In $\\S$ 2 we describe our observations and data reduction techniques and our classification and clustering analysis in $\\S$ 3. In $\\S$ 4 we perform a trial, simplified investigation of triggered star formation models as a means to explain the observed distribution of young stars in different evolutionary states across the region. In $\\S$ 5 we present our conclusions and directions of future work. ",
        "conclusions": ""
    },
    "0808/0808.4001_arXiv.txt": {
        "abstract": "We determine the mass of the black hole at the center of the spiral galaxy NGC 4258 by constructing axisymmetric dynamical models of the galaxy. These models are constrained by high spatial resolution imaging and long-slit spectroscopy of the nuclear region obtained with the {\\em Hubble Space Telescope}, complemented by ground-based observations extending to larger radii. Our best mass estimate is $\\MBH = (3.3 \\pm 0.2) \\times 10^7 \\MSun $ for a distance of 7.28 Mpc (statistical errors only).  This is within 15\\% of $ (3.82\\pm 0.01) \\times 10^7 \\MSun$, the mass determined from the kinematics of water masers (rescaled to the same distance) assuming they are in Keplerian rotation in a warped disk.  The construction of accurate dynamical models of NGC 4258 is somewhat compromised by an unresolved active nucleus and color gradients, the latter caused by variations in the stellar population and/or obscuring dust. Depending on how these effects are treated, as well as on assumptions about the ellipticity and inclination of the galaxy, we obtain black hole masses ranging from $2.4 \\times 10^7 \\MSun$ to $3.6\\times10^7 \\MSun$. This spread is mainly due to uncertainties in the stellar mass profile inside the central 2\\arcsec\\ ($\\sim$70 pc). Obscuration of high-velocity stars by circumnuclear dust (possibly associated with the masing disk) could lead to an underestimate of the black hole mass which is hard to correct. These problems are not present in the $\\sim 30$ other black hole mass determinations from stellar dynamics that have been published by us and other groups; thus, the relatively close agreement between the stellar dynamical mass and the maser mass in NGC 4258 enhances our confidence in the black hole masses determined in other galaxies from stellar dynamics using similar methods and data of comparable quality. ",
        "introduction": "The characterization of massive dark objects (hereafter ``black holes'') at the centers of nearby galaxies has advanced considerably in the past decade.  Nonetheless the dynamical modeling of both stellar and gas kinematics remains subject to uncertainties. These can arise from, among other things, systematic uncertainties and noise amplification in the deprojection of the observables from two to three dimensions, the presence of dust that can affect both the mass-to-light ratio ($M/L$) and the observed kinematics, variations in $M/L$ due to stellar population gradients, and the presence of nonstellar point and/or continuum sources in the central galactic regions which may swamp the light from the surrounding stars or gas. Gas kinematics can also be affected by nongravitational forces, random velocities and pressure support, which are difficult to model accurately, while the complex phase-space structure of stellar orbits can make it difficult to properly sample initial conditions and the phase-space density distribution of the orbits in (stellar) galactic equilibria. Reverberation mapping presents a third possibility for determining black hole masses but comes with its own set of limitations \\citep[e.g.,][]{BlaMcK82,Pet93,NetPet97,GebEtal00b}.  It is therefore desirable to test the reliability of these methods. Unfortunately, comparing stellar dynamical estimates to gas dynamical estimates of black hole masses is not usually possible, as few galaxies have properties that permit more than one method to be applied with confidence.  Even in those cases where it is possible, a detailed comparison is not always illuminating due to uncertainties in the measurements.  One case where gas dynamical and stellar dynamical mass measurements are available for the same galaxy is not reassuring: the gas dynamical mass for the black hole in IC 1459 is a factor of $\\sim 5$ smaller than the stellar dynamical mass \\citep{VerEtal00,CapEtal02}. Better agreement, within a factor of $\\sim 2$, was found between the mass estimate of the black hole in Centaurus A by \\cite{SilEtal05} from stellar dynamics, and by \\cite{MarEtal06} from gas dynamics. Even in the case of the supermassive black hole in the nucleus of our own Milky Way Galaxy, the mass estimates from individual stellar orbits \\citep{GheEtal05} and from a statistical analysis of radial velocities and proper motions \\citep{ChaSah01} differ by about a factor of two.  NGC 4258 presents a unique opportunity to test the reliability of black hole mass estimators. \\citet{MiyEtal95} and \\citet{HerEtal99} used the Very Long Baseline Array (VLBA) to observe water maser emission in a thin, warped, and nearly edge-on gas annulus located at 0.16--0.28 pc from the center of this galaxy, and determined that a central black hole with a mass $\\MBH = (3.9\\pm0.1) \\times 10^7 \\MSun$ is needed to account for the observed velocity profile, proper motions, and accelerations. This black hole mass is thought to be quite reliable for several reasons. (i) The maser emission lines are both strong and narrow, lending themselves to VLBA observations of much higher resolution than possible with the {\\slshape Hubble Space Telescope (HST)} or ground-based telescopes. (ii) The masers seem to lie in a disk that is very thin (height/radius $\\approx 0.02$), so the corrections to the rotation curve due to pressure or random velocities are negligible. (iii) Most importantly, the maser velocities scale very nearly as $v(r) \\propto r^{-1/2}$ with distance $r$ from the center, so that they exhibit an almost perfect Keplerian motion, which makes the deduction of the black hole mass easy and almost free of assumptions. (iv) Finally, the measurements of the mass and distance from the radial velocities, proper motions, and accelerations of the masers are all consistent. More recently, \\citet{HerEtal05}, using a maximum likelihood analysis of the maser positions and velocities, quantified the deviation from Keplerian motion. A number of different models that can explain this discrepancy yield central masses in the range (3.59--3.88) $\\times 10^7 \\MSun$ rescaled to a distance of 7.28 Mpc. The warped disk model, which the authors consider the most probable, implies $\\MBH = (3.82 \\pm 0.01) \\times 10^7 \\MSun$ at this distance.  In this paper, we estimate the mass of the central black hole of NGC 4258 via modeling of the stellar dynamics of the galaxy. This is accomplished using the orbit superposition tools developed by our collaboration (the ``Nuker'' team) \\citep{GebEtal00a,BowEtal01,GebEtal03,ThoEtal04}. Alternative orbit superposition tools have been developed by \\cite{MarEtal98}, \\cite{CreEtal99}, and \\cite{ValEtal04}. Because the maser black hole mass estimate seemed very accurate, we intended this to be a thorough, end-to-end test of the accuracy of our method, at least for this galaxy. Unfortunately, we encountered a number of problems that hindered our ability to determine unambiguously the mass density profile in the central region of the galaxy. In particular, (i) it was not possible to accurately subtract the light of the unresolved active galactic nucleus (AGN) at the center of the galaxy to recover the underlying stellar luminosity profile; (ii) it was not possible to reliably determine the stellar $M/L$ profile within $\\sim$2\\arcsec\\ of the nucleus due to a color gradient, presumably caused by diffuse dust and/or a spatial variation in the stellar population; and (iii) it was not possible to estimate the extinction caused by a compact dust patch near the nucleus, which may be obscuring light from fast-moving stars very close to the black hole, thus causing a systematic underestimate of the black hole mass.  Despite these difficulties, our ``most trusted'' estimate, $\\MBH = (3.3 \\pm 0.2) \\times 10^7 \\MSun$, is in reasonable agreement with the maser mass estimate. The quoted error margin corresponds to the ``1$\\sigma$'' confidence band, and reflects statistical errors only. Uncertainties in the mass profile discussed in the previous paragraph widen the error margin to $2.4 \\times 10^7 \\MSun < \\MBH < 3.6 \\times 10^7 \\MSun$. This range still does not include the potential effect of obscuring circumnuclear dust, which cannot be easily estimated, or any other systematic effects such as deviations from axisymmetry that cannot be properly treated by our dynamical modeling method. These factors are discussed in detail in the following sections.  \\citet{past07} estimate the mass of the black hole in NGC 4258 using gas kinematics, but their 95\\% confidence intervals span a factor of ten in mass so this is not a strong test of the reliability of black-hole mass estimates.   For this work, we use 7.28 Mpc for the distance to NGC 4258 \\citep{TonEtal01}, which is in excellent agreement with the maser-derived distance of 7.2 Mpc. At that distance, $1\\arcsec \\simeq 35$ pc.  In \\S\\ref{obs} we present the photometric and kinematic datasets, which consist of {\\em HST} WFPC2 and NICMOS imaging and STIS spectroscopy, complemented by ground-based imaging and long-slit spectroscopy that extend data coverage to larger radii. We discuss the large-scale morphology of the galaxy in \\S\\ref{ObsMorph}, mainly as it pertains to the determination of the contribution of the stars to the gravitational potential. In \\S\\ref{VarML} we consider the consequences of the color gradient in the central 2\\arcsec. \\S\\ref{DynMod} presents the dynamical models.  We conclude in \\S\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  The determination of the central black hole mass in NGC 4258 from stellar dynamics proved harder than anticipated. The main source of difficulty (and uncertainty in $\\MBH$) comes from the presence of a $V-J$ color gradient in the central $\\sim$2\\arcsec\\ of the galaxy, which prevents us from determining the stellar mass profile near the center with confidence. The model that we trust most (``base model'' in Table \\ref{ModelTable}) yields $\\MBH = 3.31^{+0.22}_{-0.17} \\times 10^7 \\MSun$ and assumes that the stellar mass is traced by the ``corrected'' $V$-band profile, or \\Vcor\\ (Figure \\ref{CumPhotomCenter}). This is identical to the $V$-band profile except for the removal of a small luminous ``bump'' in the central $\\sim$0\\farcs04, which we attribute to the AGN. Using the uncorrected $V$-band profile yields a very similar $\\MBH = 3.25^{+0.22}_{-0.14} \\times 10^7 \\MSun$ (model P4 in Table \\ref{ModelTable}). This small black hole mass ``deficit,'' compared to the base model, corresponds approximately to the mass of stars in the luminosity ``bump'' and is a consequence of the degeneracy between stellar mass and black hole mass very near the center (\\S\\ref{VarMLla01}; Figure \\ref{MassProfile}).  Using the steeper $J$-band (NIR) profile as the mass tracer (model P1), we obtain $\\MBH = 2.49^{+0.09}_{-0.10}\\times10^7 \\MSun$, a factor of 25\\% lower than in the base case. Although $V$-band light is more likely to be compromised as an indicator of the stellar mass distribution than $J$-band light, due to extinction and metallicity gradients, we have greater confidence in the base model because (i) the $V$-band profile is better determined near the center owing to the higher resolution of the $V$ images, (ii) the AGN is considerably fainter in $V$ than in $J$, making it easier to subtract, and (iii) the minimum $\\chi^2$ for the base model is 280, significantly less than the minimum $\\chi^2$ for the P1 model using $J$-band data (which was 303).  The spread in $\\MBH$ introduced by uncertainties in the deprojection parameters (models D1 and D2) is only of order 10\\%, so it is the uncertainties in the stellar mass profile that dominate the errors.  After all these sources of error are taken into account, our black hole mass determination is $\\sim$15\\% lower than the ``preferred'' maser determination $[(3.82\\pm0.01) \\times 10^7 \\MSun$] and only 7\\% lower than the lowest maser determination $[(3.59\\pm0.01) \\times 10^7 \\MSun]$ (rescaled to our distance for the galaxy). We view this level of agreement (2$\\sigma$) as evidence that stellar dynamical mass determinations using similar methods to those in this paper, and with data of comparable quality, are accurate at the 2$\\sigma$ level or better.  Our previous work has sought to avoid galaxies with dusty centers such as NGC 4258.  In this case obscuration of high-velocity stars by the compact nuclear dust patch (\\S\\ref{RedVIData}), possibly associated with the masing disk, could be responsible for an underestimate of the black hole mass in our work. To these problems should be added possible systematic effects due to deviations from axisymmetry, which cannot be properly treated by our axisymmetric modeling code. We also note that the discrepancy between the mass determination from stellar dynamics and the maser mass determination in NGC 4258 is smaller than the discrepancy between the stellar dynamical mass determination and the mass determination from individual orbits for our own Galaxy, which amounts to about a factor of two \\citep{ChaSah01,GheEtal05}.  Most concerns about black hole mass estimates from stellar dynamics have stressed the likelihood that the masses are overestimated, whereas here the stellar dynamical mass is smaller than the maser mass.  Finally, although we regard the maser mass in NGC 4258 as the gold standard in black hole masses, it is conceivable, however unlikely, that some unrecognized effect has led to an overestimate of the mass determined from maser kinematics."
    },
    "0808/0808.2282_arXiv.txt": {
        "abstract": "Using SDSS Data Release 6, we construct two independent samples of candidate stellar wide binaries selected as i) pairs of unresolved sources with angular separation in the range $3\\arcsec - 16\\arcsec$, ii) common proper motion pairs with $5\\arcsec - 30\\arcsec$ angular separation, and make them publicly available. These samples are dominated by disk stars, and we use them to constrain the shape of the main-sequence photometric parallax relation $M_r(r-i)$, and to study the properties of wide binary systems. We estimate $M_r(r-i)$ by searching for a relation that minimizes the difference between distance moduli of primary and secondary components of wide binary candidates. We model $M_r(r-i)$ by a fourth degree polynomial and determine the coefficients using Markov Chain Monte Carlo fitting, independently for each sample. Both samples yield similar relations, with the largest systematic difference of 0.25 mag for F0 to M5 stars, and a root-mean-square scatter of 0.13 mag. A similar level of agreement is obtained with photometric parallax relations recently proposed by \\citet{jur08}. The measurements show a root-mean-square scatter of $\\sim0.30$ mag around the best fit $M_r(r-i)$ relation, and a mildly non-Gaussian distribution. We attribute this scatter to metallicity effects and additional unresolved multiplicity of wide binary components. Aided by the derived photometric parallax relation, we construct a series of high-quality catalogs of candidate main-sequence binary stars. These range from a sample of $\\sim17,000$ candidates with the probability of each pair to be a physical binary (the ``efficiency'') of $\\sim65\\%$, to a volume-limited sample of $\\sim1,800$ candidates with an efficiency of $\\sim90\\%$. Using these catalogs, we study the distribution of semi-major axes of wide binaries, $a$, in the $2,000 < a < 47,000$~AU range. We find the observations to be well described by the \\\"Opik distribution, $f(a)\\propto 1/a$, for $a<a_{break}$, where $a_{break}$ increases roughly linearly with the height $Z$ above the Galactic plane ($a_{break}\\propto12,300\\,Z{\\rm [kpc]}^{0.7}$~AU). The number of wide binary systems with $100 \\, {\\rm AU} < a < a_{break}$, as a fraction of the total number of stars, decreases from 0.9\\% at $Z=0.5$ kpc to 0.5\\% at $Z=3$ kpc. The probability for a star to be in a wide binary system is independent of its color. Given this color, the companions of red components seem to be drawn randomly from the stellar luminosity function, while blue components have a larger blue-to-red companion ratio than expected from luminosity function. ",
        "introduction": "\\label{introduction}  Binary systems can be roughly divided into close (semi-major axes $a\\la10$ AU) and wide (semi-major axes $a\\ga100$ AU, \\citealt{cha07}) pairs. Close binary systems have long been recognized as useful tools for studies of stellar properties. For example, the stellar parameters such as the masses and radii of individual stars are readily determined to high confidence using eclipsing binaries \\citep{and91}. Wide binary systems have proven to be a tool for studies of star formation processes, as well as an exceptionally useful tracer of local potential and tidal fields through which they traverse. Specifically, they were used to place the constraints on the nature of halo dark matter \\citep{ycg04} and to explore the dynamical history of the Galaxy \\citep{aph07}. A further comprehensive list of current applications of wide binaries can be found in \\citet{cha07}.  Close binaries, owing to their relatively short orbital periods and equally short timescales of brightness or spectrum fluctuations, are fairly easy to detect. Unambiguous identification of wide binary systems, on the other hand, requires accurate astrometry on much longer timescales, as these systems have orbital periods $\\ga10,000$ years. However, instead of requiring unambiguous identification, large samples of {\\em candidate} wide binaries can be selected by simply assuming that pairs of stars with small angular separation are also gravitationally bound \\citep{bs81,gou95}, or by searching for common proper motion pairs \\citep{luy79,pov94,aph00,gs03,cg04,lb07}. The angular separation method is simple to apply, but it also introduces a relatively large number of false candidates due to chance association of nearby pairs. The contamination by random associations can be reduced by imposing constraints, such as the common proper motion, or by requiring that the stars are at similar distances. The distances can be inferred through a variety of means, one of which is the use of an appropriate photometric parallax\\footnote{Also known as ``color-luminosity relation''.} relation.  The photometric parallax relation provides the absolute magnitude of a star given that star's color and metallicity. There are a number of proposed photometric parallax relations for main sequence stars in the literature that differ in the methodology used to derive them, photometric systems, and the absolute magnitude and metallicity range in which they are applicable. Not all of them are mutually consistent, and most exhibit significant intrinsic scatter of order a half a magnitude or more (see Figure~2 in \\citealt[hereafter J08]{jur08}).  Instead of using an existing relation to select wide binaries, we propose a novel method that {\\em simultaneously} derives the photometric parallax relation {\\em and} selects a sample of wide binary candidates. The method relies on the fact that components of a physical binary have equal distance moduli ($m_1 - M_1 = m_2 - M_1$) and therefore $\\delta \\equiv \\Delta M - \\Delta m \\equiv (M_2 - M_1) - (m_2 - m_1)  = 0$. Assuming that both stars are on the main sequence, and the {\\em shape} of the adopted photometric parallax relation is correct, the difference in absolute magnitudes $\\Delta M=M_2-M_1$ calculated from the parallax relation must equal the measured difference of apparent magnitudes, $\\Delta m=m_2-m_1$. The $\\Delta M = \\Delta m$ equality for binaries must be valid irrespective of color, and therefore represents a test of the validity of the adopted photometric parallax relation or, alternatively, a way to estimate the parallax relation.  In practice, the distribution of $\\delta$ will not be a delta-function both due to instrumental (finite photometric precision) and physical effects (true vs.~apparent pairs). However, for {\\em true} wide binaries, the distribution of $\\delta$ is expected to be narrow, strongly peaked at zero, and the individual $\\delta$ values are expected to be uncorrelated with color. In contrast, the distribution of $\\delta$ values for randomly associated stellar pairs (hereafter random pairs) should be much broader even when the correct photometric parallax relation is adopted, reflecting the different distances of components of projected binary pairs. This dichotomy can be used to assign a probability to each candidate, of whether it is a true physical binary or a result of chance projection on the sky.  The paper is organized as follows. In Section~\\ref{data}, we give an overview of the SDSS imaging data, and describe the selection, completeness and population composition of two initial, independent samples of candidate binaries. In Section~\\ref{method} we describe the photometric parallax estimation method, compare the best-fit photometric parallax relations to the J08 relation, and analyze the scatter in predicted absolute magnitudes. The properties of wide binaries, such as the color and spatial distributions, are analyzed in Section~\\ref{properties}. Finally, the results and their implications for future surveys are discussed in Section~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  We have presented a novel approach to photometric parallax estimation based on samples of candidate wide binaries selected from the Sloan Digital Sky Survey (SDSS) imaging catalog. Our approach uses the fact that binary system's components are at equal distances and estimates the photometric parallax relation for main-sequence stars by minimizing the difference of their distance moduli. While this method is similar to constraints on photometric parallax relation obtained from globular clusters in that it does not require absolute distance estimates, it has the advantage that it extends to redder colors than available for globular clusters observed by the SDSS, and it implicitly accounts for the metallicity effects.  The derived best-fit photometric parallax relations represent metallicity-averaged relations and thus provide an independent confirmation of relations proposed by J08 in their study of the Galactic structure. An important result of this work is our estimate of the expected error distribution for absolute magnitudes determined from photometric parallax relations (a root-mean-square scatter of $\\sim$0.3 mag, see Section~\\ref{scatter}), which is in good agreement with modeling assumptions adopted by J08. The mildly non-Gaussian error distribution is consistent with both the impact of unresolved binary stars, and the variation of photometric parallax relation with metallicity; we are unable to disentangle these two effects.  The best-fit photometric parallax relations enabled the selection of high-efficiency samples of disk wide binaries with $\\sim22,000$ candidates, that include about 14,000 true binary systems (efficiency of $\\sim2/3$). Using the photometric measurements and angular distance of the two components, samples with efficiency exceeding 80\\% can be constructed (see Section~\\ref{samples}). Such samples could be used as a starting point to further increase the selection efficiency with the aid of radial velocity measurements. Spectral observations of systems where the brighter component is an F/G star, for which it is easy to estimate metallicity, could be used to calibrate both spectroscopic and photometric methods for estimating metallicity of cooler K and M dwarfs. Compared to the state-of-the-art catalogs of wide binaries by \\citet{cg04} and \\citet{lb07}, the samples discussed here represent a significant increase in the number of potential binaries, and probe larger distances (to $\\sim4$ kpc). To facilitate further studies of wide binaries, we make the catalog publicly available\\footnote{The catalog can be downloaded from \\url[HREF]{http://www.astro.washington.edu/bsesar/SDSS\\_wide\\_binaries.tar.gz}}.  Using the high-efficiency subsamples, we analyzed their dynamical and physical properties. We find that the spatial distribution of wide binaries follows the distribution of single stars to within a factor of 2, and that the probability for a star to be in a wide binary system is independent of its color. However, given this color, the companions of red components seem to be drawn randomly from the stellar luminosity function, while blue components have a larger blue-to-red companion ratio than expected from luminosity function (see Section~\\ref{color}). These results are consistent with recent results by \\citet{lb07}, and provide strong constraints for the scenarios describing the formation of such systems (e.g., \\citealt{gie06} and references therein; \\citealt{cla07}; \\citealt{hur07}).  We also study the semi-major axis distribution of wide binaries in the $2,000-47,000$ AU range (see Section~\\ref{spatial}). The observed distribution is well described by the \\\"Opik distribution, $f(a)\\propto 1/a$, for $a<a_{break}$, where $a_{break}$ increases roughly linearly with the height above the Galactic plane ($a_{break}\\sim 12,300$ AU at $Z=1$ kpc). Alternatively, the $a_{break}$ correlates with the local number density of stars as $a_{break} \\propto \\rho^{-1/4}$, but we are unable to robustly identify the dominant correlation ($Z$ and $\\rho$ are highly correlated).  The distribution of semi-major axes for wide binaries was also discussed by \\citet{cg04}. They used a sample of wide binaries selected using common proper motion from the rNLTT catalog \\citep{gs03}, and found $f(a)\\propto 1/a^{1.6}$, with no evidence of a turnover at $a \\la 3000$. Their sample extended to larger angular separations than ours, and probed smaller distances. On the other hand, \\citet{pah07} used wide binaries from the same \\citet{cg04} sample, and detected \\\"Opik distribution, $f(a)\\propto 1/a$, for $a<3,000$, consistent with the result of \\citet{cg04}. In a recent study,  L\\'epine \\& Bongiorno searched for faint common proper motion companions of Hipparcos stars and detected a turnover from \\\"Opik distribution to a steeper distribution around $a\\sim3,000$ AU. Their sample also probed much smaller distances than ours. We compare these results in Figure~\\ref{Fig:compare}. As evident, the variation of $a_{break}$ with distance from the Galactic plane detected here (approximately with distance, as shown in Figure~\\ref{Fig:compare}, since stars in our sample are mostly at high galactic latitudes), is consistent with the above results that are based on more local samples. In particular, this comparison of different studies suggests that the flattening of $f(a)$ for small $a$ that ``puzzled'' Chanam\\'e \\& Gould (see their section 4.3) is probably due to a combination of selection effects and the approach of the domain where \\\"Opik distribution is valid in their sample.  The \\\"Opik distribution suggests that the process of star formation produces multiple stars, which evolve towards binaries after ejecting one or more single stars \\citep{pah07}. The departure from the \\\"Opik distribution may be evidence for disruption of wide binaries over long periods of time by passing stars, giant molecular clouds, massive compact halo objects (MACHOs), or disk and Galactic tides \\citep{heg75,wsw87,ycg04}. However, we note that the $a_{break} \\propto \\rho^{-1/4}$ correlation (see Figure~\\ref{logZ_plots}) is outside the expected range discussed by \\citet{ycg04} ($a_{break}\\propto \\rho^{-2/3}$ for close strong encounters, and $a_{break}\\propto \\rho^{-1}$ for weak encounters).  The samples presented here can be further refined and enlarged. First, the SDSS covers only a quarter of the sky. Upcoming next-generation surveys, such as the SkyMapper \\citep{kel07}, the Dark Energy Survey \\citep{fla07}, Pan-STARRS \\citep{kai02} and the Large Synoptic Survey Telescope (\\citealt{ive08b}, LSST hereafter), will enable the construction of such samples over most of the sky. Due to fainter flux limits (especially for the Pan-STARRS and LSST), the samples will probe a larger distance range and will reach the halo-dominated parts of the Galaxy. Furthermore, due to improved photometry and seeing (e.g., for the LSST, by about a factor of two), the selection will be more robust. We scale the 20,000 candidates discussed here, assuming $log(N) = C + 0.4\\,r$, to the LSST depth that enables accurate photometric metallicity ($r<23$; I08a) and predict a minimum sample size of $\\sim$400,000 candidate wide binary systems in 20,000 deg$^2$ of sky. It is likely that the sample would include more than a million systems due to the increase of the stellar counts close to the Galactic plane.  Another important development will come from the Gaia mission \\citep{per01,wil05}, which will provide direct trigonometric distances for stars with $r<20$. With trigonometric distances, accurate photometric parallax relation can be used to provide strong constraints on the incidence and color distribution of unresolved multiple systems. Until then, a radial velocity survey of candidate binaries assembled here could help with pruning the sample from random associations, and with better characterization of various selection effects."
    },
    "0808/0808.2141_arXiv.txt": {
        "abstract": "Using direct simulations, weakly nonlinear theory and nonlinear mean-field theory, it is shown that the quenched velocity field of a saturated nonlinear dynamo can itself act as a kinematic dynamo. The flow is driven by a forcing function that would produce a Roberts flow in the absence of a magnetic field. This result confirms an analogous finding by F.\\ Cattaneo \\& S.\\ M.\\ Tobias (arXiv:0809.1801) for the more complicated case of turbulent convection, suggesting that this may be a common property of nonlinear dynamos; see also the talk given also online at the Kavli Institute for Theoretical Physics (http://online.kitp.ucsb.edu/online/dynamo\\_c08/cattaneo). It is argued that this property can be used to test nonlinear mean-field dynamo theories. ",
        "introduction": "The magnetic fields of many astrophysical bodies displays order on scales large compared with the scale of the turbulent fluid motions that are believed to generate these fields via dynamo action. A leading theory for these types of dynamos is mean-field electrodynamics \\citep{Mof78,KR80}, which predicts the evolution of suitably averaged mean magnetic fields. Central to this theory is the mean electromotive force based on the fluctuations of velocity and magnetic fields. This mean electromotive force is then expressed in terms of the mean magnetic field and its first derivative with coefficients $\\alpha_{ij}$ and $\\eta_{ijk}$. The former represents the $\\alpha$ effect and the latter the turbulent magnetic diffusivity.  Under certain restrictions the coefficients $\\alpha_{ij}$ and $\\eta_{ijk}$ can be calculated using for example the first order smoothing approximation, which means that nonlinearities in the evolution equations for the fluctuations are neglected. Whilst this is a valid approach for small magnetic Reynolds numbers or short correlation times, it is not well justified in the astrophysically interesting case when the magnetic Reynolds number is large and the correlation time comparable with the turnover time. However, in recent years it has become possible to calculate $\\alpha_{ij}$ and $\\eta_{ijk}$ using the so-called test-field method \\citep{S05,S07}. For the purpose of this paper we can consider this method essentially as a ``black box'' whose input is the velocity field and its output are the coefficients $\\alpha_{ij}$ and $\\eta_{ijk}$. This method has been successfully applied to the kinematic case of weak magnetic fields in the presence of homogeneous turbulence either without shear \\citep{Sur_etal08,Betal08b} or with shear \\citep{B05,Betal08}.  More recently, this method has also been applied to the nonlinear case where the velocity field is modified by the Lorentz force associated with the dynamo-generated field \\citep{Betal08c} In that case the test-field method consists still of the same black box, whose input is only the velocity field, but now this velocity field is based on a solution of the full hydromagnetic equations comprising the continuity, momentum, and induction equations. We emphasize that the magnetic field is quite independent of the fields that appear in the test-field method inside the black box.  Our present work is stimulated by an interesting and relevant numerical experiment performed recently by \\cite{CT08}. They considered a solution of the full hydromagnetic equations where the magnetic field is generated by turbulent convective dynamo action and has saturated at a statistically steady value. They then used this velocity field and subjected it to an independent induction equation, which is equal to the original induction equation except that the magnetic field $\\BB$ is now replaced by a passive vector field $\\tilde{\\BB}$, which does not react back on the momentum equation. Surprisingly, they found that $|\\tilde{\\BB}|$ grows exponentially, even though the velocity field is already quenched by the original magnetic field.  One might have expected that, because the velocity is modified such that it produces a statistically steady solution to the original induction equation, $\\tilde{\\BB}$ should decay or also display statistically steady behavior. The argument sounds particularly convincing for time-independent flows because, if a growing $\\tilde{\\BB}$ were to exist, one would expect this alternative field to grow and replace the initial field. This view is supported by recent simulations in which the flow field from a geodynamo simulation in a spherical shell was used as velocity field in kinematic dynamo computations, and no growing $\\tilde{\\BB}$ was found \\citep{Tilgner08}. However, it turns out that this reasoning is not correct in general. One finds counterexamples even within the confines of mean field MHD using analytical tools. The existence of a growing $\\tilde{\\BB}$ thus is not tied to chaotic flows or fluctuating small-scale dynamos.  This finding of \\cite{CT08} is interesting in view of the applicability of the test-field method to the nonlinear case. Of course, the equations used in the test-field method are different from the original induction equation. (The equations used in the test-field method include an inhomogeneous term and the mean electromotive force is subtracted, but they are otherwise similar to the original induction equation.) Given the seemingly unphysical behaviour of the induction equation in the presence of a vector field different from the actual magnetic field, it would be tempting to argue that one should choose test fields whose shape is rather close to that of the actual magnetic field \\citep{CH08}. On the other hand, the $\\alpha_{ij}$ and $\\eta_{ijk}$ tensors should give the correct response to all possible fields, not just the $\\BB$ field that grew out of a particular initial condition, but also the passive $\\tilde{\\BB}$ field that obeys a separate induction equation. It is therefore important to choose a set of test fields that are orthogonal to each other, even if none of the fields are solutions of the induction equation. One goal of this paper is to show that the $\\alpha_{ij}$ and $\\eta_{ijk}$ tensors obtained in this way provide not only interesting diagnostics of the flow, but they are also able to explain the surprising result of \\cite{CT08} in the context of a simpler example. However, let us begin by repeating the numerical experiment of \\cite{CT08} using the simpler case of a Roberts flow. Next, we consider a weakly nonlinear analysis of this problem and turn then to its mean-field description. ",
        "conclusions": "The most fundamental question of dynamo theory beyond kinematic dynamo theory is ``how do magnetic fields saturate?''. In the simplest picture, the velocity field reorganizes in response to the Lorentz force such that all magnetic fields decay except one which has zero growth rate and which is the one we observe. This picture is already questionable for chaotic dynamos. In a chaotic system, nearby initial conditions lead to exponentially separating time evolutions. If one takes a magnetic field $\\BB$ with (on time average) zero growth rate which is the saturated solution of a chaotic dynamo, and solves the kinematic dynamo problem for a passive vector field $\\tilde{\\BB}$ with initial conditions different from $\\BB$, one is prepared to find growing $\\tilde{\\BB}$. Examples for growing $\\tilde{\\BB}$ in chaotic dynamos have been given by \\cite{CT08}.  For a time-independent saturated dynamo, on the other hand, the simple picture seems to be adequate at first. However, we have shown in this paper that growing $\\tilde{\\BB}$ also exist in the time-independent Roberts dynamo. The origin of the growing $\\tilde{\\BB}$ can in this case be understood with the help of weakly nonlinear theory. The growing $\\tilde{\\BB}$ has the same shape as the saturated dynamo field but is translated in space.  What was wrong with the naive intuition invoked above? It was based on a stability argument (the equilibrated magnetic field should be replaced by $\\tilde{\\BB}$ if there is a growing $\\tilde{\\BB}$). However, the linear stability problem for a solution of the full dynamo equations is different from the kinematic dynamo problem for $\\tilde{\\BB}$, because in the latter, the velocity field is fixed. Both problems are closely related eigenvalue problems, but standard mathematical theorems do not provide us with a relation between the spectra of both problems. The numerical computation above gives an example of a stable Roberts dynamo, showing that the linear stability problem for the solution found there has only negative eigenvalues. Solving the same eigenvalue problem with velocity fixed, which is the kinematic problem for $\\tilde{\\BB}$, can very well lead to positive eigenvalues, and indeed, it does. The dynamo is therefore only stable because the velocity field is able to adjust to perturbations in the magnetic field. The magnetic field on its own is unstable."
    },
    "0808/0808.2188_arXiv.txt": {
        "abstract": "High-resolution submillimetre (submm) imaging of the high-redshift radio galaxy (HzRG), 4C\\,60.07, at $z=3.8$, has revealed two dusty components of roughly equal integrated flux. {\\em Spitzer} imaging shows that one of these components (`B') is coincident with an extremely red active galactic nucleus (AGN), offset by $\\sim$4\\,arcsec ($\\sim$30\\,kpc) from the HzRG core. The other submm component (`A') -- resolved by our synthesised beam and devoid of emission at 3.6--8.0-$\\mu$m -- lies between `B' and the HzRG core. Since the radio galaxy was discovered via its extremely young, steep-spectrum radio lobes and the creation of these lobes was likely triggered by the interaction, we argue that we are witnessing an early-stage merger, prior to its eventual equilibrium state. The interaction is between the host galaxy of an actively-fueled black hole (BH), and a gas-rich starburst/AGN (`B') marked by the compact submm component and coincident with broad CO(4--3) emission. The second submm component (`A') is a plume of cold, dusty gas, associated with a narrow ($\\sim$150\\,km\\,s$^{-1}$) CO feature, and may represent a short-lived tidal structure. It has been claimed that HzRGs and submm-selected galaxies (SMGs) differ only in the activity of their AGNs, but such complex submm morphologies are seen only rarely amongst SMGs, which are usually older, more relaxed systems. Our study has important implications: where a galaxy's gas reservoir is not aligned with its central BH, CO may be an unreliable probe of dynamical mass, affecting work on the co-assembly of BHs and host spheroids. Our data support the picture wherein close binary AGN are induced by mergers. They also raise the possibility that some supposedly jet-induced starbursts may have formed co-evally (yet independently of) the radio jets, both triggered by the same interaction.  Finally, we note that the HzRG host would have gone unnoticed without its jets and its companion, so there may be many other unseen BHs at high redshift, lost in the sea of $\\sim5\\times 10^8$ similarly bright IRAC sources -- sufficiently massive to drive a $>$10$^{27}$-W radio source, yet practically invisible unless actively fueled. ",
        "introduction": "Radio galaxies are believed to host actively-fueled, spinning BHs which power their immense radio luminosities and fashion their characteristic double lobes \\citep*{fr74,rees78,bp82,mccarthy93}. High-redshift examples of the phenomenon (HzRGs) have often been identified as ultra-steep-spectrum (USS) emitters in radio surveys\\footnote{USS sources, with $\\alpha<-1$ where $S_{\\nu}\\propto \\nu^{\\alpha}$, were often found to lack identifications on photographic plates \\citep*[e.g.][]{tielens79}.} \\citep[e.g.][]{bornancini07} and are thus selected during their extreme youth \\citep*[$\\le$10\\,Myr --][]{br99}. Nowadays, a wealth of unequivocal evidence also links HzRGs with very massive galaxies in the early Universe. Near- and mid-infrared (-IR) observations show that HzRGs are associated with the most massive stellar populations at any given redshift \\citep*{blr97, seymour07}.  Direct evidence of vast reservoirs of atomic and molecular hydrogen has also been established, via observations of strong H\\,{\\sc i} absorption against luminous, morphologically complex Ly$\\alpha$ halos -- often extending 100--200\\,kpc from the central radio galaxy -- and via detections of CO and dust \\citep*{hm93, papadopoulos00, archibald01, reuland03, reuland04, debreuck03, debreuck05, klamer05, villarmartin06}.  If HzRGs pinpoint the most massive galaxies out to the highest redshifts then we expect these systems to signpost the most massive peaks in the primordial Universe and to evolve into today's brightest galaxies -- cD-type ellipticals\\footnote{Defined as being centrally located in clusters and very much larger and brighter than other galaxies in the cluster, the {\\em c} being analogous to the notation for supergiant stars in stellar spectroscopy \\citep{mms64}.}. In accord with these expectations, the environments of HzRGs are over-dense in a rich variety of galaxy types, including SMGs. The first example, still the most striking, is the field of 4C\\,41.17 at $z=\\rm 3.8$ \\citep{ivison00, greve07}, and \\citet{stevens03} presented low-resolution submm imaging of a further six HzRGs. Some of these displayed evidence of extended ($\\sim$5--20\\,arcsec, $\\sim$30--150\\,kpc) dust emission \\citep[and UV emission --][]{hatch08}, suggesting that starbursts in HzRGs differ from the compact events seen in local ultraluminous IR galaxies (ULIRGs) and, indeed, in the general $z\\sim\\rm 1-3$ SMG population as probed by high-resolution CO and submm--radio continuum imaging \\citep[$\\ls$0.5\\,arcsec, or $\\ls$4\\,kpc --][]{chapman04, tacconi06, tacconi08, bi08, younger08b}, i.e.\\ that the mechanism for the formation of these very massive galaxies may be fundamentally different to that of $\\gs$L$_{\\star}$ galaxies forming from SMGs \\citep{smail04}.  Probing the size and morphology of the rest-frame far-IR emission in HzRGs is thus a key piece in the puzzle of galaxy formation. With the $\\sim$14-arcsec ({\\sc fwhm}) spatial resolution usually available, we have been unable to distinguish between entirely different modes for the formation of the most massive galaxies. One could imagine HzRG host galaxies assembling via a massive, widespread burst with a low overall star-formation density, or within multiple gas-rich sub-components destined to merge. Evidence from mm interferometric imaging has been enlightening, though the need for imaging on a range of spatial scales is often evident: \\citet{debreuck05} found evidence for two gas components in the 4C\\,41.17 system whilst \\citet{papadopoulos00} argued for a galaxy-wide starburst in another HzRG, 4C\\,60.07, with gas being consumed on a scale of $\\sim$30\\,kpc, and for a gas reservoir more commensurate with local ULIRGs in another, 8C\\,1909+722.  4C\\,60.07 is a powerful double-lobed FR\\,{\\sc ii} radio galaxy -- one of the brightest of the HzRGs observed by \\citet{archibald01} and \\citet{reuland04} in the submm waveband. The radio spectrum of 4C\\,60.07 is unremarkable at low frequencies, with a spectral index, $\\alpha=-0.9$ between 38 and 178\\,MHz (where $S_{\\nu} \\propto \\nu^{\\alpha}$), but its spectrum steepens at higher frequencies, reaching $\\alpha=-1.6$ between 4.9 and 15\\,GHz, hence its selection as an USS radio emitter and its eventual identification with an $R_{\\rm Vega}\\sim 23.2$ galaxy at $z=\\rm 3.7887\\pm 0.0007$ \\citep{chambers96, rottgering97, reuland07}.  In this paper we present submm interferometry of 4C\\,60.07, utilising the Submillimetre Array\\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.} \\citep[SMA --][]{ho04} at a wavelength of 890\\,$\\mu$m with a spatial resolution of $\\sim$2\\,arcsec and high astrometric precision -- the first high-resolution submm continuum image of an unmistakably massive galaxy at $z\\sim\\rm 4$, sensitive to dust on the scales predicted by existing single-dish imaging. The primary goal was to determine the size and nature of the emission. Is it diffuse, or is it due to multiple discrete, compact objects whose total flux density can account for the high $S_{\\rm 850\\mu m}$ seen by SCUBA?  In the next section we describe an extensive set of observations and the associated analysis. Presentation of the reduced images follows in \\S\\ref{discussion}, with our interpretation of those data and discussion of their implications in \\S\\ref{model} and \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}   New, high-resolution submm imaging of the HzRG, 4C\\,60.07, at $z=3.8$, reveals two clumps of emission separated by 3.3\\,arcsec ($\\sim$25\\,kpc). Although one clump is resolved by our $\\sim$2-arcsec beam, both are relatively compact ($\\ls$15\\,kpc) and there is little evidence for a galaxy-wide starburst on larger scales from our observations.  Previously thought to comprise two gas-rich starbursts, one centred on the radio galaxy, we discover an extra level of complexity due to the misalignment of the submm emission and the radio core.  Given the extreme youth of HzRGs, we interpret the 4C\\,60.07 system as an early-stage merger, with its radio jets and a nearby gas-rich starburst/AGN both induced by the interaction. We interpret a second clump of dust, resolved by our submm imaging and coincident with very narrow CO emission, as a plume of cold, dusty gas in a tidal stream, seen in a non-equilibrium state. It is separated spatially from the radio galaxy, which may have exhausted its supply of fuel.  Recent work on the $z=4.7$ quasar BR\\,1202$-$0725 \\citep{feain04}, on LH\\,850.7 \\citep{ivison02} and on Minkowski's Object \\citep{croft06} has suggested that some starbursts are induced by the interaction of a radio galaxy jet with a large gas reservoir. The story we are seeing unfold in the 4C\\,60.07 system raises the possibility that some of these jets may have been triggered by the interaction that led to the starburst, rather than the starburst being triggered by the jet.  When a BH is spatially offset from the gas and dust of its host system, as seen in 4C\\,60.07, this has implications for the calculation of its dynamical mass (via the extent of its CO, spatially and in velocity) and related claims about the co-assembly of BHs and host spheroids at high redshift \\citep[e.g.][]{walter03, walter04}, regardless of whether the system in question is a disk or a merger.  Our observations of 4C\\,60.07 bring to mind scaled-up versions of the NGC\\,4038/4039 (`Antennae' -- \\citealt{wang04}) and VV\\,114 \\citep{frayer99} systems -- early-stage mergers where the bulk of the star formation occurs in an obscured `overlap region' between two gas-rich nuclei. In NGC\\,4038/4039 this region displays a heavily enhanced star-formation rate yet contains little stellar mass \\citep[$\\sim$10 per cent of the total --][]{wang04}. In 4C\\,60.07, the galaxy nuclei correspond to the radio core and `B', and the overlap region is `A'. There are also similarities between 4C\\,60.07 and HE\\,0450$-$2958, one of the objects originally thought to have been caught during its transition from a ULIRG into an optically bright quasar \\citep{hn88, sanders88, cs01} but now believed to be a gas-rich spiral in collision with -- and stimulating optical activity in -- an unobscured quasar hosted by a gas-poor elliptical \\citep{papadopoulos08}.  At $z\\gs3$, where quasars are becoming rarer, a binary quasar with a small separation ($<$1\\,arcmin) has yet to be found \\citep[at $2<z<3$, tens are known --][]{mortlock99, kochanek99, hennawi06}. The 4C\\,60.07 system contains two AGN and, although not strictly speaking a `binary quasar', it does support the argument made by \\citet{djorgovski91} that the high number of binary quasars on small angular separations -- orders of magnitude more than that predicted by an extrapolation of the quasar correlation function power law -- is due to the triggering of AGN during mergers \\citep[see also][]{smail03, alexander03, iono06}.  We would not have been aware of the BH that presumably drives the immense radio luminosity of 4C\\,60.07 were it not for the interaction that triggered both the FR\\,{\\sc ii} radio activity and the nearby starburst. Is there an unseen population of BHs at high redshift -- invisible unless actively fueled, yet sufficiently massive to drive a $>$10$^{27}$-W radio source? Aside from its rare radio properties, the SED of 4C\\,60.07's host galaxy is not particularly unusual, and such objects would be impossible to sift from the $\\sim5\\times 10^8$ similarly bright IRAC sources in the sky.  \\citet{reuland07} claim that HzRGs and SMGs differ only in the activity of their AGNs.  Why, then, are complex submm morphologies the exception rather than the rule in typical SMGs? The answer lies in the youth of the 4C\\,60.07 system, as set by its selection as a USS radio emitter. We see 4C\\,60.07 in the first throes of a violent interaction, probably within a few Myr of its latest burst of activity commencing, whereas analysis of the rest-frame optical properties of SMGs show that they are typically $\\ge$10$\\times$ older than HzRGs -- $\\sim$100\\,Myr \\citep{smail04,swinbank06} -- and are therefore more likely to display more relaxed morphologies. This provides a ready explanation for the extended nature of the submm emission seen towards the \\citeauthor{stevens03} sample of HzRGs, which is at odds with the compact emission seen for most SMGs \\citep{younger07, younger08, bi08}.  Since these HzRGs will all be youthful, we expect non-equilibrium morphologies to be commonplace, though not necessarily the complex dual-AGN signature seen in 4C\\,60.07. It is not the activity of their AGNs that mark HzRGs as different from SMGs, it is their youth. The possibility that HzRGs differ from SMGs has many implications. It would be desirable to make similar observations of a larger sample of HzRGs at a resolution similar or better than that employed here.  Recent studies of rest-frame far-IR emission, either directly \\citep{tacconi06, younger07, younger08} or via the far-IR/radio correlation \\citep{chapman04,bi08}, have highlighted the importance of high spatial resolution. Single-dish observations -- with SCUBA2 \\citep{holland06}, {\\em Herschel} \\citep{griffin07} or {\\em SPICA} \\citep{swinyard08} -- will often provide an ambiguous or misleading picture. Our work re-inforces this view: we have spent a decade assuming that the radio core of 4C\\,60.07 dominates the rest-frame far-IR emission when the emission is, in fact, associated with a gas-rich companion. We must await sensitive interferometers such as the Atacama Large Millimetre Array \\citep[ALMA --][]{wootten03} to probe the interactions that drive many aspects of galaxy evolution, and design future missions, such as the {\\em Far-Infrared Interferometer} \\citep[{\\em FIRI} --][]{hi08}, with the ability to probe a large range of spatial scales.  \\vspace*{5mm} \\noindent {\\it Facilities:} {SMA, \\em Spitzer}, VLA, JCMT, WHT."
    },
    "0808/0808.2051_arXiv.txt": {
        "abstract": "Results from the first two years of data from the Taiwanese-American Occultation Survey (TAOS) are presented. Stars have been monitored photometrically at 4~Hz or 5~Hz to search for occultations by small ($\\sim$3~km) Kuiper Belt Objects (KBOs).  No statistically significant events were found, allowing us to present an upper bound to the size distribution of KBOs with diameters 0.5~km~$<D<$~28~km. ",
        "introduction": "\\setcounter{footnote}{0}  The study of the Kuiper Belt has exploded since the discovery of 1992 QB1 by \\citet{1993Natur.362..730J}.  The brightness distribution of objects with $R$ magnitude brighter than $\\sim$26 is relatively well-established by many surveys, most recently by \\citet[and references therein]{2008Icar..195..827F} .  The brightness distribution is adequately described by a simple cumulative luminosity function \\mbox{$\\Sigma(<R) = 10^{\\alpha(R-R_0)}$~deg$^{-2}$}, where $R_0 \\sim 23$ and $\\alpha \\sim 0.6$, for objects with magnitude $R<26$.  There is clear evidence for a break to a shallower slope for fainter objects: the deepest survey, conducted using the Advanced Camera for Surveys on the \\emph{Hubble Space Telescope} \\citep{2004AJ....128.1364B} extended to $R = 28.5$, and found a factor of $\\sim$25 fewer objects than would be expected if the same distribution extended into this range.  The size distribution of Kuiper Belt Objects (KBOs) is believed to reflect a history of \\emph{agglomeration} during the planetary formation epoch, when relative velocities between particles were low and collisions typically resulted in particles sticking together, followed by \\emph{destructive collisions} when the relative velocities were increased by dynamical processes after the giant planets formed \\citep{1996AJ....112.1203S, 1997Icar..125...50D, 1997AJ....114..841S, 1999AJ....118.1101K, 1999ApJ...526..465K, 2004AJ....128.1916K, 2005Icar..173..342P}. The slope of the distribution function for larger objects reflects the early phase of agglomeration, while the shallower distribution for smaller objects reflects a subsequent phase of destructive collisions. The location of the break moves to larger sizes with time, while the distribution for smaller objects is expected to evolve towards a steady state collisional cascade \\citep{2004AJ....128.1916K, 2005Icar..173..342P}. Models for the spectrum of small bodies differ between \\citet{2005Icar..173..342P}, who derived a double power-law distribution, and \\citet{2004AJ....128.1916K}, whose simulations show more structure, depending on material properties.   Thus, the size spectrum encodes information about the history of planet formation and dynamics. However, the size spectrum for small KBOs is not constrained by the imaging surveys because the objects of interest are too faint for direct detection using presently available instruments. These small objects may, however, be detected indirectly when they pass between an observer and a distant star \\citep{1976Natur.259..290B, 1992QJRAS..33...45D, 1992ASPC...34..171A, 1997MNRAS.289..783B, 2000Icar..147..530R, 2003ApJ...587L.125C, 2007AJ....134.1596N}. The challenge confronting any survey exploiting this technique is the combination of very low anticipated event rate and short duration of the events (typically $<$ 1~second).  \\begin{figure*}[bt] \\plottwo{f1a.eps}{f1b.eps} \\caption[]{Demonstration of detrending by the filtering process. The left panel shows a raw lightcurve set $f(t)$, and the right panel shows the lightcurve set $h(t)$ after filtering. The total number points in the lightcurve set is 26,622, corresponding to 89~minutes. Here only one of every 10~points is plotted for readability.} \\label{fig:lcdemo} \\end{figure*}  Other groups are attempting similar occultation surveys. \\citet{2006AJ....132..819R} reported three events in 10~star-hours of photometric data sampled at 45~Hz, which they modeled as objects at 15~AU, 140~AU, and 210~AU, respectively, placing the inferred objects outside the Kuiper Belt.  \\citet{2008AJ....135.1039B} reported results of 5~star-hours of data sampled at 40~Hz, during which no events were detected.  \\citet{2006Natur.442..660C} reported a surprisingly high rate of possible occultation events in \\emph{RXTE} x-ray observations of Sco X-1, but many of these events have since been attributed to instrumental effects \\citep{2008ApJ...677.1241J, 2007MNRAS.378.1287C}.  We report here the first results of the Taiwanese American Occultation Survey (TAOS). TAOS differs from the previously reported projects primarily in the extent of the photometric time series, a total of $1.53 \\times 10^5$~star-hours, and in that data are collected simultaneously with three telescopes. Some compromises have been made in regard to signal-to-noise (SNR), which is typically lower than in previously reported surveys, and in cadence, which is 4~Hz or 5~Hz, in contrast to the higher rates mentioned above.  The substantial increase in exposure more than compensates for the lower cadence and SNR, and we are able to probe significant ranges of the model space for small KBOs.  We have also developed a statistical analysis technique which allows efficient use of the multi-telescope data to detect brief occultation events that would be statistically insignificant if observed with only one telescope. ",
        "conclusions": "We have surveyed the sky for occultations by small KBOs using the three telescope TAOS system. We have demonstrated that a dedicated occultation survey using an array of small telescopes, an innovative statistical analysis of multi-telescope data, and a large number of star-hours, can be used as a powerful probe of small objects in the Kuiper Belt, and we are thus able to place the strongest upper bound to date on the number of KBOs with $0.5~\\mathrm{km} < D < 28$~km. We continue to operate TAOS, soon with an additional telescope, and will report more sensitive survey results in the future."
    },
    "0808/0808.2098_arXiv.txt": {
        "abstract": "The dynamics of expansion of the Universe and evolution of scalar perturbations are discussed for the quintessential scalar fields $Q$ with the classical Lagrangian $L=\\frac{1}{2}Q_{;i}Q^{;i}-U(Q)$ satisfying the additional condition $w=const$ or $c^2_a=0$. Both quintessential fields are studied for the same cosmological model. It is shown that the accelerated expansion of the Universe is caused by the effect of rolling down of the field to minimum. At the early epoch the contribution to dynamics of the quintessence with $w=const$ is negligible (like that of cosmological constant) while quintessence with $c^2_a=0$ mimics dust matter. In future the scalar field with $c^2_a=0$ will mimic cosmological constant.  The systems of evolution equations for gauge-invariant perturbations of metric, matter and quintessence have been analysed analyticaly for the early stage of the Universe life and numerically up to the present epoch. It is shown that amplitudes of the adiabatic matter density perturbations grow similarly in both models (and like in $\\Lambda$CDM-model), but time dependences of different amplitudes of the quintessence perturbations are varied: gauge-invariant variables $D_g^{(Q)}$ and $D_s^{(Q)}$ decay from initial constant value after the particle horizon entry while $D^{(Q)}$ and $V^{(Q)}$ grow at the early stage before the horizon entry and decay after that -- in the quintessence-dominated epoch, when gravitational potential starts to decay -- so, that at the current epoch they are approximately two orders lower than matter ones on the supercluster scales. Therefore, on the subhorizon scales the quintessential scalar fields are smoothed out while the matter clusters.  It is also shown that both quintessential scalar fields suppress the growth of matter density perturbations and the amplitude of gravitational potential. In these QCDM-models -- unlike $\\Lambda$CDM ones -- such suppression is scale dependent and more visible for the quintessence with $c^2_a=0$. ",
        "introduction": "Cosmological observations of the last decade surely assert that the main part of the energy density of the Universe -- more than 70\\% --  belongs to the unknown essence, called ''dark energy''. Its cosmological mission is to provide the accelerated expansion of the Universe, revealed from exploration of SN Ia's in the distant galaxies and temperature fluctuation power spectrum of cosmic microwave background. The cosmological $\\Lambda$CDM-model, based on the Einstein equations with cosmological constant (see \\cite{apunevych2007,spergel2007,komatsu2008} and references therein), describes very well almost whole set of the observational data on dynamics of expansion of the Universe and formation of its large-scale structure. But physical interpretation of the cosmological constant is rather problematic \\cite{peebles1993,sahni2000,carrol2001,peebles2003,bousso2008,shapiro2008}. Therefore alternative approaches  -- new physical fields (classical scalar field -- quintessence, tachyon field, k-essence, phantom field, quintom field), Chaplygin gas, gravity and general relativity modifications, multidimensional gravity, branes and others -- are intensively analysed (see reviews \\cite{sahni2000,carrol2001,padmanabhan2003,peebles2003,copeland2006,bludman2007,frieman2008} and special issue of Gen. Relativ. Gravit., 2008, v.40) now. Up to now none of them has crucial preferability from observational or theoretical point of view. Therefore each of them must be comprehensively studied. Here we restrict ourselves to quintessential scalar fields with classical Lagrangian $L=Q_{;i}Q^{;i}/2-U(Q)$ in the dark energy -- matter dominated Universe.  The quintessence model can be defined by setting of the appropriate potential $U(Q)$ or equation of state (EoS) parameter $w_{Q}\\equiv p_{Q}/c^2\\rho_{Q}$. There is a dozen or more physically-motivated shapes of the potential $U(Q)$: exponential, double exponential, exponential with inverse power, power-law, etc. The dynamics of such scalar fields is intensively studied (see review \\cite{copeland2006}). The EoS parameter of dark energy completely defines the background dynamics as well as the evolution of cosmological perturbations \\cite{hu1998,hu1999,linder2008}. Since observational data on SN Ia magnitude -- redshift relation and cosmic microwave anisotropy give relatively narrow ranges of dark energy density and EoS parameter values, it looks quite attractive to establish the potential $U(Q)$ using these data and analyse the background dynamics and perturbative properties of such scalar field which are not studied widely enough.  In our previous papers we have constructed the potentials of scalar fields with classical and tachyonic Lagrangian leading to the constant EoS parameter $w_{Q}=const$ \\cite{sergijenko2008a} and analysed the background dynamics and perturbative properties of such scalar fields \\cite{sergijenko2008b}. It was shown that cosmological model with cold dark matter and such types of the scalar field ($QCDM$-model) agrees slightly better with the accessible today observable data than the $\\Lambda$CDM-model. But difference of quantitative merits of goodness is not large enough to pick out one of them at confidential level of  $1\\sigma$. Since the degeneracies between model parameters of dark energy and cosmological parameters \\cite{kunz2007,durrer2008,hlozek2008,linder2008,zhdanov2008} exist for the background dynamics, the complete analysis of linear density perturbations in both dark matter and dark energy components is important for improvement of dark energy observational tests. Among the large number of free quintessence parameters and unknown initial values of quintessence perturbation modes there is only small part of models, for which the evolution of perturbations has been studied. The general conclusion is that magnitudes of dark energy density perturbations on scales smaller than horizon are essentially lower than corresponding magnitudes of matter density ones. But character of their evolution depends strongly on the scalar field model (its potential, time variation of EoS parameter, sound speed, etc.), initial conditions, scale of perturbations and gauge (see for example  \\cite{dave2002,doran2003,dedeo2003,bean2004,bartolo2004,liu2004,kulinich2004,dutta2007,unnikrishnan2008}).  Here the special attention should be paid to the EoS parameter of dark energy $w_{Q}$, which can be constant or varying in time. The temporal variation of the dark energy EoS parameter is often presented by linear fitting formula with two \\cite{chevallier2001} or three \\cite{komatsu2008} parameters to be estimated. Other functional dependences of $w_Q$ on scale factor or redshift can be found in \\cite{copeland2006,linder2006,sahni2006}. Here we study the parametrization of the equation of state, which needs only 1 additional quantity with clear physical meaning -- the adiabatic speed of sound $c_a^2\\equiv \\dot{p}_Q/c^2\\dot{\\rho}_Q$ (the analysis of generalized dark sector components can be found in the early works \\cite{hu1998,hu1999}). In general, $c_a^2$ is the unknown function of time. However, taking into account the simplicity we restrict ourselves to $c_a^2=const$, so it is regarded only as the second physical parameter defining the equation of state of dark energy (the first one -- the present value of $w_Q$).  In this paper we undertake the comparative analysis the evolution of gauge-invariant variables of the scalar perturbations in the model with non-relativistic matter ($p_{M}\\ll c^2\\rho_{M}$) and scalar field which we define by classical Lagrangian with potential constructed for two cases ($w_{Q}=const$ and $c^2_a=0$) in the concordance cosmological models. These cases have been chosen because they allow us to obtain analytical solutions, which seems to look very attractive in the world of numerical computations. We assume the adiabatic initial conditions for matter and dark energy scalar perturbations. ",
        "conclusions": "\\begin{figure} \\includegraphics[width=8cm]{gpsi.eps} \\caption{Evolution of ratios $D^{(M)}a_{init}/D^{(M)}_{init}a$ and $\\Psi/\\Psi_{init}$ for linear perturbations with scales $k=0.0001$, $0.001$, $0.01$ and $0.1$ $Mpc^{-1}$ (from top to bottom) in the models with non-relativistic matter and quintessence ($c^2_a=0$ -- solid line, $w=const$ -- dashed line).} \\label{fig7} \\end{figure} The dynamics of expansion of the Universe and evolution of scalar perturbations are studied for the quintessential scalar fields $Q$ with the classical Lagrangian $L=\\frac{1}{2}Q_{;i}Q^{;i}-U(Q)$ satisfying the additional condition $w=const$ or $c^2_a=0$. For both quintessential scalar fields the potential $U(Q)$ and time dependence of $Q$ are constructed for the same cosmological model and it is shown that the accelerated expansion of the Universe is caused by the effect of rolling down of the potential to minimum (Fig.\\ref{fig3}). In QCDM-model with $w=const$ the ratio $\\rho_M/\\rho_Q\\rightarrow\\infty$ when $a\\rightarrow 0$, while in QCDM-model with $c^2_a=0$  $\\rho_M/\\rho_Q\\rightarrow (1-\\Omega_Q)/(1+w_0)\\Omega_Q$ and $w\\rightarrow 0$ when $a\\rightarrow 0$. At the early epoch $w$-quintessence is dynamically unsubstantial like cosmological constant while $c^2_a$-quintessence mimics dust matter ($w\\approx 0$) at $a\\ll 1$ and cosmological constant ($w=-1$) at $a\\gg 1$. The dependence of acceleration parameter on redshift is a bit different for them (Fig.\\ref{fig1}) but close to $\\Lambda$CDM-model one and indistinquishable observationally now.  Asymptotic analysis of the systems of evolutionary equations for gauge-invariant perturbations has shown that adiabatic initial conditions for non-relativistic matter and $w$- and $c^2_a$-quintessence are allowed.  The numerical integration of these systems give time dependences of gauge-invariant variables for matter and quintessence scalar perturbations (Fig.\\ref{fig4}, \\ref{fig5}). The main conclusion deduced from them is following: the magnitudes of the adiabatic matter density perturbations grow like in $\\Lambda$CDM-model, while for quintessence $D_g^{(Q)}$, $D_s^{(Q)}$ are constant and $D^{(Q)}$, $V^{(Q)}$ grow before the particle horizon entry but all variables decay after that in such way, that at the current epoch they are approximately two orders lower than the corresponding quantities for dust matter on supercluster scales. Therefore, on subhorizon scales the quintessential scalar field is smoothed out while the matter is clustered.  The quintessential scalar fields studied here suppress the growth of matter density perturbations and the  magnitude of gravitational potential (Fig.\\ref{fig7}). In these QCDM-models -- unlike $\\Lambda$CDM ones -- such suppression is scale dependent and more visible for $c^2_a$-quintessence. Such features of quintessence are important for calculations of the matter density power spectrum at different redshifts and the power spectrum of CMB temperature fluctuations in the range of scales of the late integrated Sachs-Wolfe effect. That can be used for interpretation of data of current and planned experiments in order to identificate the nature of dark energy."
    },
    "0808/0808.2860_arXiv.txt": {
        "abstract": "We describe three-dimensional general relativistic magnetohydrodynamic simulations of a geometrically thin accretion disk around a non-spinning black hole. The disk has a thickness $h/r\\sim0.05-0.1$ over the radial range $(2-20)GM/c^2$.  In steady state, the specific angular momentum profile of the inflowing magnetized gas deviates by less than 2\\% from that of the standard thin disk model of \\citet{nt73}.  Also, the magnetic torque at the radius of the innermost stable circular orbit (ISCO) is only $\\sim2\\%$ of the inward flux of angular momentum at this radius.  Both results indicate that magnetic coupling across the ISCO is relatively unimportant for geometrically thin disks. ",
        "introduction": "The recent development of general relativistic magnetohydrodynamic (GRMHD) codes (e.g., \\citealt{gmt03,dev03}) has finally allowed realistic numerical simulations of magnetized accretion disks around black holes (BHs).  This has led to a better understanding of the inner regions of these disks (e.g. \\citealt{khh05}) and of their role in launching relativistic jets (e.g. \\citealt{M06}).  One of the interesting new results is the recognition that magnetic fields alter the structure of the accretion flow near and inside the innermost stable circular orbit (ISCO).  This may result in large deviations from the traditional picture of a vanishing torque at the ISCO \\citep{krolik99,gammie99,khh05}.  \\citet{pac00} and \\citet{ap03} suggested that the zero-torque condition is likely to be a good approximation for geometrically thin disks.  This was confirmed by \\citet{snm08} who, using a global height-integrated model, showed that modifications to the stress profile are negligibly small for disk thicknesses $h$ less than about a tenth of the local radius $r$.  Their work was, however, based on a hydrodynamic model with $\\alpha$-viscosity and did not explicitly include magnetic fields.  Numerical MHD simulations by \\citet{kh02} using a pseudo-Newtonian potential, and by \\citet{khh05} using a GRMHD code, indicated that magnetic torques are indeed important inside the ISCO.  In particular, these authors found no evidence for a ``stress edge,'' leading them to argue that simple disk models based on the zero-torque condition (e.g., \\citealt{ss73}; \\citealt{nt73}, hereafter NT73) may be seriously wrong. If true this would undermine recent efforts to estimate of the spin parameters of BHs using the NT73 model \\citep{shafee06, mcc06, ddb06,liu08}.  The MHD simulations carried out so far have considered non-radiating accretion flows that are geometrically rather thick.  The one exception is the recent work of \\citet{rb08} who considered a thin disk with $h/r\\sim 0.05$ in a pseudo-Newtonian potential.  In this paper, we describe global 3D GRMHD simulations of a thin disk ($h/r\\sim0.05$) around a non-spinning BH and compare our simulated model with the NT73 model. ",
        "conclusions": "For a geometrically thin accretion disk with $h/r \\sim0.05-0.1$ (Fig. 1) around a non-spinning BH, we find that the specific angular momentum profile $\\bar{\\ell}_{\\rm in}(r)$ of the accreting magnetized gas is quite similar to that assumed in the idealized model of NT73 (Fig. 2). Specifically, $\\bar{\\ell}_{\\rm in}$ is slightly smaller than, but nearly equal to, the Keplerian value of $\\bar{\\ell}$ for radii outside the ISCO, and is almost independent of $r$ inside the ISCO.  At the BH horizon, the value of $\\bar{\\ell}_{\\rm in}$ deviates from the NT73 value by only 2.0\\%.  This result suggests that the NT73 model is a good approximation for thin disks.  However, fast magnetosonic waves from radii as small as $5.1M$ can escape the plunging region and communicate magnetic stresses between the disk and the plunging regions, which means that the zero-torque condition assumed in the NT model is not perfectly valid.  Nevertheless, the normalized outward flux of angular momentum $\\bar{\\ell}_{\\rm out}$ at the ISCO is only $\\sim2.2\\%$ of the inward flux $\\bar{\\ell}_{in}$.  Thus, the magnetic coupling across the ISCO is quantitatively weak.  As a result, we estimate the luminosity of the disk to be enhanced by no more than a few percent.  We find that turbulent activity is pronounced at radii beyond about $10M$, and that the flow is nearly laminar inside the ISCO.  This suggests that the bulk of the dissipation occurs outside the ISCO in the disk proper, not in the plunging region. In our next paper we plan to discuss the energy dissipation profile of 3D GRMHD simulations, which is ultimately what determines the observed radiation.  Future studies should also investigate the dependence of these results on $h/r$, magnetic field geometry, black hole spin (e.g. \\citealt{gsm04}), cooling, resistivity and viscosity, and mass distribution.  The results reported here are in qualitative agreement with those obtained by \\citet{rb08}.  They carried out a 3D non-relativistic MHD simulation in a pseudo-Newtonian potential, whereas ours is a fully relativistic GRMHD simulation.  Also, their initial torus had its inner radius at the ISCO, $r_{\\rm in}=6M$, whereas our torus initially had $r_{\\rm in}=20M$.  This might explain some differences in our results, e.g., their turbulence edge is at $5.5M$ whereas ours is at $6.8M$ (compare their Fig.~3 with our Fig.~4).  Indeed, we find that our results are sensitive to the initial conditions if we place the inner edge of the torus much inside $r_{\\rm in}\\sim 20M$.  Despite these differences the two simulations agree on the main result, viz., for a geometrically thin disk, the ISCO does behave like a physical boundary separating the disk proper from the plunging region. \\citet{beckwith08b} suggest a larger deviation from NT73, but they considered a torus with a larger value of $h/r$ and did not use a self-consistent dissipation model. A study of the actual dissipation in the simulation, the goal of our next paper, will hopefully resolve these differences."
    },
    "0808/0808.3111_arXiv.txt": {
        "abstract": "Recent observations of carbon, sulphur, and titanium isotopes at redshifts z$\\sim$1 and in the local stellar disc and halo have opened a new window into the study of isotopic abundance patterns and the origin of the chemical elements.  Using our Galactic chemical evolution code {\\tt GEtool}, we have examined the evolution of these isotopes within the framework of a Milky Way-like system. We have three aims in this work: first, to test the claim that novae are required, in order to explain the carbon isotope patterns in the Milky Way; second, to test the claim that sulphur isotope patterns at high-redshift require an initial mass function biased towards massive stars; and third, to test extant chemical evolution models against new observations of titanium isotopes that suggest an anti-correlation between trace-to-dominant isotopes with metallicity. Based upon our dual-infall galactic chemical evolution modelling of a Milky Way-like system, and the subsequent comparison with these new and unique datasets, we conclude the following: novae are not required to understand the evolution of $^{12}$C/$^{13}$C in the solar neighbourhood; a massive star-biased initial mass function is consistent with the low ratios of $^{12}$C/$^{13}$C and $^{32}$S/$^{34}$S seen in one high-redshift late-type spiral, but the consequent super-solar metallicity prediction for the interstellar medium in this system seems highly unlikely; and deficient isotopes of titanium are predicted to correlate positively with metallicity, in apparent disagreement with the new datasets; if confirmed, classical chemical evolution models of the Milky Way (and the associated supernovae nucleosynthetic yields) may need a substantial overhaul to be made consistent. ",
        "introduction": "Galactic chemical evolution models are employed to study the spatial and temporal evolution of elements and isotopes throughout the Universe. When coupled to phenomenological representations of galaxy assembly, such models can be compared directly with exquisite elemental and isotopic abundance patterns observed locally in the Milky Way. From these comparisons, conclusions can be drawn regarding the veracity of the underlying micro-physics governing stellar evolution and nucleosynthesis, in addition to the macro-physics governing the assembly of galaxies, the redistribution of the interstellar medium (ISM) over galactic timescales, and the relative birthrate of stars of various masses (the so-called initial mass function, or IMF).  While most galactic chemical evolution models to date have concentrated on predictions related to the total elemental abundance patterns, in unique circumstances, the availability of detailed isotopic patterns can enhance the predictive power of these models; such isotopic patterns afford additional leverage for discriminating between the various origin sites for the chemical elements.  A full literature review of the field would be unwieldy, but we refer the reader to any number of comprehensive reviews and the many important references therein - e.g. Timmes, Woosley \\& Weaver (1995, hereafter TWW95), Prantzos, Aubert \\& Audouze (1996), Fenner \\& Gibson (2003), Romano \\& Matteucci (2003, hereafter, RM03), and Chiappini et~al. (2008).  Recent observational work has opened a new window into the study of isotopic abundance patterns - specifically, the identification (for the first time) of carbon and sulphur isotopes at redshifts z$\\sim$1 (Muller et~al. 2006, hereafter M06; Levshakov et~al. 2006), coupled with the recent determination of titanium isotopic abundances in the local stellar disc and halo (Chavez 2008). The abundances of each of these isotopes, and their evolution with redshift, hold clues as to the relative importance of supernovae versus asymptotic giant branch stars versus novae in seeding the Universe with these important elements.  In this work, we present predictions for isotopic ratios of carbon, sulphur, and titanium within the framework of a classical Milky-Way-like disc galaxy model. The paper is organised as follows: the fundamental nucleosynthetic origins of the relevant carbon, sulphur, and titanium isotopes are first (briefly) reviewed in \\S~2; in \\S~3, we introduce the chemical evolution code ({\\tt GEtool}) and the four stellar yield compilations employed in our modelling; finally, the results are presented and summarised in \\S~4 and \\S~5, respectively. ",
        "conclusions": "We have explored the evolution of carbon, sulphur, and titanium isotopes in both the local- (Milky Way) and high-redshift Universe, using a suite of chemical evolution models generated with {\\tt GEtool}.  We examined the need for a novae contribution to explain the carbon isotope patterns in the Milky Way, the evidence for a massive star-biased IMF at high-redshift based upon sulphur isotope ratios, and the impact of new observations of titanium isotope patterns in nearby stars upon our picture of galactic chemical evolution.  We have found:   \\begin{itemize} \\item In contrast to earlier studies, the necessity for a significant contribution of $^{13}$C from novae is ameliorated when employing contemporary models of asymptotic giant branch stars.\\footnote{It should be noted that novae, or some additional source, \\it will \\rm be required though to explain isotopic ratios such as $^{14}$N/$^{15}$N and $^{16}$O/$^{18}$O; the KL07 yields do not assist in this regard.} \\item A massive star-biased IMF at high-redshift results in a significant decrease in the predicted $^{12}$C/$^{13}$C and $^{32}$S/$^{34}$S, consistent with those observed, but at the expense of predicting highly super-solar metallicity in otherwise normal-looking spiral galaxies. \\item Earlier studies which suggested that $^{50}$Ti was significantly underproduced were, in part, led to this conclusion by the adoption of older monolithic-style collapse models of galactic evolution.  Our dual-infall model eliminates these apparent problems, although it remains true that $^{47}$Ti is problematic (i.e., underproduced). \\item Our chemical evolution models predict a positive correlation between trace-to-dominant titanium isotope ratios and metallicity, while the data is more suggestive of a lack of correlation. \\end{itemize}"
    },
    "0808/0808.3327_arXiv.txt": {
        "abstract": "We report on the discovery of a Cepheid population in the Sculptor Group spiral galaxy NGC 247 for the first time. On the basis of wide-field images collected in photometric surveys in V and I bands which were conducted with three different telescopes and cameras, 23 Cepheid variables were discovered with periods ranging from 17 to 131 days. We have constructed the period-luminosity relations from these data and obtain distance moduli to NGC 247 of 28.20 $\\pm$ 0.05 mag (internal error) in V, 28.04 $\\pm$ 0.06 mag in I, and 27.80 $\\pm$ 0.09 mag in the reddening-independent Wesenheit index. From our optical data we have determined the total mean reddening of the Cepheids in NGC 247 as E(B-V)=0.13 mag, which brings the true distance modulus determinations from the V and I bands into excellent agreement with the distance determination in the Wesenheit index. The best estimate for the true distance modulus of NGC 247 from our optical Cepheid photometry is 27.80 $\\pm$0.09 (internal error) $\\pm$ 0.09 mag (systematic error) which is in excellent agreement with other recent distance determinations for NGC 247 from the Tip of the Red Giant branch method, and from the Tully-Fisher relation. The distance for NGC 247 places this galaxy at twice the distance of two other Sculptor Group galaxies, NGC 300 and NGC 55, yielding supporting evidence for the filament-like structure of this group of galaxies. The reported distance value is tied to an assumed LMC distance modulus of 18.50 mag. \\\\ ",
        "introduction": "In order to improve the extragalactic distance scale using stellar standard candles observable in nearby galaxies, the Araucaria Project\\footnote{http://ezzelino.ifa.hawaii.edu/$\\sim$bresolin/Araucaria/index.html} (Gieren et al. 2001) has been investigating the most promising stellar distance indicators: Cepheid variables, RR Lyr{\\ae} stars, red clump stars, blue supergiants  and the tip of the red giant branch (TRGB). The main  goal of our project is to obtain an accurate calibration, better than 5\\%, of these primary distance indicators as a function of the environmental properties of the host galaxy: metallicity, age and star formation history. To achieve this goal, the Araucaria Project studies galaxies of different morphological types, ages and metallicities in the Local Group, and in the more distant Sculptor Group which contains a number of spiral galaxies easily accesible from the southern hemisphere.  One of the Araucaria target galaxies is the spiral galaxy NGC 247 (Fig. 1), member of the Sculptor Group (Jerjen et al. 1998; Karachentsev 2005). This spiral classified as SAB(s)d in the NASA/IPAC Extragalactic Database  and located in the constellation of Cetus near the Galactic South Pole and is far from the Galactic Plane, as indicated by its galactic coordinates  $\\ell = 113.95 \\deg$, $b =-83.56 \\deg$. In the following, we give a brief overview about previous studies of NGC 247 which have measured the most important properties of this galaxy.\\\\ The first authors conducting  a  photoelectric study of bright stars in NGC 247 in the U, B, V, R and I bands were Alcaino \\& Liller (1984) who provided a fundamental reference for the subsequent photometric calibration of resolved stellar populations of this galaxy.\\\\ >From Schmidt plates, Carignan (1985) studied the surface brightness of NGC 247 and computed a distance modulus of 27.01 mag for the galaxy. It was also possible to establish that NGC 247 has an inclination of 75.4 $\\deg$, indicating the difficulty involved in the study of its stellar populations due to the problems of crowding and blending accentuated by this high inclination.  Through a study of the properties of the neutral hydrogen and the mass distribution using VLA data, Carignan \\& Puche (1990) established that the HI disk of NGC 247 is relatively limited in extension. Based on their distance of 2.52 Mpc (Carignan 1985), these authors determined the total mass in HI as (8.0$ \\pm$ 0.3) $\\times 10^{8} M_{\\odot}$, which corresponds to a mass-luminosity ratio of $M_{H_I}/L_B=$0.34.  Str\\\"assle et al. (1999) made a study in the X-ray energy range between 0.25 and 2 Kev using  ROSAT data. They found that the mass of the hot gas of NGC 247, which emits mainly in the band of 0.25 Kev, is 0.7 $\\times 10^{8} M_{\\odot}$. Taking the model for dwarf galaxies of Burlak (1996), which integrates three components of the galaxy: stellar disk, gaseous disk and dark halo, these authors found that the total mass of NGC 247 is 7.5 $\\times 10^{10} M_{\\odot}$, a value consistent with their measurements of the rotation curve for this galaxy. Based on these data, Str\\\"assle et al. (1999) found a mass-luminosity ratio of $M/L_x$=28.4, which is in agreement with determinations of the same ratio for other Sculptor galaxies.  Using the MOSAIC II camera at the 4.0 m Blanco telescope at Cerro Tololo, Olsen et al. (2004) conducted a modern study of globular clusters in the Sculptor galaxies. These authors were able to detect a very small population of three globular clusters in NGC 247.  Karachentsev et al. (2006) were the first to measure the distance to NGC 247 from the Tully-Fisher method and obtained a value of 27.87 $\\pm$ 0.21 mag, in significant disagreement with the earlier, shorter distance value obtained by Carignan. In their infrared study of the central regions of nearby galaxies, Davidge \\&  Courteau (2002) had adopted a distance modulus of  27.3 mag for NGC 247. The most recent study reporting a measurement of the distance of NGC 247 from a stellar population was performed by Davidge (2006). Based on the I-band magnitude of the TRGB, he obtained a distance modulus of 27.9 $\\pm$ 0.1 mag , which implies a distance of 3.80 $\\pm$ 0.18 Mpc, making this galaxy one of the most distant targets of the Araucaria Project. Taking into account the inclination of the NGC 247 disk, Davidge stressed that variable internal extinction is expected, since the galaxy inclination favors the possibility that the line of sight passes through dust clouds more than once. Freedman et al. (1988), and Catanzarite et al. (1994) were the first to mention the discovery of a number of Cepheid variables in NGC 247. They reported that twelve Cepheids had been found from  optical surveys in B, V, R and I bands conducted at CTIO and LCO. Unfortunately, a catalog and/or photometric data for these Cepheids were never published.  This paper is the first study which reports  the discovery of a sizeable Cepheid population in NGC 247 in a verificable catalog. The paper is structured as follows: details of the observations, data reductions, and calibrations are described in \\S 2. The catalog of photometric properties of the 23 Cepheid variables in NGC 247 we have discovered is presented in \\S 3. The determination of the distance of NGC 247 from the observed period-luminosity relations in the V and I bands, and in the reddening-independent Wesenheit index is the subject of \\S 4. We discuss our results in \\S 5 and summarize our conclusions in \\S6. ",
        "conclusions": "The main conclusions of this paper can be summarized as follows:  1. We have conducted an extensive wide-field imaging survey in the Sculptor Group spiral galaxy NGC 247 in the optical V and I bands which has led to the discovery of 23 Cepheid variables with periods between 17 and 131 days. These are the first Cepheid variables with published data reported in this galaxy.  2. From the periods and mean magnitudes of the Cepheids derived from our photometry, we have constructed PL relations in the V, I and ${\\rm W}_{\\rm I}$ bands. The data define tight PL relations which are well fitted with the slopes of the corresponding LMC Cepheid PL relations adopted from the OGLE-II project. We have derived reddened distances for the 3 bands, which lead to a determination of the total mean reddening of the NGC 247 Cepheids of E(B-V) = 0.13 mag. With an adopted foreground reddening towards NGC 247 of 0.02 from Schlegel et al. (1998), we determine a reddening value of 0.11 produced {\\it inside} NGC 247 and affecting its Cepheid population. We adopt a final true distance modulus of 27.80 mag for NGC 247 from the reddening-free Wesenheit magnitude PL relation, in very good agreement with the reddening-corrected values coming from the V and I bands. The total uncertainty on this distance value, including the contribution from systematic errors, is estimated to be $\\pm$ 6\\%.  3. The Cepheid distance to NGC 247 derived in this paper agrees very well, within the respective 1 $\\sigma$ uncertainties, with the recent determinations of Davidge (2006) from the I-band TRGB method, and of Karachentsev et al. (2006) from the Tully-Fisher method.  4. The distance of NGC 247 of 3.63 Mpc is about twice as large as the respective distances of two other Sculptor Group galaxies, NGC 300 (Gieren et al. 2005a) and NGC 55 (Gieren et al. 2008a). This supports the evidence that the Sculptor Group has a filament-like structure (Jerjen et al. 1998) with a large depth extension in the line of sight."
    },
    "0808/0808.3812.txt": {
        "abstract": "%%% Abstract to run on from here. The shape of the main arc formed by the Canis Major clouds has been suggested to result from a supernova explosion possibly triggering the recent star formation activity. The presence of dozens of OB stars and reflection nebulae forms the CMa OB1/R1 associations. More than a hundred emission line stars are found in this region, including the famous Z CMa, a binary system containing a Herbig Be star and a FUor companion. Several embedded infrared clusters with different ages are associated with the CMa clouds. The main characteristics of the region in terms of cloud structure, stellar content, age of associated young clusters, distance, and X-ray emission are presented in this chapter. Some of the arguments in favor and against the hypothesis of induced star formation are discussed in the last section. ",
        "introduction": "%%% Top level section head (remove \"%\" symbol) %\\subsection{}   %%% Second level section head (remove \"%\" symbol) %\\subsubsection{}   %%% Lowest level section head (remove \"%\" symbol) %\\section*{}\t%%% Unnumbered top level section head (remove \"%\" symbol) %\\subsection*{}   %%% Unnumbered second level section head (remove \"%\" symbol) ",
        "conclusions": ""
    },
    "0808/0808.0658.txt": {
        "abstract": "In this Chapter we deal with the attempts to measure  the  general relativistic gravitomagnetic Lense-Thirring effect with the Satellite Laser Ranging (SLR) technique applied to the existing LAGEOS and LAGEOS II terrestrial satellites and to the recently approved LARES. According to general relativity, the rotation of the central body which acts as source of the gravitational field  changes the spatial orientation of the orbit of test particles determined by the longitude of the ascending node $\\Omega$. Such a shift, which is cumulative in time, is of the order of 2 m yr$^{-1}$ in the case of the LAGEOS satellites. Extracting this signature from the data is a demanding task because of many classical orbital perturbations having the same pattern as the gravitomagnetic one, like those induced by the centrifugal oblateness of the Earth which represents a major source of systematic error. The first issue addressed here is: are the so far published evaluations of the systematic bias induced by  the uncertainty in the even zonal harmonic coefficients $J_{\\ell}$ of the multipolar expansion of the Earth's geopotential  reliable and realistic? The answer is negative. Indeed, if the difference $\\Delta J_{\\ell}$ among the even zonals estimated in different global solutions (EIGEN-GRACE02S, EIGEN-CG03C, GGM02S, GGM03S, ITG-Grace02s, ITG-Grace03s, JEM01-RL03B, EGM2008) is assumed for the uncertainties $\\delta J_{\\ell}$  instead of using their more or less calibrated covariance sigmas $\\sigma_{J_\\ell}$, it turns out that the systematic error $\\delta\\mu$ in the Lense-Thirring test with the nodes of LAGEOS and LAGEOS II is about 3 to 4 times larger than in the evaluations so far published based on the use of the sigmas of one model at a time separately, amounting up to $37\\%$ for the pair EIGEN-GRACE02S/ITG-Grace03s. The comparison among the other recent GRACE-based models yields bias as large as about $25-30\\%$.  The major discrepancies still occur for $J_4, J_6$ and $J_8$, which are just the zonals the combined LAGEOS/LAGOES II nodes are most sensitive to. The second issue is the possibility of reaching a realistic total accuracy of $1\\%$ with LAGEOS, LAGEOS II and LARES, which will be launched in the near future.   While LAGEOS and LAGEOS II fly at altitudes of about 6000 km, LARES will be likely placed at an altitude of 1200 km. Thus, it will be sensitive to much more even zonals than LAGEOS and LAGEOS II. Their corrupting impact has been evaluated with the standard Kaula's approach up to degree $\\ell=60$ by using $\\Delta J_{\\ell}$; it turns out that it may be as large as some tens percent. The different orbit of LARES may also have some consequences on the non-gravitational orbital perturbations affecting it which might further degrade the obtainable accuracy. ",
        "introduction": "In the weak-field and slow motion approximation, valid when the magnitude of the gravitational potentials $U$ and velocities $v$ characteristic of the problem under examination are smaller with respect to the speed of light $c$, i.e. for $U/c^2, v/c\\lll 1$, the Einstein field equations of general relativity get linearized resembling to the Maxwellian equations of electromagntism. As a consequence, a gravitomagnetic field, induced by the off-diagonal components $g_{0i}, i=1,2,3$ of the space-time metric tensor related to the mass-energy currents of the source of the gravitational field, arises \\citep{MashNOVA}; it has no classical counterparts in Newtonian mechanics. The gravitomagnetic field affects orbiting test particles, precessing gyroscopes, moving clocks and atoms and propagating electromagnetic waves \\citep{Rug,Scia04}. Perhaps, the most famous gravitomagnetic effects are the precession  of the axis of a gyroscope \\citep{Pugh,Schi} and the Lense-Thirring\\footnote{According to an interesting historical analysis recently performed in \\citep{Pfi07}, it would be more correct to speak about an Einstein-Thirring-Lense effect.} precessions \\citep{LT} of the orbit of a test particle, both occurring in the field of a central slowly rotating mass like, e.g., our planet.   Direct, undisputable measurements of such  fundamental predictions of general relativity are not yet available.   The measurement of the gyroscope precession in the Earth's gravitational field has been the goal of the dedicated space-based\\footnote{See on the WEB http://einstein.stanford.edu/} GP-B mission \\citep{Eve, GPB} launched in 2004 and carrying onboard four superconducting gyroscopes; its data analysis is still ongoing. The target accuracy was originally $1\\%$, but it is still unclear if the GP-B team will succeed in reaching such a goal because of some unmodelled effects affecting the gyroscopes: 1) a time variation in the polhode motion of the gyroscopes and 2) very large classical misalignment torques on the gyroscopes.  In this paper we  will focus on the measurement of the Lense-Thirring effect in the gravitational field of the Earth. It consists of a secular rate of the longitude of the ascending node $\\Omega$ \\begin{equation}\\dot\\Omega_{\\rm LT} = \\rp{2 G J}{c^2 a^3 (1-e^2)^{3/2}},\\lb{let}\\end{equation} and of the argument of pericentre $\\omega$ \\begin{equation}\\dot\\omega_{\\rm LT} = -\\rp{6 G J\\cos i}{c^2 a^3 (1-e^2)^{3/2}},\\lb{LT_o}\\end{equation} of the orbit of  a test particle. In \\rfr{let} and \\rfr{LT_o} $G$ is the Newtonian constant of gravitation, $J$ is the proper angular momentum of the central body, $a$ and $e$ are the semimajor axis and the eccentricity, respectively, of the test particle's orbit and $i$ is its inclination to the central body's equator.  The semimajor axis $a$ determines the size of the ellipse, while its shape is controlled by the eccentricity $0\\leq e<1$; an orbit with $e=0$ is a circle. The angles $\\Omega$ and $\\omega$ fix the orientation of the orbit in the inertial space and in the orbital plane, respectively. $\\Omega$, $\\omega$ and $i$ can be viewed as the three Euler angles which determine the orientation of a rigid body with respect to an inertial frame. In Figure \\ref{plo} we illustrate the geometry of a Keplerian orbit. % % % % % \\begin{figure} \\includegraphics[width=\\columnwidth]{FIGURE_7.eps} \\caption{Keplerian orbit. The longitude of the ascending node $\\Omega$ is counted from a reference X direction in the reference $\\{$XY$\\}$ plane, chosen coincident with the equatorial plane of the rotating body of mass $M$ and proper angular momentum $J$, to the line of the nodes, i.e. the intersection of the orbital plane with the reference plane. The argument of pericentre $\\omega$ is an angle in the orbital plane counted from the line of the nodes; the location of the pericentre is marked with $\\Pi$. The time-dependent position of the moving test particle of mass $m$ is given by the true anomaly $f$, counted anticlockwise from the pericentre; $\\phi$ is an azimuthal angle in the $\\{$XY$\\}$ plane. $L$ is the orbital angular momentum, perpendicular to the orbital plane. The inclination between the orbital plane and the equatorial plane of $M$ is $i$. Courtesy by H. Lichtenegger, OEAW, Graz.} \\label{plo} \\end{figure} % % % % %  In this Chapter we will critically discuss the following two topics % \\begin{itemize} \\item The realistic evaluation in Section \\ref{grav} of the total accuracy in the test performed in recent years with the existing Earth's artificial satellites LAGEOS and LAGEOS II \\citep{Ciu04,Ciu06}. LAGEOS was put into orbit in 1976, followed by its twin LAGEOS II in 1992; they are passive, spherical spacecraft entirely covered by retroreflectors which allow for their accurate tracking through laser pulses sent from Earth-based ground stations according to the Satellite Laser Ranging (SLR) technique. They orbit at altitudes of about 6000 km ($a_{\\rm LAGEOS} =12270$ km, $a_{\\rm LAGEOS\\ II}=12163$ km) in nearly circular paths ($e_{\\rm LAGEOS}=0.0045$, $e_{\\rm LAGEOS\\ II}=0.014$) inclined by 110 deg and 52.65 deg, respectively, to the Earth's equator. The Lense-Thirring effect for their nodes amounts to about 30 milliarcseconds per year (mas yr$^{-1}$) which correspond to about 2 m yr$^{-1}$ at the LAGEOS altitudes.  The idea of measuring the Lense-Thirring node rate with the just launched LAGEOS satellite, along with the other SLR targets orbiting at that time, was put forth by \\citet{Cug78}. Tests have started to be effectively performed later by using the LAGEOS and LAGEOS II satellites \\citep{tanti}, according to a strategy  \\citep{Ciu96} involving the use of a suitable linear combination of the nodes $\\Omega$ of both satellites and the perigee $\\omega$ of LAGEOS II. This was done to reduce the impact of the most relevant source of systematic bias, i.e. the mismodelling in the even ($\\ell=2,4,6\\ldots$) zonal ($m=0$) harmonic coefficients $J_{\\ell}$ of the multipolar expansion of the Newtonian part of the terrestrial gravitational potential due to the diurnal rotation (they induce secular precessions on the node and perigee of a terrestrial satellite much larger than the gravitomagnetic ones. The $J_{\\ell}$ coefficients cannot be theoretically computed but must be measured by using dedicated satellites; see Section \\ref{grav}): the three-elements combination used allowed for removing the uncertainties in $J_2$ and $J_4$. In \\citep{Ciu98} a $\\approx 20\\%$ test was reported by using the\\footnote{Contrary to the subsequent models based on the dedicated satellites CHAMP (http://www-app2.gfz-potsdam.de/pb1/op/champ/index$\\_$CHAMP.html) and GRACE (http://www-app2.gfz-potsdam.de/pb1/op/grace/index$\\_$GRACE.html), EGM96 relies upon multidecadal tracking of SLR data of a constellation of geodetic satellites including LAGEOS and LAGEOS II as well; thus the possibility of a sort of $a-priori$ `imprinting' of the Lense-Thirring effect itself, not solved-for in EGM96, cannot be neglected.} EGM96 \\citep{Lem98} Earth gravity model; subsequent detailed analyses showed that such an evaluation of the total error budget was overly optimistic in view of the likely unreliable computation of the total bias due to the even zonals \\citep{Ior03,Ries03a,Ries03b}.  An analogous, huge underestimation turned out to hold also for the effect of the non-gravitational perturbations \\citep{Mil87} like the direct solar radiation pressure, the Earth's albedo, various subtle thermal effects depending on  the physical properties of the satellites' surfaces and their rotational state \\citep{Inv94,Ves99,Luc01,Luc02,Luc03,Luc04,Lucetal04,Ries03a}, which the perigees  of LAGEOS-like satellites are particularly sensitive to. As a result, the realistic total error budget in the test reported in \\citep{Ciu98} might be as large as $60-90\\%$ or even more (by considering EGM96 only).  The observable used in \\citep{Ciu04} with the GRACE-only EIGEN-GRACE02S model \\citep{eigengrace02s} was the following linear combination\\footnote{See also \\citep{Pav02,Ries03a,Ries03b}.} of the nodes of LAGEOS and LAGEOS II, explicitly computed in \\citep{IorMor} following the approach put forth in \\citep{Ciu96} \\begin{equation} f=\\dot\\Omega^{\\rm LAGEOS}+ c_1\\dot\\Omega^{\\rm LAGEOS\\ II }, \\lb{combi}\\end{equation} where \\begin{equation} c_1\\equiv-\\rp{\\dot\\Omega^{\\rm LAGEOS}_{.2}}{\\dot\\Omega^{\\rm LAGEOS\\ II }_{.2}}=-\\rp{\\cos i_{\\rm LAGEOS}}{\\cos i_{\\rm LAGEOS\\ II}}\\left(\\rp{1-e^2_{\\rm LAGEOS\\ II}}{1-e^2_{\\rm LAGEOS}}\\right)^2\\left(\\rp{a_{\\rm LAGEOS\\ II}}{a_{\\rm LAGEOS}}\\right)^{7/2}.\\lb{coff}\\end{equation} The coefficients $\\dot\\Omega_{.\\ell}$ of the aliasing classical node precessions \\citep{Kau} $\\dot\\Omega_{\\rm class}=\\sum_{\\ell}\\dot\\Omega_{.\\ell}J_{\\ell}$ induced by the even zonals  have been analytically worked out in, e.g. \\citep{Ior03}; they yield $c_1=0.544$. The Lense-Thirring signature of \\rfr{combi} amounts to 47.8 mas yr$^{-1}$. The combination  \\rfr{combi} allows, by construction, to remove the aliasing effects due to the static and time-varying parts of the first even zonal $J_2$. The nominal (i.e. computed with the estimated values of $J_{\\ell}$, $\\ell=4,6...$) bias due to the remaining higher degree even zonals would amount to  about $10^5$ mas yr$^{-1}$; the need of a careful and reliable modeling of such an important source of systematic bias is, thus, quite apparent. Conversely, the nodes of the LAGEOS-type spacecraft are  affected by the non-gravitational accelerations at a $\\approx 1\\%$ level of the Lense-Thirring effect \\citep{Luc01,Luc02,Luc03,Luc04,Lucetal04}. For a comprehensive, up-to-date overview of the numerous and subtle issues concerning the measurement of the Lense-Thirring effect see \\citep{IorNOVA}.  % % \\item  The possibility that the LARES mission, recently approved by the Italian Space Agency (ASI), will be able to measure the Lense-Thirring node precession with an accuracy of the order of $1\\%$ (Section \\ref{larez}).  In \\citep{vpe76a,vpe76b}  it was proposed to measure the Lense-Thirring precession of the nodes $\\Omega$ of a pair of counter-orbiting spacecraft to be launched in terrestrial polar orbits and endowed with drag-free apparatus. A somewhat equivalent, cheaper version of such an idea was put forth in 1986 by \\citep{Ciu86} who proposed to launch a passive, geodetic satellite in an orbit identical to that of LAGEOS  apart from the orbital planes which should have been displaced by 180 deg apart. The measurable quantity was, in the case of the proposal by \\citet{Ciu86}, the sum of the nodes of LAGEOS and of the new spacecraft, later named LAGEOS III, LARES, WEBER-SAT, in order to cancel to a high level of accuracy the corrupting effect of the multipoles of the Newtonian part of the terrestrial gravitational potential which represent the major source of systematic error (see Section \\ref{grav}). Although extensively studied by various groups \\citep{CSR,LARES}, such an idea was not implemented for many years. In \\citep{Ioretal02} it was proposed to include also the data from LAGEOS II by using a different observable.  Such an approach was proven in \\citep{IorNA} to be potentially useful in making the constraints on the orbital configuration of the new SLR satellite less stringent than it was originally required in view of the recent improvements in our knowledge of the classical part of the terrestrial gravitational potential due to the dedicated CHAMP and, especially, GRACE  missions. % For recent updates of the LARES/WEBER-SAT %mission, including recently added additional goals in fundamental physics and related criticisms, see %\\citep{LucPao01,Ior02,Ciu04b,Ciu06b,IorJCAP,Ior07c,Ior07d}.  Since reaching high altitudes and minimizing the unavoidable orbital injection errors is expensive, it was explored the possibility of discarding LAGEOS and LAGEOS II using a low-altitude, nearly polar orbit for LARES \\citep{LucPao01,Ciu06b}, but in \\citep{Ior02,Ior07c} it was proven that such alternative approaches are not feasible. It was also suggested that LARES would be able to probe alternative theories of gravity \\citep{Ciu04b}, but also in this case it turned out to be impossible \\citep{IorJCAP,Ior07d}.  The stalemate came to an end when ASI recently made the following official announcement  (http://www.asi.it/SiteEN/MotorSearchFullText.aspx?keyw=LARES): ``On February 8, the ASI board approved funding for the LARES mission, that will be launched with VEGA\u0092s maiden flight before the end of 2008. LARES is a passive satellite with laser mirrors, and will be used to measure the Lense-Thirring effect.''  The italian version of the announcement yields some more information specifying that LARES, designed in collaboration with National Institute of Nuclear Physics (INFN), is currently under construction by Carlo Gavazzi Space SpA; its Principal Investigator (PI) is I. Ciufolini and its scientific goal is to measure at a $1\\%$ level the  Lense-Thirring effect in the gravitational field of the Earth. Concerning the orbital configuration of LARES, the ASI website says about VEGA that (http://www.asi.it/SiteEN/ContentSite.aspx?Area=Accesso+allo+spazio): ``[...] VEGA can place a 15.000 kg satellite on a low polar orbit, 700 km from the Earth. By lowering the orbit inclination it can launch heavier payloads, whereas diminishing the payload mass it can achieve higher orbits. [...]''    In the latest communication to INFN, Rome, 30 January 2008, \\cite{INFN} writes that LARES will be launched with a semimajor axis of approximately 7600 km and an inclination between 60 and 80 deg. More precise information can be retrieved in Section 5.1, pag 9 of the document Educational Payload on the Vega Maiden Flight Call For CubeSat Proposals, European Space Agency, Issue 1 11 February 2008, downloadable at http://esamultimedia.esa.int/docs/LEX-EC/CubeSat$\\%$20CFP$\\%$20issue$\\%$201.pdf. It is  written there that LARES will be launched into a circular orbit with altitude $h=1200$ km, corresponding to a semimajor axis $a_{\\rm LARES}=7578$ km, and inclination $i=71$ deg to the Earth's equator.   \\end{itemize} ",
        "conclusions": "In this Chapter we have shown how the so far published evaluations of the total systematic error in the Lense-Thirring measurement with the combined nodes of the SLR LAGEOS and LAGEOS II satellites due to the classical node precessions induced by the even zonal harmonics of the geopotential  are  optimistic. Indeed, they are all based on the use of the covariance matrix's sigmas, more or less reliably calibrated, of various Earth gravity model solutions used one at a time separately in such a way that the model X yields an error of $x\\%$, the model Y yields an error $y\\%$, etc. Instead,  comparing the estimated values of the even zonals for different pairs of models   allows for a much more realistic evaluation of the real uncertainties in our knowledge of the static part of the geopotential. As a consequence, the bias in the Lense-Thirring effect measurement is about $3-4$ times larger than that so far claimed, amounting to various tens percent ($37\\%$ for the pair EIGEN-GRACE02S and ITG-GRACE03s, about $25-30\\%$ for the other most recent GRACE-based solutions).  Applying the same strategy to the ongoing LARES mission shows that the goal of reaching a $1\\%$ measurement of the Lense-Thirring effect with LAGEOS, LAGEOS II and LARES is optimistic. Indeed, since LARES will orbit at a lower altitude with respect to the LAGEOS satellites, a larger number of even zonal harmonics are to be taken into account. Assessing realistically their impact is not easy. Straightforward calculations up to degree $\\ell = 60$ with the standard Kaula's approach yield errors as large as some tens percent. Such an important point certainly deserves great attention. Another issue which may potentially degrade the expected accuracy is the impact of some non-gravitational perturbations which would have a larger effect on LARES than expected because of its lower orbit.  %-----------------------------------------"
    },
    "0808/0808.2671_arXiv.txt": {
        "abstract": "As a result of a survey of Itokawid meteors (i.e., meteors originated from Near Earth Asteroid (25143) Itokawa = 1998 $\\mathrm{SF_{36}}$), from among the multi-station optical meteor orbit data of $\\sim 15000$ orbits, and applying the $D$-criteria, we could find five Itokawid meteor candidates. We also analyzed corresponding mineral materials of the Itokawid candidates through their trajectory and atmospheric data. We conclude, on the basis of our investigation, that the fireball, MORP172, is the strongest Itokawid candidate. ",
        "introduction": "The problem as to whether meteor showers associated with Near Earth Asteroids (NEAs) exist or not has been discussed by many authors (e.g., Olsson-Steel 1988; Hasegawa et al. 1992; Babadzhanov 1994). The most prominent case is the association between NEA (3200) Phaethon and the Geminid stream complex. However, it should be an exceptional case since Phaethon is considered to be a dormant or extinct cometary nucleus; hence the Geminid stream complex meteoroids were presumably released from Phaethon in its active cometary phase (Gustafson 1989). Meanwhile, meteoroids of asteroidal origin certainly exist and they are usually observed and recognized as meteorites, meteoritic fireballs, or high bulk-density meteors. They are probably produced as impact ejecta by asteroidal collisions. Although there is no asteroid--meteor association as strong as the Phaethon--Geminid complex yet, most likely the asteroidal meteor (or meteorite) streams indeed exist (e.g., Halliday 1987; Terentjeva {\\&} Barabanov 2002; Spurn\\'{y} et al. 2003).  The HAYABUSA (MUSES-C) sample return mission target, NEA (25143) Itokawa = 1998 $\\mathrm{SF_{36}}$, is one of the Apollo-type NEAs, whose orbital parameters at present epoch are perihelion distance $\\sim 0.95$ AU and small eccentricity $\\sim 0.28$ along with an inclination near the ecliptic plane at $\\sim 1^\\circ.7$. The taxonomic class of Itokawa is less the cometary nature than a typical S(IV)-type with rather high albedo (e.g., Binzel et al. 2001; Sekiguchi et al. 2003; Ishiguro et al. 2003). A meteoritic analog would correspond to a LL chondrite (Binzel et al. 2001). Since Itokawa is a small asteroid of $\\sim 400$ m in mean axis, the escape speed from Itokawa will be only $\\sim 25$ cm s$^{-1}$ (Ostro et al. 2004). Consequently, impact ejecta should easily escape from Itokawa toward interplanetary space and some may become meteoroids. The Earth approaches Itokawa's orbit to within 0.05 AU from late March to early July, over a span of three months, in which we would be able to observe some meteors originated from Itokawa (hereafter, we designate such a meteor as ``Itokawid\") every year. If we can obtain the orbital and physical data of the Itokawid meteors, they must give us very important information when compared with the forthcoming analytical results of the HAYABUSA onboard scientific instruments and the sample of Itokawa.  On the basis of the above assumption, we surveyed whether there are Itokawid meteors or not from among the optical meteor orbit database, considering orbital similarity with Itokawa. As a result, we could locate Itokawid meteor candidates. We also analyzed the corresponding mineral materials of the Itokawid candidates through their trajectory and atmospheric data. ",
        "conclusions": "As a result of a survey of the Itokawid meteors from among our multi-station optical meteor orbit data of $\\sim 15000$ orbits, and applying the orbital similarity criteria, $D$-criteria, we could find five Itokawid meteor candidates. Two out of them, the meteor H7220 and the fireball MORP172, are probably associated with Itokawa, judging from their $D_{\\rm SH}$ values. It is also notable that the secular near-invariants of MORP172 in the $U$-$\\cos \\theta $ coordinates match well with those of Itokawa: this suggests both objects have a strong dynamical relationship to each other.  The $K_B$, $P_E$ and $A_L$ evaluations indicate all the Itokawid candidates evidently belong to asteroidal meteors. MORP172 is the only one classified at ordinary/carbonaceous chondrite, the others being in carbonaceous chondrites. Hence, the classified mineral matter of MORP172 shows more similarity to Itokawa's surface composition of a LL chondrite analogue than the other candidates.  We conclude, on the basis of the investigations above, that the fireball, MORP172, is the strongest Itokawid candidate. However, if you ask whether there are Itokawid meteors or not, we will answer, ``Possible\", since only a small sample like this is available so far. Therefore, more optical detections and analyses for the Itokawid meteors are desirable in the future work."
    },
    "0808/0808.0990_arXiv.txt": {
        "abstract": "Semiclassical states in isotropic loop quantum cosmology are employed to show that the improved dynamics has the correct classical limit. The effective Hamiltonian for the quantum cosmological model with a massless scalar field is thus obtained, which incorporates also the next to leading order quantum corrections. The possibility that the higher order correction terms may lead to significant departure from the leading order effective scenario is revealed. If the semiclassicality of the model is maintained in the large scale limit, there are great possibilities for $k=0$ Friedmann expanding universe to undergo a collapse in the future due to the quantum gravity effect. Thus the quantum bounce and collapse may contribute a cyclic universe in the new scenario. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.0959_arXiv.txt": {
        "abstract": "{% The role of shear in alleviating catastrophic quenching by shedding small-scale magnetic helicity through fluxes along contours of constant shear is discussed. The level of quenching of the dynamo effect depends on the quenched value of the turbulent magnetic diffusivity. Earlier estimates that might have suffered from the force-free degeneracy of Beltrami fields are now confirmed for shear flows where this degeneracy is lifted. For a dynamo that is saturated near equipartition field strength those estimates result in a 5-fold decrease of the magnetic diffusivity as the magnetic Reynolds number based on the wavenumber of the energy-carrying eddies is increased from 2 to 600. Finally, the role of shear in driving turbulence and large-scale fields by the magneto-rotational instability is emphasized. New simulations are presented and the $3\\pi/4$ phase shift between poloidal and toroidal fields is confirmed. It is suggested that this phase shift might be a useful diagnostic tool in identifying mean-field dynamo action in simulations and to distinguish this from other scenarios invoking magnetic buoyancy as a means to explain migration away from the midplane. ",
        "introduction": "Shear clearly plays an important role in amplifying toroidal fields from poloidal, but that is not all. Shear also plays a role in ``unquenching'' any dynamo effect that may play a role in producing poloidal field from toroidal, thus closing the dynamo loop. A prime example of such a dynamo effect is the $\\alpha$ effect, but other possible known effects may include the shear-current effect and the incoherent alpha-shear effect. The ``unquenching'' of dynamo effects, as well as the dynamo effects themselves, require more detailed considerations of the results available so far. This is the principal goal of this paper.  The term ``unquenching'' in connection with the $\\alpha$ effect may appear somewhat unusual, but this choice of words must be seen in contrast to the possibility of catastrophic $\\alpha$ quenching. Here, ``catastrophic'' indicates that the quenching by the mean magnetic field, $\\meanBB$, becomes more extreme as the magnetic Reynolds number, $\\Rm$, increases. Traditionally, such quenching is represented by the simplistic formula (Vainshtein \\& Cattaneo 1992) \\EQ \\alpha(\\meanBB)={\\alpha_0\\over1+\\Rm\\meanBB^2/B_{\\rm eq}^2}, \\label{QuenchSimple} \\EN where $B_{\\rm eq}=\\bra{\\mu_0\\rho\\uu^2}^{1/2}$ is the equipartition field strength with respect to the kinetic energy density, $\\uu$ is the small-scale turbulent velocity, $\\rho$ is the density, $\\mu_0$ is the vacuum permeability, $\\Rm=u_{\\rm rms}/\\eta\\kf$ is the magnetic Reynolds number, $\\eta$ is the magnetic diffusivity, $\\kf$ is the wavenumber of the energy-carrying scale, and angular brackets denote volume averaging. However, \\Eq{QuenchSimple} is really only valid under special circumstances that are quite uninteresting for dynamo action: infinite wavelength of the magnetic field, complete stationarity, and no possibility of magnetic helicity fluxes. The latter possibility is now believed to be the most important one for astrophysical dynamos, as was first suggested by Blackman \\& Field (2000). Any one of these three caveats alleviates the severity of catastrophic quenching, because they all lead to ``extra terms'' that enter in the numerator of \\Eq{QuenchSimple} with an $\\Rm$ factor in front, i.e.\\ \\EQ \\alpha(\\meanBB)={\\alpha_0+\\Rm\\times\\mbox{``extra effects''} \\over1+\\Rm\\meanBB^2/B_{\\rm eq}^2}, \\label{QuenchExtra} \\EN where we have assumed that the kinematic $\\alpha$ value, $\\alpha_0$, remains independent of time. In its essence, this equation with extra effects included goes back to the early work of Kleeorin \\& Ruzmaikin (1982), and later Kleeorin et al.\\ (1995, 2000), and has been discussed in connection with alleviating catastrophic $\\alpha$ quenching by Blackman \\& Brandenburg (2002) and Brandenburg \\& Subramanian (2005a). A more complete quenching formula with extra effects included takes the form \\EQ \\alpha={\\alpha_0+\\Rm\\left( \\eta_{\\rm t}{\\mu_0\\meanJJ\\cdot\\meanBB\\over B_{\\rm eq}^2} -{\\nab\\cdot\\meanFF_{\\rm C}\\over2 k_{\\rm f}^2B_{\\rm eq}^2} -{\\partial\\alpha/\\partial t\\over2\\eta_{\\rm t} k_{\\rm f}^2}\\right) \\over1+\\Rm\\meanBB^2/B_{\\rm eq}^2}. \\label{QuenchExtra2} \\EN Let us now discuss separately all three terms in the parenthesis of the numerator of \\Eq{QuenchExtra2}.  (i) The effect of the $\\meanJJ\\cdot\\meanBB$ term is clearly seen when considering the saturation of homogeneous dynamos in a periodic domain. In that case this formula gives \\EQ \\alpha(\\meanBB)\\rightarrow\\eta_{\\rm t}\\mu_0\\meanJJ\\cdot\\meanBB/\\meanBB^2 \\quad\\mbox{(for $\\Rm\\to\\infty$)}, \\EN so there is nothing catastrophic about this formula, unless $\\eta_{\\rm t}$ itself is catastrophically quenched. Of course, if mean fields are defined as full volume averages, $\\meanBB$ becomes completely uniform, so $\\mu_0\\meanJJ=\\nab\\times\\meanBB=0$. This is the case in numerical experiments by Cattaneo \\& Hughes (1996). The other case with a finite $\\meanJJ\\cdot\\meanBB$ term was seen in the simulations of Brandenburg (2001), where not only a large-scale magnetic field was found to saturate at super-equipartition values, but also the $\\alpha$ and $\\eta_{\\rm t}$ effects were found to be only mildly quenched.  (ii) The $\\partial/\\partial t$ term in \\Eq{QuenchExtra2} is important to explain the absence of an otherwise premature onset of quenching at resistively low field strengths where $\\meanBB^2/B_{\\rm eq}^2\\approx\\Rm^{-1}$. This is also confirmed by controlled numerical experiments where the initial field was a weak Beltrami field (Brandenburg et al.\\ 2003).  (iii) Finally, the effect of magnetic helicity fluxes was first seen in simulations with imposed fields by Brandenburg \\& Sandin (2004), and then later in dynamo simulations (Brandenburg 2005a), which brings us to the main topic of this paper. As was already seen in earlier simulations without shear, just allowing for open boundary conditions alone does not help to produce a finite magnetic helicity flux in \\Eq{QuenchExtra2} and hence does not alleviate catastrophic $\\alpha$ quenching (Brandenburg \\& Dobler 2001). However, Vishniac \\& Cho (2001) showed that in the presence of differential rotation a magnetic helicity flux can be generated and that it flows along the rotation axis. More detailed work of Subramanian \\& Brandenburg (2004, 2006) and Brandenburg \\& Subramanian (2005b) resulted in a simple formula for this particular contribution to the flux: \\EQ \\meanFF_{\\rm C}=C_{\\rm VC}(\\meanSSSS\\,\\meanBB)\\times\\meanBB, \\EN where $\\meanSSS_{ij}=\\half(\\meanU_{i,j}+\\meanU_{j,i})$ is the rate of strain matrix of the mean flow, and $C_{\\rm VC}$ is a dimensionless number of order unity. This flux is along contours of constant shear, as was demonstrated by Brandenburg et al.\\ (2005); see also \\Fig{p2vishcho}.  \\begin{figure}[t!]\\begin{center} \\includegraphics[width=\\columnwidth]{p2vishcho} \\end{center}\\caption[]{ Vectors of $\\meanFF_{\\rm C}$ together with contours of $\\meanU_y$ which also coincide with the streamlines of the mean vorticity field $\\meanWW$. Note the close agreement between $\\meanFF_{\\rm C}$ vectors and $\\meanWW$ contours. The orientation of the vectors indicates that negative current helicity leaves the system at the outer surface ($x=0$). Adapted from Brandenburg et al.\\ (2005). }\\label{p2vishcho}\\end{figure}  The sign of the magnetic helicity flux is negative in the Northern hemisphere, so vectors of positive magnetic helicity flux point away from the surface in the Northern hemisphere, then through the equator and into the Southern hemisphere. ",
        "conclusions": "\\end{figure}  Here we have used $S<0$, so we are looking for negative values of $\\eta_{21}$ for an operational shear-current effect. In all cases we find $\\eta_{21}>0$, except for large values of $\\Rm$ when $\\Pm=20$. However, the error bars are large. There is perhaps the possibility that the sign of $\\eta_{21}$ may change under other circumstances, e.g.\\ for the more complicated shear profile shown in \\Fig{p2vishcho}. Yet another possibility is that in the presence of helicity the sign may change. Some evidence to this effect has been provided by Mitra et al.\\ (2008). It turns out that in nonhelical shear flow turbulence, $\\eta_{21}$ can become negative in the saturated state with helicity. However, given that there is still an $\\alpha$ effect, the relevance of the shear-current effect is less obvious in these simulations.  Obviously, with helicity there is also an $\\alpha$ effect, so the shear-current effect would not be the sole cause of large-scale dynamo action. This leaves us with the question what causes large-scale magnetic field generation in non-helical turbulence with shear? In addition to the simulations of Brandenburg (2005a), discussed above, there are also simulations of Yousef et al.\\ (2008), where a large-scale magnetic field is generated. They find a growth rate that is proportional to the shear rate $S$. The authors argue that such a result would not be consistent with the shear-current effect, because the growth rate would then be proportional to the product of $S$ and $\\eta_{21}$, where $\\eta_{21}$ itself would be proportional to $S$.  The alternative proposal by Brandenburg et al.\\ (2008a) is that an incoherent $\\alpha$-shear or $\\alpha\\Omega$ dynamo is at work. The occurrence of an incoherent $\\alpha$-shear dynamo (cf.\\ Vishniac \\& Brandenburg 1997; Proctor 2007; Kleeorin \\& Rogachevskii 2008) was quantified by estimating the rms values of the fluctuations of all components of $\\alpha_{ij}$ and $\\eta_{ij}$. \\FFig{kinshear_Re14_alprms} shows how these rms values vary with increasing values of $\\Rm$. We recall that the average of all components of $\\alpha_{ij}$ is zero, so there is no regular $\\alpha$ effect. The onset of incoherent dynamo action depends on the value of the dynamo number \\EQ D_{\\alpha S}^{\\rm incoh}=\\alpha_{\\rm rms} S/\\eta_{\\rm T}k_1^3, \\EN where $\\alpha_{\\rm rms}$ is the rms value of the streamwise component $\\alpha_{22}$ (but all components are found to have the same rms value), $\\eta_{\\rm T}=\\eta_{\\rm T}+\\eta$ is the sum of turbulent and microscopic magnetic diffusivity, and $k_1=2\\pi/L$ is the lowest wavenumber in the $z$ direction of the box.  \\begin{figure}[t!] \\centering\\includegraphics[width=\\columnwidth]{kinshear_Re14_alprms}\\caption{ Dependences of the rms values of the temporal fluctuations $\\alpha_{\\rm rms}$ (normalized by $\\eta_{\\rm t0}k_{\\rm f}$), $\\eta_{\\rm t}^{\\rm rms}$, $\\eta_{21}^{\\rm rms}$, and $\\eta_{12}^{\\rm rms}$ (normalized by $\\eta_{\\rm t0}$), on $\\Rm$ for $\\Rey=1.4$ and $\\mbox{Sh}=-0.6$. Adapted from Brandenburg et al.\\ (2008a). }\\label{kinshear_Re14_alprms}\\end{figure}  The simulations of Brandenburg et al.\\ (2008a) give fluctuations of $\\alpha_{\\rm rms}$ that correspond to $D_{\\alpha S}^{\\rm incoh}\\approx4$. This is to be compared with the marginal value of $\\approx2.3$ obtained numerically using a single-mode approximation. There is in principle also the incoherent shear-current effect, based on the random occurrence of values of $\\eta_{21}$ with suitable sign. However, the effect is found to be subdominant compared with the dynamo number of the incoherent $\\alpha$-shear dynamo. Yet another possibility for generating large-scale magnetic fields in the presence of shear is given by additive noise in the electromotive force due to small-scale dynamo action (Blackman 1998). However, in Yousef et al.\\ (2008) no small-scale dynamo was excited, so this type of noise would not explain the field seen in their simulation.  In all cases one would expect random reversals of the toroidal magnetic field on a turbulent diffusion time scale. However, simulations usually give longer time scales, which may be explained by a tendency to conserve magnetic helicity. Indeed, using  magnetic helicity conservation, Brandenburg et al.\\ (2008a) found that in a closed domain the reversal time increases with $\\Rm$ to the 1/2 power. The mean reversal time might therefore well be much larger than a turbulent diffusion time.    In this paper we have illustrated the important dual role played by shear: helping to produce strong toroidal field and helping to unquench the $\\alpha$ effect. We have tried to make the case that successful large-scale dynamos must have an opportunity to shed small-scale magnetic helicity through the boundaries. This is what the Sun does (e.g., D\\'emoulin et al.\\ 2002), and this is also what many simulations can do, although the degree of realism to which simulations can do this varies. So far, there is no simulation that goes all the way to including at least a simplified way of modeling coronal mass ejections as the final stage in shedding small-scale magnetic helicity. This should obviously be tried in the near future using simulations in spherical shell segments, possibly with an additional outer layer where wind acceleration can occur (Brandenburg et al.\\ 2007).  In this connection it might be helpful to explain the special meaning of ``small scale'' in connection with large-scale dynamos. In solar physics, one is used to associating active regions with large scales, while small scales would often refer to the resolution limit of $100\\km$ or less. Obviously, in connection with the solar cycle, relevant time and length scales would be on the order of years and several hundred megameters, respectively -- well encompassing the duration and size of active regions. In the present work we have used averages that do not a priori separate between large and small scales. By assuming an average over one or two coordinate directions (e.g.\\ an azimuthal average for global simulations of the Sun, or horizontal averages in local box simulations), there could still be residual fluctuations on short time and small length scales, but they should go away with increasing system size in the azimuthal or horizontal directions like one over the square root of the number of eddies that are being averaged over. Exactly this fact is behind the operation of the incoherent $\\alpha$-shear dynamo discussed at the end of \\Sec{Incoherent}. Clearly a finite horizontal extent is not only a numerical restriction of local boxes. Indeed, a finite azimuthal extent is bare reality even for the Sun. This point has been made by Hoyng (1993) in order to explain limits to the phase and amplitude stability of the solar cycle, but this very mechanism could even suffice to drive a dynamo in shearing systems where there is no $\\alpha$ effect present. It should be emphasized that the incoherent $\\alpha$ effect as such is independent of shear, but it is able to produce coherent large-scale fields only in the presence of shear. Obviously, more work is needed to sharpen the analytical and numerical tools to understand such systems better. Finally, let us emphasize that incoherent $\\alpha$-effect dynamos too have a nonlinear stage that is controlled by magnetic helicity at some level. In Brandenburg et al.\\ (2008a) we only used a very simplistic one-mode truncation to make this point, but the whole dynamical quenching formalism applies otherwise just as well."
    },
    "0808/0808.1499_arXiv.txt": {
        "abstract": "{} {We analyze the  first X-ray observations with XMM-Newton of 1RXS~J070407.9+262501 and 1RXS~180340.0+401214, in order to characterize their broad-band temporal and spectral properties, also in the UV/optical domain, and to confirm them as Intermediate Polars.} {For both objects, we performed a timing analysis of the X-ray and UV/optical light curves to detect the white dwarf spin pulsations and study their energy dependence. For 1RXS~180340.0+401214 we also analyzed optical spectroscopic data to determine the orbital period. X-ray spectra were analyzed in the 0.2--10.0~keV range to characterize the emission properties of both sources.} {We find that the X-ray light curves of both systems are energy dependent and are dominated, below 3--5~keV, by strong pulsations at the white dwarf rotational periods (480~s for 1RXS~J070407.9+262501 and 1520.5~s for 1RXS~180340.0+401214). In 1RXS~180340.0+401214 we also detect an X-ray beat variability at 1697~s which, together with our new optical spectroscopy, favours an orbital period of 4.4~hr that is longer than previously estimated. Both systems show complex spectra with a hard (up to 40~keV) optically thin and a soft (85--100~eV) optically thick components heavily absorbed by material partially covering the X-ray sources.} {Our observations confirm the two systems as Intermediate Polars and also add them as new members of the growing group of 'soft' systems which show the presence of a soft X-ray blackbody component. Differences in the temperatures of the blackbodies are qualitatively explained in terms of reprocessing over different sizes of the white dwarf spot. We suggest that systems showing cooler soft X-ray blackbody components also possess white dwarfs irradiated by cyclotron radiation.} ",
        "introduction": "Intermediate polars (IPs) are a subclass of magnetic cataclysmic variables (CVs) in which the rotational period $P_{\\omega}$ of the white dwarf (WD) is not synchronized with the orbital period $P_{\\Omega}$ of the binary system ($P_{\\omega} \\ll P_{\\Omega}$). The detection of a modulation at the WD spin period in the X-ray regime represents a signature of magnetic accretion. The material lost from the secondary star generally passes through a disc and is channelled onto the magnetic polar regions of the WD where a strong shock develops \\citep{aizu73}. Hard X-ray radiation is emitted by the hot post-shock gas that cools via thermal bremsstrahlung while it settles onto the WD surface. Because of the highly asynchronous WD rotation and of the lack of detectable polarized emission in the optical and near-IR in most systems, IPs were believed to possess weakly magnetized WDs ($B \\sim 0.1-10\\ \\mathrm{MG}$, \\citet{warner}). A distinct soft ($k T \\sim 20-50\\ \\mathrm{eV}$) X-ray blackbody emission was discovered by the \\emph{ROSAT} satellite in a handful of IPs \\citep{mason92,haberletal94,haberlmotch95,burwitzetal96}, some of them also showing polarized optical radiation indicating field strengths up to 20--30~MG (e.~g. \\citet{piirolaetal93,piirolaetal08}). This spectral component resembles that observed in the X-ray spectra of Polars, the other subclass of magnetic CVs, that instead possess synchronously rotating and  strongly magnetized ($B \\sim 10 - 230\\ \\mathrm{MG}$) WDs.  Our current understanding of the origin of soft X-ray emission and energy balance in Polars has greatly improved since the early model of \\citet{lamb_masters79}. It has become evident that the soft and hard X-ray emissions in at least these systems are largely decoupled \\citep[see][]{beuermann99,beuermann04}, most of the soft X-rays arising from accretion of dense blobs (`blobby accretion') which penetrate deep in the WD atmosphere. Irradiation of the WD atmospheric polar regions by hard X-rays and cyclotron radiation emitted in a stand-off shock has been recently discussed by \\citet{konig06}, yelding to the conclusion that most of the reprocessed light appears in the far-UV rather than in the soft X-rays.  The similarity with the Polars raised the question on whether the `soft IPs' are  the `true' progenitors of low-field Polars (see also \\citet{norton04}). However, very recently a soft and rather absorbed X-ray emission is being found in a growing number of IPs \\citep{haberl02,demartino04,demartino06a,demartino06b,demartino08,evanshellier07,staude08}. Whether all IPs possess such a component and whether it balances the primary emission from the post-shock flow is a new aspect that deserves investigation.  Furthermore, in recent years the number of candidate members of the IP group has grown quite rapidly \\citep{woudtwarner04,gan05} almost doubling the number of previously known systems as well as enlarging the range of asynchronism. While this could suggest an evolution towards synchronism \\citep{norton04}, the confirmation of new candidates is essential to understand the evolutionary link between IPs and Polars.  1RXS~J070407.9+262501 (hereafter RXJ0704) was listed in the ROSAT All-Sky Survey Source Catalogue (RASS, \\citet{vog99}) as a faint X-ray source, with a count rate of $0.19\\ \\mathrm{cts\\ s}^{-1}$ in the 0.1--2.4~keV PSPC range and ROSAT hardness ratios~\\footnote{ROSAT hardness ratios 1 and 2 are defined as $\\mathrm{HR1} = \\frac{(0.40 - 2.40) - (0.07 - 0.40)}{(0.07 - 2.40)}$ and $\\mathrm{HR2} = \\frac{(1.00 - 2.40) - (0.40 - 1.00)}{(0.40 - 2.40)}$, where $(A - B)$ is the raw count rate in the $A - B$ energy range expressed in keV \\citep{motchetal96}.} $\\mathrm{HR1} = -0.40$ and $\\mathrm{HR2} = +0.55$ that suggested the presence of a soft component in the X-ray spectrum. RXJ0704 was then identified as a CV by \\citet{gan05} and proposed to be an IP from the detection of optical photometric and spectroscopic periodicities at 480.71~s and $\\sim 250\\ \\mathrm{min}$, respectively. The short period variability was ascribed to the spin period of the WD, while the other was identified as the orbital period of the binary system. The optical light curve folded at the putative WD rotational period was double peaked, suggesting the presence of two equally contributing accreting poles. However, no other X-ray observation was performed on this source in order to study its X-ray properties and hence to confirm its membership in the IP class.  1RXS~180340.0+401214 (hereafter RXJ1803) was also discovered in the RASS as a faint (PSPC count rate $0.18\\ \\mathrm{cts\\ s}^{-1}$) source with a very soft X-ray spectrum ($\\mathrm{HR1} = -0.18$ and $\\mathrm{HR2} = -0.02$). Later, it was also identified as an IP candidate by \\citet{gan05} from its optical photometric variability at a 1520.5~s period that was interpreted as the rotational period of the accreting WD. A spectroscopic period of 160.2~min was inferred from the analysis of radial velocity variations in the $\\mathrm{H}_{\\alpha}$ emission line wings and identified as the binary orbital period, thus placing this CV in the 2--3~hr period gap. The two periods would then indicate a system with a weak degree of asynchronism. Also for this source no further X-ray observations were available.  In this work we report XMM-Newton observations of RXJ0704 and RX1803 aiming at determining for the first time their X-ray variability as well as their X-ray broad-band spectral properties. For RXJ1803 we further complement the XMM-Newton data with optical spectroscopy acquired at the Observatoire de Haute Provence (OHP, France). ",
        "conclusions": "XMM-Newton observation of RXJ0704 and RXJ1803 confirm these CVs as true members of the IP class. Both systems are dominated by strong pulsations at the WD rotational periods which dominate the soft portion of their spectrum, i.~e. below 3-5~keV. With $P_{\\omega} = 480\\ \\mathrm{s}$, RXJ0704 joins the group of fast rotators and possess a relatively high degree of asynchronism ($P_{\\omega} / P_{\\Omega} = 0.04$). RXJ1803 is, instead, a slow rotator $P_{\\omega} = 1520.5\\ \\mathrm{s}$. From our new optical spectroscopy and the X-ray data of this system we infer an orbital period of 4.4~hr, that is longer than previously determined. This gives a spin-to-orbit period ratio of $\\sim 0.1$, typical for IPs.  The detection of strong spin pulses indicate that accretion occurs preferentially via a disc in both systems. However, while RXJ0704 only shows X-ray variability at the spin period, RXJ1803 also displays a weak variability at the orbital sideband ($\\omega -\\Omega$), implying that accretion also occurs via disc-overflow.  Their X-ray spectra are compatible with composite emission consisting of an optically thin component ($\\sim 11-25\\ \\mathrm{keV}$), which in RXJ1803 is better described by a multi-temperature plasma, reaching values up to 30--40~keV, as well as by blackbody emission with a relatively high temperature (85--100~eV) heavily absorbed by dense material partially covering the X-ray source. This adds these two systems to the growing group of `soft IPs', which indicates that the presence of a soft blackbody component is not solely a characteristic of the Polars. However, a major difference is in the temperatures of the blackbodies in the IPs which cover a wider range than that observed in the Polar systems. We qualitatively try to explain the differences in terms of reprocessing over different sizes of the WD spot. We suggest that IPs showing cooler soft X-ray blackbody components also possess WDs irradiated by cyclotron radiation. A further systematic observational approach addressing the detection of polarized optical/IR emission in IPs is highly desiderable."
    },
    "0808/0808.3987_arXiv.txt": {
        "abstract": "Transiting planets like HD209458b offer a unique opportunity to scrutinize their atmospheric composition and structure. Transit spectroscopy probes the transition region between the day and night sides, called limb. We present a re-analysis of existing HST/STIS transmission spectra of HD209458bs atmosphere. From these observations we: Identify H2 Rayleigh scattering, derive the absolute Sodium abundance and quantify its depletion in the upper atmosphere, extract a stratospheric T-P profile with a temperature inversion and explain broad band absorptions with the presence of TiO and VO molecules in the atmosphere of this planet. ",
        "introduction": "Because of the wavelength-dependent opacities of absorbing species, measurement of relative changes in eclipse depth as a function of wavelength during primary transit has the potential to reveal the presence (or absence) of specific chemical species (Seager et al.\\ 2000a, Hubbard et al.\\ 2001, Brown et al.\\ 2001).  In the case of HD~209458b's atmosphere, the transmission spectroscopy method led  to the detection of sodium (Charbonneau et~al. 2002, Sing et~al. 2008b). In the UV, absorptions of several percents for H\\,{\\sc i} Lyman-$\\alpha$, O\\,{\\sc i} and C\\,{\\sc ii} have been measured in the hydrodynamically escaping upper atmosphere (Vidal-Madjar et al.\\ 2003, 2004, 2008, D\\'esert et al.\\ 2004, Lecavelier\\ 2007, Ehrenreich et al. 2007).  We use public archived {\\it STIS} spectra obtained during planetary transit at two spectral resolutions (low and medium). Both datasets are combined to extend the measurements over the entire optical regime to quantify possible absorbers appearing in the transmission spectrum (Sing et al. 2008a).   In this poster we present the identification of Rayleigh scattering by H$_2$ molecules (Lecavelier et~al. 2008b), the depletion of NaI as well as the extraction of a \\TP\\ with inversion (Sing et~al. 2008b). Finally, we derived the upper limits for the TiO/VO molecular abundances in the atmosphere of the planet (D\\'esert et al. 2008). ",
        "conclusions": ""
    },
    "0808/0808.0710.txt": {
        "abstract": "%\\begin{description}",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.3725_arXiv.txt": {
        "abstract": "The space-borne antimatter experiment PAMELA has recently reported a surprising rise in the positron to electron ratio at high energies. It has also recently been  found that electromagnetic radiative corrections in some cases may boost the gamma-ray yield from supersymmetric dark matter annihilations in the galactic halo by up to three or four orders of magnitude, providing distinct spectral signatures for indirect dark matter searches to look for. Here, we investigate whether the same type of corrections can also lead to sizeable enhancements in the positron yield. We find that this is indeed the case, albeit for a smaller region of parameter space than for gamma rays; selecting models with a small mass difference between the neutralino and sleptons, like in the stau coannihilation region in mSUGRA, the effect becomes more pronounced. The resulting, rather hard positron spectrum with a relatively sharp cutoff may potentially fit the rising positron ratio measured by the PAMELA satellite. To do so, however, very large ``boost factors'' have to be invoked that are not expected in current models of halo structure. If the predicted cutoff would also be confirmed by later PAMELA data or upcoming experiments, one could either assume non-thermal production in the early universe or non-standard halo formation to explain such a spectral feature as an effect of dark matter annihilation. At the end of the paper, we briefly comment on the impact of radiative corrections on other annihilation channels, in particular antiprotons and neutrinos. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.4105_arXiv.txt": {
        "abstract": "We present grids of stellar models and their associated oscillation frequencies that have been used by the \\corot\\ \\textit{Seismology Working Group} during the scientific pre\\-paration of the \\corot\\ mission. The stellar models have been calculated with the \\cesam\\ stellar internal structure and evolution code while the oscillation frequencies have been obtained from the \\cesam\\ models by means of the \\adipls\\ adiabatic oscillation programme. The grids cover a  range of masses, chemical compositions and evolutionary stages corresponding to those of the \\corot\\ primary targets. The stellar models and oscillation frequencies are available on line through the \\textit{Evolution and Seismic Tools Activity} (ESTA) web site. ",
        "introduction": "During the scientific preparation of the \\corot\\ mission, a reference frame in the H--R diagram had to be provided for the \\corot\\ potential target stars. In particular, the critical selection of the primary targets and the census and characterisation of the stellar content in the fields surrounding the potential primary targets have required to locate all considered stars (or stellar systems) in the H--R diagram and eventually to estimate their mass, evolutionary stage and expected oscillation spectra. For that purpose we have calculated the internal structure and evolution of stars in a range of masses, chemical compositions and evolution stages of interest for the modelling of \\corot\\ targets \\citep{2005sf2a.conf..283M}. We have used the {\\cesam}\\footnote{Code d'Evolution Stellaire Adaptatif et Modulaire} stellar evolution code \\citep{Morel97,pm-apss}. We have selected several models along each evolution sequence and we have calculated their oscillation frequencies by means of the \\adipls\\footnote{Aarhus Adiabatic Pulsation Package), available at \\tt\\small http://astro.phys.au.dk/jcd/adipack.n} programmes \\citep{1982MNRAS.199..735C, jcd2-apss}.  In this paper we briefly present these grids of models and oscillation frequencies. Section~\\ref{sec:models} presents the {\\cesam} numerical code and the input physics and initial parameters used to calculate the grids while Section~\\ref{sec:osc} presents the oscillation data obtained from the {\\adipls} adiabatic oscillation code. The present models and associated oscillation frequencies have been made available on the ESTA web site\\footnote{\\tt\\small http://www.astro.up.pt/corot/models/ and also on the Paris Observatory web site at \\tt\\small http://wwwusr.obspm.fr/$\\sim$lebreton/Modeles/CESAM\\_COROT.html}. ",
        "conclusions": "\\label{sec:conclusion}  { On the {\\esta} web site we have provided grids of stellar models and oscillation frequencies that have been used by the \\corot\\ \\textit{Seismology Working Group} during the scientific preparation of the \\corot\\ mission. The stellar models are provided for a range of masses between $1.0$ and $10.0$ {\\msol} and chemical compositions [Fe/H]=$0.0$ and $-1.0$ dex representative of the main  {\\corot} targets (solar-type stars, $\\delta$ Scuti stars and $\\beta$ Cephei stars). Oscillation frequencies are provided all along the main-sequence track for a solar-type oscillator ($M=1.2$ {\\msol}) and a $\\delta$ Scuti pulsator ($M=1.8$ {\\msol}). Up to now these reference grids have been used to locate {\\corot} potential targets in the H--R diagram in the process of target selection \\citep[see Fig. 1 in][]{yl1-apss} and to study some $\\delta$ Scuti candidates for {\\corot} \\citep{2005AJ....129.2461P}.  On the {\\esta} web site, other grids of stellar models and oscillation frequencies can be found. These stellar models have been calculated either with the {\\CESAM} or the {\\CLES} code and the oscillation frequencies have been calculated either with the {\\ADIPLS}, {\\POSC} or {\\LOSC} code. All this material (codes and grids)  is described in this volume by \\citet{yl1-apss, jm1-apss,mmf-apss}. We point out that the thorough comparisons of the stellar evolutionary codes and oscillations codes performed under {\\esta} have shown very good agreement between models calculated by the {\\CESAM} and {\\CLES} codes which gives us confidence in the use of the present grids \\citep[see][]{yl3-apss,jm2-apss,moya-apss}. }"
    },
    "0808/0808.2179_arXiv.txt": {
        "abstract": "Particle-in-cell (PIC) simulations of collisionless magnetic reconnection are performed to study asymmetric reconnection in which an outflow is blocked by a hard wall while leaving sufficiently large room for the outflow of the opposite direction. This condition leads to a slow, roughly constant motion of the diffusion region away from the wall, the so-called `X-line retreat'. The typical retreat speed is $\\sim$0.1 times the Alfv\\'{e}n speed. At the diffusion region, ion flow pattern shows strong asymmetry and the ion stagnation point and the X-line are not collocated. A surprise, however, is that the reconnection rate remains the same unaffected by the retreat motion. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.2981_arXiv.txt": {
        "abstract": "We perform a series of simulations of a Galactic mass dark matter halo at different resolutions, our largest uses over three billion particles and has a mass resolution of $1000 M_\\odot$. We quantify the structural properties of the inner dark matter distribution and study how they depend on numerical resolution. We can measure the density profile to a distance of 120 pc (0.05\\% of $R_{\\rm vir}$) where the logarithmic slope is -0.8 and -1.4 at (0.5\\% of $R_{\\rm vir}$). We propose a new two parameter fitting function that has a linearly varying logarithmic density gradient which fits the GHALO and VL2 density profiles extremely well. Convergence in the density profile and the halo shape scales as $N^{-1/3}$, but the shape converges at a radius three times larger at which point the halo becomes more spherical due to numerical resolution. The six dimensional phase-space profile is dominated by the presence of the substructures and does not follow a power law, except in the smooth under-resolved inner few kpc. ",
        "introduction": "Over twenty five years ago the theoretical framework for the evolution of a cold dark matter (CDM) dominated universe was established \\citep{1982ApJ...263L...1P}. The hierarchical and violent growth of structure in this model begins at a scale of $~10^{-6} M_\\odot$ \\citep{2005Natur.433..389D} until the most massive clusters of galaxies form that are many orders of magnitude more massive. The assumption that the dark matter is cold remains to be verified, yet numerical simulations that follow the hierarchical formation of CDM haloes have given several fundamental and robust predictions for the structural and substructure properties of the dark matter distribution within virialised haloes \\citep{1991ApJ...378..496D,1996ApJ...462..563N,1999MNRAS.310..527A,2001MNRAS.321..559B,2001ApJ...557..533F}. These results are widely used to compare with observational data and to assist comparisons with analytic models.   The first CDM halo simulated with enough resolution to resolve substructure used $10^6$ particles \\citep{1998ApJ...499L...5M}, resolving the density profile to about one percent of the virial radius \\citep{1998MNRAS.300..146G,1999MNRAS.310.1147M,2003MNRAS.338...14P,2004MNRAS.349.1039N, 2004MNRAS.353..624D}. Whilst such simulations find numerous substructures in the outer halo, they find few or none within the inner 20\\% of $R_{\\rm vir}$ and no obvious structure in phase-space in the central halo regions \\citep{2001PhRvD..64f3508M}. Advances in algorithms and supercomputing power have recently allowed us to increase this resolution by over two orders of magnitude with the Via Lactea II (VL2) simulation \\citep{2008arXiv0805.1244D}.   There are several reasons why we wish to carry out further studies at a higher resolution: (i) There are many old and forthcoming observational tests that constrain the structure of dark matter haloes on scales well within $0.001R_{\\rm vir}$. These include high resolution rotation curve data and the kinematics of stars at the centres of dwarf galaxies. Future proper motions of these inner stars with GAIA or SIM will provide even tighter constraints. The close binary nuclei in galaxies such as VCC128 constrains the dynamics on even smaller scales \\citep{2008arXiv0806.1951G}. (ii) As large surveys have pushed the surface brightness limits and detection efficiencies, many extremely faint satellite galaxies have been found orbiting the Milky Way. The completeness of current surveys is debated, and it has been argued that many hundreds of additional systems may be found in the coming years \\citep{2008arXiv0806.4381T}. Simulations that can resolve and follow the survival of substructure within $10\\%$ of $R_{\\rm vir}$ are necessary to compare with these data. (iii) Dark matter detection, either directly on Earth or indirectly via detection of annihilation relics, is the ultimate way to determine its nature. These experiments rely on accurate predictions for the phase-space structure of dark matter at the position of the Earth's orbit and the abundance and inner properties of substructure throughout the Galactic halo. (iv) Understanding the equilibrium structure resulting from violent relaxation is the ultimate challenge for galactic dynamicists. There is no compelling theory that can explain universal density and phase-space density profiles  \\citep{2001ApJ...563..483T}, or correlations such as between the local density profile and the anisotropy parameter \\citep{2006JCAP...05..014H}.  Given this motivation, we have carried out a sequence of simulations of a single Galactic mass dark matter halo, which at our highest resolution contains over a billion particles within its virial radius. In this letter we report on its inner structure and convergence properties.   \\begin{figure} \\label{pic} \\includegraphics[width=\\hsize]{bone} \\caption{The density of dark matter within the inner 200 kpc of GHALO$_2$. There are about 100,000 subhaloes that orbit within the virial radius. Each bright spot in this image is an individual, bound, dark matter subhalo made up of many thousands of particles (there are far more particles than pixels here).} \\end{figure} ",
        "conclusions": "The GHALO$_2$ simulation has achieved an unprecedented spatial and mass resolution within a CDM halo, resolving thousands of subhalos within a radius corresponding to the galactic disk and a rich phase-space structure of streams beyond a radius of $\\sim 8$ kpc. Whilst there are more detailed analyses of this simulation in progress, we have reported here on the global inner properties of density and phase-space density profiles and halo shape. Using a sequence of simulations of the same halo at difference resolutions, from $10^5$ -- $10^9$ particles, we confirm that the convergence radii for the density profile and shape scales as $N_{\\rm vir}^{-1/3}$. The logarithmic slope of the radial density profile is close to a power law, gradually turning over to a slope of $-0.8$ at our innermost resolved region ($0.05\\%$ of $R_{\\rm vir}$). We have proposed a new two parameter fitting function that has a linearly varying logarithmic gradient which provides the best fit to the inner part of the GHALO and VL2 haloes. A larger sample of haloes, such as \\citet{2001MNRAS.321..559B} and \\citet{2007MNRAS.378...55M}, would be required to determine if this functional form provides a universal fit. We find that the convergence radius of the density is a factor of three smaller than the convergence of halo shape. GHALO is prolate, yet becomes spherical within a region where orbits are most likely innacurately followed due to the effects of finite particle number, relaxation and softening.  All functional forms fit to density profiles, whether 2 or 3 parameters are empirical fits, even those based on properties (the Dehnen-McLaughlin) of the phase-space density profile whose origin is still poorly understood. Therefore the only current confidence can be given to those profiles which have been fit to the highest resolution simulations and over the widest range of halos encountered in N-body simulations. Clearly these two criteria are in conflict since simulations at the resolution of GHALO are too expensive to allow a broader study. Therefore, the results presented here should be considered as guides only, whose generality remains to be tested. Never the less we can consider economy of parameters and simplicity of functional form as guiding principles in the search for suitable profile functions to describe the end state of gravitational collapse. All the profiles we fit here (Table~1 and residuals in Figure~2) meet these subjective criteria, having at most 3 free parameters and simple functional forms.  While the phase-space density estimated by $\\rho\\sigma^{-3}$ is observed to follow a power law in radius of slope $-1.84$, its meaning is less clear since at small radii it is limited by resolution of the estimator and at larger radii it becomes dominated by subhalos. Using the more sophisticated EnBiD PSD estimator we find that the radial profile is steeper with an index of about $-2$, but that it is not as perfect a power-law as seen in $\\rho\\sigma^{-3}(r)$.  As a final comment, we note that in large galaxies, the inner structure and shape of the dark matter halo has likely been altered over time by the baryons via a range of physical effects, including dissipation, energy transfer from sinking massive objects, binary black holes, bar-halo interactions, turbulent gas motions and more. Simulations that follow the baryonic components together with the dark matter will resolve these additional questions in the coming years."
    },
    "0808/0808.2386_arXiv.txt": {
        "abstract": "{TW Hya is a classical T Tauri star that shows significant radial-velocity variations in the optical regime. These variations have been attributed to a 10 M$_\\mathrm{Jup}$ planet orbiting the star at 0.04\\,AU. } {The aim of this letter is to confirm the presence of the giant planet around TW Hya by (i) testing whether the observed RV variations can be caused by stellar spots and (ii) analyzing new optical and infrared data to detect the signal of the planet companion. } {We fitted the RV variations of TW Hya using a cool spot model. In addition, we obtained new high-resolution optical \\& infrared spectra, together with optical photometry of TW Hya and compared them with previous data. } {Our  model shows that a cold spot covering 7\\% of the stellar surface and located at a latitude of 54$^\\circ$ can reproduce the reported RV variations. The model also predicts a bisector semi-amplitude variation $<$ 10 m/s, which is less than the errors of the RV measurements discussed in Setiawan et al. (2008). The analysis of our new optical RV data, with typical errors of 10 m/s, shows a larger RV amplitude that varies depending on the correlation mask used. A slight correlation between the RV variation and the bisector is also observed although not at a very significant level. The infrared H-band RV curve is almost flat, showing a small variation ($<$35 m/s) that is not consistent with the published optical orbit. All these results support the spot scenario rather than the presence of a hot Jupiter. Finally, the photometric data shows a 20\\% (peak to peak)  variability, which is much larger than the 4\\% variation expected for the modeled cool spot. The fact that the optical data are correlated with the surface of the cross-correlation function points towards hot spots as being responsible for the photometric variability. } {We conclude that the best explanation for the RV signal observed in TW\\,Hya is the presence of a cool stellar spot and not an orbiting hot Jupiter. } ",
        "introduction": "TW Hya (spectral type K7, M$_*$ =0.7$M_{\\sun}$, age$\\sim$10\\,Myr) is an extensively studied classical T Tauri star (CTTS) \\citep[e.g.][]{1983A&A...121..217R}. With a distance of 56$\\pm$7 pc \\citep{Esa1997}, it is probably one of the closest laboratories for studying of our own solar system in its early stages of formation.  Recently, \\citet[][SHL08, hereafter]{Seti2008} announced the discovery of a giant planet ($\\sim$10\\,M$_\\mathrm{Jup}$) orbiting TW Hya at a separation of 0.04\\,AU.  Their discovery is based on high-resolution optical spectroscopic observations:  the analysis of two independent datasets obtained with a difference of $\\sim$ 3 months shows a significant radial velocity (RV, hereafter) variation with an amplitude of 196$\\pm$61 m/s and a periodicity of 3.56$\\sim$0.02 days.  This, together with the lack of correlation between the RV variation and the cross correlation-function bisector (BIS), has been interpreted as proof of the existence of a planetary-mass object orbiting the star.  In this letter, we study whether the RV variations in TW Hya can be caused by stellar spots and not by a planetary-mass companion.  With this in mind, we have tried to reproduce SHL08 data with a cool spot model. In addition, we obtained almost simultaneous optical \\& near-infrared high resolution spectroscopy and optical photometry of the source in order to detect the planet signal with a different dataset. ",
        "conclusions": "The analysis of new optical and infrared spectra of TW Hya allows us to conclude that the best explanation for the observed RV signal is the presence of a cold spot and not a  hot Jupiter orbiting the star. First, we find a clear dependence of the optical RV amplitude on the CCF mask used and, second, the IR RV curve is almost flat, with a small RV scatter ($\\sim$35 m/s) that is inconsistent with the optical, orbital solution. In fact, we show that a simple cool spot model can reproduce the RV observations presented by SHL08. On the other hand, the photometric data are correlated with the surface of the CCF, which points towards a hot spot as the cause of the photometric variability, as reported in previous studies.   Our result shows that searching for planets around young stars is highly demanding and cumbersome. Only the concerted use of the different techniques and instruments presented here were able to clarify the nature of the signal observed in TW\\,Hya. It is, for instance, remarkable that no phase shifts were found in the RV signal between  different epochs (Sect. 3.1), which shows that very stable cool spots exist around CTTSs. This must be taken with care when searching for planets around this kind of objects."
    },
    "0808/0808.0260.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary { } % aims heading (mandatory) {Intrinsic absorption is a fundamental physical property to understand the evolution of active galactic nuclei (AGN). Here a sample of 1290 AGN, selected in the 2--10 keV band from different flux-limited surveys with very high optical identification completeness is studied.} % methods heading (mandatory) {The AGN are grouped into two classes, unabsorbed (type-1) and absorbed (type-2), depending on their optical spectroscopic classification and X--ray absorption properties, using hardness ratios. Utilizing the optical to X--ray flux ratios, a rough correction for the $\\sim8\\%$ redshift incompleteness still present in the sample is applied. Then the fraction of absorbed sources is determined as a function of X-ray luminosity and redshift.} % results heading (mandatory) {A strong decrease of the absorbed fraction with X-ray luminosity is found. This can be represented by an almost linear decrease from $\\sim80\\%$ to $\\sim20\\%$ in the luminosity range log L$_X$=42--46 and is consistent with similar derivations in the optical and MIR bands. Several methods are used to study a possible evolution of the absorption fraction. A significant increase of the absorbed fraction with redshift is found, which can be described by a power law with a slope $\\sim(1+z)^{0.62\\pm0.11}$, saturating at a redshift of z$\\sim$2. A simple power law fit $\\sim(1+z)^{0.48\\pm0.08}$ over the whole redshift iis also marginally consistent with the data.} % conclusions heading (optional), leave it empty if necessary {The variation of the AGN absorption with luminosity and redshift is described with higher statistical accuracy and smaller systematic errors than previous results. The findings have important consequences for the broader context of AGN and galaxy co-evolution. Here it is proposed that the cosmic downsizing in the AGN population is due to two different feeding mechanisms: a fast process of merger driven accretion at high luminosities and high redshifts versus a slow process of gas accretion from gravitational instabilities in galactic disks rebuilding around pre-formed bulges and black holes.} ",
        "introduction": "In recent years it has become obvious that supermassive black holes at the centers of galaxies must play an important role in the evolution of galaxies. There is a strong correlation between the black hole mass and global properties of its host galaxy spheroid, like the bulge luminosity (Kormendy \\& Richstone, \\cite{Kormendy1995}; Magorrian et al., \\cite{Magorrian1998}) and the stellar velocity dispersion, i.e. the M$_{BH}$-$\\sigma$ relation (Ferrarese \\& Merritt \\cite{Ferrarese2000}; Gebhardt et al. \\cite{Gebhardt2000}). At the same time, evidence is mounting that active galactic nuclei (AGN) and galaxies in general undergo very similar evolution patterns. The peaks of AGN activity and star formation occur in the same redshift range (z = 1.5-2) and there is a similar dramatic decline towards low redshift. Moreover, the mass density of local dormant supermassive black holes in galaxy centers is consistent with the mass density accreted by AGN throughout the history of the Universe (Marconi et al. \\cite{Marconi2004}; Merloni \\cite{Merloni2004}), yielding further evidence for a tight link between nuclear black hole activity and the growth of galaxy bulges. Many recent theoretical models propose that feedback from the growing supermassive black holes plays an important role in linking the properties and evolution of central black holes and their host galaxy (Silk \\& Rees, \\cite{Silk1998}; Di Matteo et al., \\cite{DiMatteo2005}). The feedback from powerful AGN can quench star formation and inhibit the further growth of the most massive galaxies (Scannapieco \\& Oh, \\cite{Scannapieco2004}), and this effect is indeed required to match the results of hydrodynamical simulations to the observed bright end of the galaxy luminosity functions (Springel et al., \\cite{Springel2006}; Hopkins et al., \\cite{Hopkins2006}).  A strong dependence of the AGN space density evolution on X-ray luminosity has been found, the so called \"luminosity-dependent density evolution\" (LDDE) with a clear increase of the peak space density redshift with increasing X--ray luminosity, both in the soft X-ray (0.5--2 keV) and the hard X-ray (2--10 keV) bands (Miyaji et al., \\cite{Miyaji2000}; La Franca et al., \\cite{LaFranca2002}; Cowie et al., \\cite{Cowie2003}; Ueda et al., \\cite{Ueda2003}; Fiore et al., \\cite{Fiore2003}; Hasinger et al., \\cite{Hasinger2005}). This \"AGN cosmic downsizing\" evolution, which recently has also been confirmed in the radio and optical bands (Cirasuolo et al., \\cite{Cirasuolo2005}, Bongiorno et al., \\cite{Bongiorno2007}) is very similar to the \"downsizing\" of the faint galaxy population which shows a smooth decline of their maximum luminosity with decreasing redshift for z$<$1 (Cowie et al., \\cite{Cowie1996}). The observed \"cosmic downsizing\" indicates that active star formation and black hole growth shift to lower mass galaxies throughout the evolution of the Universe, which is somewhat counter-intuitive to the standard Cold Dark Matter scenario, where large scale structure evolves hierarchically from small to larger entities. The AGN results indicate a dramatically different evolutionary picture for low--luminosity AGN compared to the high--luminosity QSOs. While the rare, high--luminosity objects can form and feed very efficiently, possibly by multiple mergers rather early in the universe (see e.g. Li et al., \\cite{Li2007}), with their space density declining more than two orders of magnitude at redshifts below z=2, the bulk of the AGN has to wait much longer to grow or to be activated, with a decline of space density by less than a factor of 10 below a redshift of one. The late evolution of the low--luminosity Seyfert population is very similar to that which is required to fit the mid--infrared source counts and background (Franceschini et al. \\cite{Franceschini2002}) and also the bulk of the star formation in the Universe (Madau \\cite{Madau1996}), while the rapid evolution of powerful QSOs traces more closely the history of formation of massive spheroids (Franceschini et al. \\cite{Franceschini1999}). This could indicate two different modes of accretion and gas supply for the black hole growth with substantially different accretion efficiency (see e.g. the models of Cavaliere \\& Vittorini \\cite{Cavaliere2000}; Di Matteo et al., \\cite{DiMatteo2003}; Merloni \\cite{Merloni2004}; Menci et al., \\cite{Menci2004}).  A large fraction of the accretion in the Universe is obscured by intervening gas and dust clouds. Indeed, the spectrum of the X-ray background can very well be described by a combination of absorbed and unabsorbed AGN evolving through cosmic time (e.g. Comastri et al., \\cite{Comastri1995}; Gilli et al., \\cite{Gilli2007}, hereafter GCH07). However, one of the major uncertainties in the study of AGN evolution and black hole growth is the unknown distribution of absorbing column densities and its dependence on AGN luminosity and on cosmic time. Indeed, studies of local Seyfert 2 galaxies have shown that a large fraction of these ($\\sim$40\\%) is Compton thick (Risaliti et al., \\cite{Risaliti1999}) and thus practically absent in X--ray surveys. Therefore, the luminosity--dependent AGN evolution picture could be significantly biased by systematic selection effects (see e.g. Treister et al., \\cite{Treister2004}). Ueda et al. (\\cite{Ueda2003}) have for the first time found a significant decrease of the fraction of obscured AGN with increasing X--ray luminosity, and similar results have been obtained almost simultaneously from independent X--ray selected AGN samples by Steffen et al. (\\cite{Steffen2003}) and Hasinger (\\cite{Hasinger2004}). Recently, these trends have been confirmed in optically selected AGN observed in the SDSS (Simpson \\cite{Simpson2005} and with {\\em Spitzer} in the MIR (Maiolino et al., \\cite{Maiolino2007}; Treister et al. \\cite{Treister2008}), indicating a break-down of the standard AGN unification model (Antonucci \\cite{Antonucci1993}). The possible evolution of the fraction of obscured sources with redshift is still a matter of debate in the literature. A tentative redshift dependence of the obscured fraction was reported by La Franca et al. (\\cite{LaFranca2005}). Other authors (e.g. Ueda et al. \\cite{Ueda2003}, GCH07) did not find a significant redshift dependence. Recently, Treister and Urry (\\cite{Treister2006}) claimed a shallow increase of the absorbed fraction with redshift. The different results could indicate both systematic and statistical effects in the analysis.  In this paper the best available AGN samples selected in the 2-10 keV X-ray band with the highest redshift completeness possible have been compiled, resulting in 1406 objects, of which more than 92\\% have redshift information, mainly from optical spectroscopy. This allows for the so far best combination of object statistics and completeness in any hard X-ray selected AGN sample. In section 2 the hard X-ray selected AGN sample is described in detail. Section 3 discusses the classification criteria used to discriminate between type-1 (unabsorbed) and type-2 (absorbed) AGN. Section 4 presents the number counts of the different source populations and section 5 elaborates on a method to correct for the small, but significant redshift incompleteness in the sample. Section 6 is the main body of the paper and discusses the luminosity- and redshift dependence of the absorbed AGN fraction. Finally, the results are discussed in section 7 and summarized in section 8. Throughout this work a cosmology with $\\Omega_m$=0.3, $\\Omega_\\Lambda$=0.7 and H$_0$=70 km s$^{-1}$ Mpc$^{-1}$ is used.   \\begin{table*}[ht] \\begin{center} \\caption[]{The hard X--ray sample} \\begin{tabular}{@{}lccccccc@{}} \\hline Survey       &Solid Angle& $S_{lim}$&$N_{\\rm tot}$&$N_{\\rm unid}$&$N_{\\rm zph}$&$N_{\\rm AGN1}$&$N_{\\rm AGN2}$ \\\\ & [deg$^2$] &[erg cm$^{-2}$ s$^{-1}$] &               &                &               && \\\\ \\hline HEAO1/Grossan    & 2658--22000  & (1.0--2.7)   $\\cdot$ 10$^{-11}$ &   50 &   0 &   0 &  41 &   9\\\\ AMSS/ALSS    & $\\sim$4.8--90.8  & (0.11--2.1)  $\\cdot$ 10$^{-12}$ &  139 &   1 &   0 &  88 &  37\\\\ HBSS         & 25.17            & 2.15 $\\cdot$ 10$^{-13}$         &  67 &  2 &   0 & 42 &  20\\\\ XMS          & 2.46--3.80       & (3.30--9.36) $\\cdot$ 10$^{-14}$ &  144 &  13 &   0 & 111 &  25\\\\ CLASXS       & 0.4              &          1.3 $\\cdot$ 10$^{-14}$ &  112 &  19 &   0 &  67 &  41\\\\ HELLAS2XMM   & 0.21--0.90       & (1.0--3.7)   $\\cdot$ 10$^{-14}$ &  171 &  27 &   0 &  115 &  51\\\\ SEXSI        & 0.046--0.979     & (0.51--5.72) $\\cdot$ 10$^{-14}$ &  252 &  40 &   0 & 151 &  90\\\\ LH/XMM       & 0.1056           & 5.1          $\\cdot$ 10$^{-15}$ &   64 &   4 &  11 &  40 &  20\\\\ CDF--S       & 0.012--0.0873    & (0.44--1.49) $\\cdot$ 10$^{-15}$ &  239 &   4 &  66 &  80 & 131\\\\ CDF--N       & 0.0075--0.0873   & (0.18--1.49) $\\cdot$ 10$^{-15}$ &  284 &   8 & 111 &  99 & 141\\\\ \\hline Total        &                  &                                 & 1406 & 112 & 188 & 759 & 531\\\\ \\hline \\end{tabular}\\label{tab:samp} \\end{center} \\end{table*} ",
        "conclusions": "\\begin{enumerate} \\item A meta--sample of hard X--ray selected AGN has been compiled from the literature, providing an unprecedented combination of statistical quality and spectroscopic as well as photometric redshift completeness. Systematic differences in the selection of this rather heterogeneous set of samples were taken care of as far as possible and should not significantly bias the results.  \\item The small, but significant redshift incompleteness could be corrected for using the correlation between luminosity and X--ray to optical flux ratio observed in non-broad-line AGN. A comparison between the analysis with and without the inclusion of crude redshifts carried throughout the paper shows that the crude redshifts tend to enhance the high-redshift high-luminosity bins, but the main results do not depend strongly on this choice.  \\item The X--ray selected AGN could be classified using both optical spectroscopy and X--ray hardness ratios. This classification scheme appears more robust than either optical or X--ray classification alone.  \\item A strong decrease of the fraction of absorbed AGN was found with redshift. This trend confirms similar results obtained previously in the X-ray, but also optical and MIR bands. As expected, systematic selection effects are present between different bands, but do not dominate the results. This signals a break--down of the strong unified AGN model and indicates that high--luminosity AGN can clean out their environment.  \\item The same decreasing trend with luminosity is found in all individual redshift shells up to z$\\sim$3. However, the data indicate a significant evolution of the absorbed AGN fraction, which increases by about a factor of 2 up to z$\\sim$2 and remains approximately constant thereafter.  \\item This rather shallow evolution implies that the luminosity--dependent AGN density evolution observed in the X--ray, optical and radio bands is a real property of the parent population and not caused by systematic selection effects.  \\item  A new scenario is proposed, in which the cosmic down-sizing of the AGN population is due to two different fuelling mechanisms, on one hand the efficient growth of luminous QSOs by galaxy mergers early in the universe, on the other hand the re-juvenation of pre-formed bulges and black holes by slow gas accretion over cosmic time.  \\end{enumerate}"
    },
    "0808/0808.1100_arXiv.txt": {
        "abstract": "Integral field spectroscopy of galaxies at redshift $z\\sim2$ has revealed a population of early-forming, rotationally-supported disks. These high-redshift systems provide a potentially important clue to the formation processes that build disk galaxies in the universe. A particularly well-studied example is the $z=2.38$ galaxy $\\BzKgal$,  which was shown  by  \\cite{genzel2006a} to be a rotationally supported disk despite the fact that its high star formation rate and short gas consumption timescale require a very rapid acquisition of mass. Previous kinematical analyses have suggested that $z\\sim2$ disk galaxies like $\\BzKgal$ did not form through mergers because their line-of-sight velocity fields display low levels of asymmetry. We perform the same kinematical analysis on a set of simulated disk galaxies formed in gas-rich mergers of the type that may be common at high redshift, and show that the remnant disks display low velocity field asymmetry and satisfy the criteria that have been used to classify high-redshift galaxies as disks observationally. Further, we compare one of our remnants to the bulk properties of $\\BzKgal$ and show that it has a star formation rate, gas surface density, and a circular velocity-to-velocity dispersion ratio that matches $\\BzKgal$ remarkably well. We suggest that observations of high-redshift disk galaxies like $\\BzKgal$ are consistent with the hypothesis that gas-rich mergers play an important role in disk formation at high redshift. ",
        "introduction": "\\label{section:introduction}  Understanding the formation of disk galaxies remains a primary but elusive goal in cosmology.  The standard scenario for disk formation entails the quiescent, dissipational collapse of rotating gas clouds within virialized dark matter halos \\citep{white1978a,fall1980a,blumenthal1984a}, and this picture successfully explains the angular momentum content, size, and kinematical structure of observed systems \\citep[e.g.,][]{mo1998a}.  However, CDM-based cosmological simulations of disk galaxy formation have had difficulty producing high angular momentum disks without large bulges \\citep[e.g.,][]{navarro1994a,navarro2000a}, almost certainly in part due to the fact that mergers are common in Cold Dark Matter (CDM) cosmologies \\citep[][and references therein]{stewart08a}. More recent simulations have faired better by including  strong feedback, which suppresses star formation and thereby reduces the collisionless heating associated with stellar mergers \\citep[e.g.,][]{abadi2003a,sommer-larsen2003a,robertson2004a,governato2004a}.  Motivated by these results, \\citet[][hereafter, R06a]{robertson2006a} presented a supplemental ``merger-driven'' scenario for the cosmological disk galaxy formation where extremely gas-rich mergers at high-redshift lead to rapidly-rotating remnant systems with large gaseous and stellar disks. Building on work by \\cite{barnes2002a}, \\cite{brook2004a}, and \\cite{springel2005a}, R06a used a suite of hydrodynamical simulations to study gas rich mergers of various gas fractions, mass ratios, and orbital properties, and concluded that hierarchical structure formation and the regulation of star formation may work in concert to build disk galaxies through gas-rich mergers. Subsequent cosmological simulations  have reported the formation of Milky Way-like disk galaxies in gas-rich mergers \\citep{governato2007a}, and \\cite{hopkins2008a} have examined merger remnants in order to gain a phenomenological understanding of disk survival in gas-rich mergers. Given that high gas fractions are the essential ingredient in this scenario, observations at high-redshift (when gas fractions were higher and mergers more common) provide an important testing ground for models of disk galaxy formation.   Only recently have detailed observations of high-redshift disks been possible. \\cite{erb2003a} measured the rotation curves of UV-selected $z\\sim2$ galaxies \\citep[the BM/BX sample,][]{adelberger2004a,steidel2004a} with $\\Ha$ slit spectroscopy \\citep[see also][]{erb2006a}. \\cite{forster-schreiber2006a} used the SINFONI integral field spectrograph \\citep{eisenhauer2003a} to measure spatially-resolved kinematics of 14 BM/BX galaxies, and found rotating systems with large specific angular momenta but significant random motions ($v/\\sigma\\sim2-4$).  Among the best-studied high-redhift disks is  $\\BzKgal$ at $z=2.38$, which was found using the $BzK$ photometric selection technique \\citep{daddi2004a}. \\citet[][hereafter G06]{genzel2006a} used the SINFONI spectrograph with  adaptive optics on the Very Large Telescope (VLT) to measure the kinematic structure $\\BzKgal$ with an incredible spatial resolution of 0.15\" ($\\sim1.2\\,\\kpc$ at $z \\simeq 2.4$). These observations revealed massive ($\\Mstar \\sim 8\\times10^{10}\\Msun$, $\\Mgas \\sim 4\\times10^{10}\\Msun$) rotating disk galaxy with a large star formation rate ($\\SFR\\sim140\\Msun\\yr^{-1}$) and substantial random motions ($v/\\sigma\\sim3$). The inferred gas consumption timescale for $\\BzKgal$ is only $\\sim285~\\Myr$, or about one-tenth the age of the universe at $z\\sim2.38$,  and implies a remarkably rapid acquisition of its mass.  Nonetheless G06 conclude that $\\BzKgal$ is not a merger remnant because it displays velocity fields similar to those expected for quiescent disks.   The existence of high-redshift disks that have not experienced large mergers may be surprising in light of $\\Lambda$CDM expectations. The merger rate per dark matter halo is predicted to be a factor of $\\sim15$ times higher at $z \\sim 2.4$ than at $z\\sim0$ \\citep{fakhouri2008a}. Approximately $10 \\%$ of galaxy-size halos at $z \\simeq 2.4$ should have undergone a nearly one-to-one merger ( $>1:1.25$)  in the last $\\sim500~\\Myr$, and  $\\sim 50 \\%$ should have experienced a $>1:3$ merger in the last $\\sim 1$ Gyr (Stewart, Bullock et al., in preparation). However, these mergers are expected to be gas rich.   Observations of galaxies at $z\\sim2$ imply gas fractions $\\sim10-20$ times higher than for nearby galaxies at fixed stellar mass \\citep{erb2006a,gavazzi2008a}.   These two pieces of information motivate us to consider gas-rich mergers as a means to explain $z\\sim 2$ galaxies with young stellar populations, short gas exhaustion timescales, and disk-like kinematics.  Any model that attempts to explain the properties of high-redshift disks must not only reproduce their bulk properties (stellar masses, rotational velocities, gas fractions, star formation rates, etc.) but also their detailed kinematic properties. An important kinematic metric for classifying high-redshift galaxies as disks was developed by \\citet[][hereafter S08]{shapiro2008a}, who extended the ``kinemetry'' technique of \\cite{krajnovic2006a} in order to examine the velocity fields of $z\\sim2$ galaxies from \\cite{forster-schreiber2006a} and G06. Specifically, S08 provided a straightforward means to classify $z\\sim2$ galaxies as disks based on the symmetry of their velocity fields.   The purpose of this {\\em Letter} is to determine whether galaxies formed in a gas-rich mergers can match the observed properties of $z \\sim 2$ disks. We describe our simulations in \\S \\ref{section:methodology} and the kinematical analysis in \\S \\ref{section:kinemetry}. In \\S \\ref{section:results} we demonstrate that a set of gas-rich merger remnants satisfy the disk galaxy kinematical classification developed by S08, and compare one of the remnants in particular to the detailed properties of  $\\BzKgal$. Overall, we find good agreement between the properties of simulated disk galaxies formed in gas-rich mergers and observed $z\\sim2$ disks like $\\BzKgal$. We assume a $\\LCDM$ cosmology with $\\Omega_{m}\\simeq0.3$, $\\Omega_{\\Lambda}\\simeq0.7$, and a Hubble parameter of $H_{0} = 70 \\km\\s^{-1}\\Mpc^{-1}$. ",
        "conclusions": "\\label{section:results}  The bottom row of Figure \\ref{fig:vel_field} shows the SMD, $\\vlos$, and $\\slos$ maps used to perform the kinematical analysis of the remnant disk.  The average kinemetric parameters $\\vasym$ and $\\sasym$ (Equation 2) are measured from the velocity maps.  For the disk merger remnant plotted in Figure \\ref{fig:vel_field} we find $\\vasym=0.076$ and $\\sasym=0.063$, which reflect the relative symmetry the velocity fields. The combined asymmetry parameter $\\Kasym\\equiv\\sqrt{\\vasym^{2} + \\sasym^{2}}$ provides a global characterization of asymmetry in the kinematical structure of the galaxy, and S08 suggest $\\Kasym<0.5$ as an observational criterion for distinguishing between disk galaxies and mergers.  The disk remnant examined in Figure \\ref{fig:vel_field} satisfies this criterion with $\\Kasym=0.1$, even though it was formed in an equal mass, gas-rich merger. We have performed the same analysis on a range of other simulated merger remnant disk galaxies from R06a. In each case, we examine the merger product soon after the merger and $150~\\Myr$ after the time of the final coalescence. We find similar results for disk remnants formed in gas-rich major polar mergers (R06a simulation GCoF, $\\Kasym=0.14$), major inclined mergers (R06a simulation GCoO, $\\Kasym=0.15$), and minor (1:8 mass ratio) coplanar mergers (R06a simulation FCm, $\\Kasym=0.16$). If real galaxy analogues of these simulated galaxies were observed with SINFONI, a kinematical analysis would likely classify them (correctly) as disks. We emphasize that in each case the gas disk develops rapidly, with $v/\\sigma$ and $\\Kasym$ both declining by factors of $\\sim2-3$ within $\\sim100~\\Myr$ of the height of merger coalescence.  While both the simulated disk merger remnants and observed $z\\sim2$ disk galaxies satisfy the observational criterion of S08 designed to identify disk galaxies, the simulations should also closely match other observed properties of $z\\sim2$ disks.  Consider the galaxy $\\BzKgal$ observed by  G06, which has an inferred star formation rate of $\\SFR\\approx140^{+110}_{-80}\\Msun \\yr^{-1}$ and a dynamical mass of $M_{\\mathrm{dyn}}=1.1 \\pm 0.1 \\times10^{11}\\Msun$ within a radius of $r\\lesssim 9\\kpc$.  The simulated disk remnant analyzed in Figure \\ref{fig:vel_field} has an average star formation rate of $\\SFR\\approx150 \\Msun \\yr^{=1}$ over the previous $100~\\Myr$ and a dynamical mass of $M_{\\mathrm{dyn}}=1.2\\times10^{11}\\Msun$.  The inferred stellar and gas masses for $\\BzKgal$ are $\\Mstar = 7.7^{+3.9}_{-1.3}\\times10^{10}\\Msun$ and $\\Mgas = 4.3\\times10^{10}\\Msun$, while the disk remnant has a stellar mass of $\\Mstar = 7.3\\times10^{10}\\Msun$ and a gas mass of $\\Mgas = 3.8\\times10^{10}\\Msun$.  Table \\ref{table:properties} compares the salient properties of $\\BzKgal$ and the simulated disk remnant, and they are similar in each case. Overall, the galaxy $\\BzKgal$ displays properties that are remarkably similar to the example simulated disk galaxy merger remnant shown in Figure \\ref{fig:vel_field}.  We note that the properties of $z\\sim2$ disks like $\\BzKgal$ are rather extreme compared with local, isolated disk galaxies. The gas surface density, star formation rate, and random motions ($v/\\sigma$) of $\\BzKgal$ are roughly an order of magnitude larger than for the Milky Way, even as their rotational velocities are similar and their stellar masses differ by less than 50\\% \\citep[e.g.,][]{blitz1980a,rana1991a,kent1991a,klypin2002a,levine2006a}. Our simulated disk remnant displays similar properties (see Table \\ref{table:properties}) owing to the dynamical effects of the merger from which it formed. It is possible that such a pre-formed (hot) disk would tend to settle into a more classical Milky-Way type thin disk at late times if the pre-formed stars adjust to to the infall of (quiescent) gas-disk material, as might be expected in a more classical stage of disk formation. Subsequent gaseous infall from the intergalactic medium may also help prolong star formation in the remnant and increase the gas disk scale length.  While the observed properties of $z\\sim2$ disks resemble simulated disk systems formed in gas-rich mergers, other viable explanations for their origin exist. Notably, \\cite{bournaud2007a,bournaud2008a} have suggested that high-redshift disk galaxies undergo a ``clump-cluster'' phase where large-scale gravitational instability leads to a turbulent, clumped gaseous distribution \\citep[see also][]{noguchi1998a,noguchi1999a}. An analysis of clump-cluster galaxy models using kinemetry method of \\cite{krajnovic2006a} and S08 would provide a useful test of this scenario.  As the quantity and quality of kinematical data for $z\\sim2$ disk galaxies improve, the observational constraints will better distinguish between various models for their formation.  These observations will be essential for determining how disk galaxies are assembled over the history of cosmic structure formation."
    },
    "0808/0808.0330_arXiv.txt": {
        "abstract": "We present a unified approach to neutrino processes in nucleon matter based on Landau's theory of Fermi liquids that includes one- and two-quasiparticle-quasihole pair states as well as mean-field effects. We show how rates of neutrino processes involving two nucleons may be calculated in terms of the collision integral in the Landau transport equation for quasiparticles. Using a relaxation time approximation, we solve the transport equation for density and spin-density fluctuations and derive a general form for the response functions. We apply our approach to neutral-current processes in neutron matter, where the spin response function is crucial for calculations of neutrino elastic and inelastic scattering, neutrino-pair bremsstrahlung and absorption from strongly-interacting nucleons. We calculate the relaxation rates using modern nuclear interactions and including many-body contributions, and find that rates of neutrino processes are reduced compared with estimates based on the one-pion exchange interaction, which is used in current simulations of core-collapse supernovae. ",
        "introduction": "Neutrino emission, absorption and scattering processes in nucleon matter play a crucial role for the physics of stellar collapse, supernova explosions and neutron stars~\\cite{Raffelt,Prakashreview}. Since the leptons in these processes interact weakly, the neutrino rates can be expressed compactly in terms of the response of nuclear matter to axial and vector probes. In many situations, the axial response is the more important, and in this paper we concentrate on this case, which for a system of nonrelativistic nucleons amounts to the spin or spin-isospin response. These responses have been calculated by a number of groups~\\cite{Sawyer,IP,Prakash1,Prakash2,BS} allowing for single nucleon quasiparticle-quasihole pair states.\\footnote{The basic single-particle-like excitations we work with are quasiparticles and quasiholes that have properties quantitatively different from those of free particles or holes. However, for brevity, we shall refer to these excitations simply as particles and holes.} However, this is insufficient for rates of neutrino processes involving two nucleons, such as neutrino-pair bremsstrahlung and absorption, and modified Urca reactions, in which two particle-hole pair states are necessary. The possible importance of two particle-hole pair states for neutrino inelastic scattering, in particular for energy exchange and the formation of the neutrino spectra, has been emphasized by Raffelt {\\it et al.}~\\cite{Raffelt1,Raffelt2,Raffelt3}. Bounds on the magnitude of the two particle-hole pair weight have been investigated in Ref.~\\cite{OP} and it has been shown how the two-pair response is directly related to the collision term in Landau's transport equation for quasiparticles~\\cite{LOP}.  Noncentral contributions to nuclear interactions, such as tensor forces from pion exchanges and spin-orbit forces, are essential for the two particle-hole pair response, as is clear from calculations of neutrino-pair bremsstrahlung and the modified Urca processes~\\cite{FM} and from general considerations based on conservation laws~\\cite{OP}. Neutrino-pair bremsstrahlung and absorption change the number of neutrinos and are key for equilibrating muon and tau neutrino number densities in supernovae. The standard rates for bremsstrahlung are based on the one-pion exchange model for nucleon-nucleon interactions~\\cite{FM} (in the context of supernovae, see for example Ref.~\\cite{Raffelt3}). This is a reasonable starting point, since it represents the long-range part and the leading noncentral contribution in chiral effective field theory for nuclear forces~\\cite{chiralEFT}. However, the tensor force from pion exchange is singular at short distances, which in free space requires iteration in the spin-triplet channels~\\cite{tensor}. In addition, subleading noncentral contributions to nuclear interactions are important for reproducing nucleon-nucleon scattering for the relevant channels and energies~\\cite{Epelbaum}.  The aim of this paper is to give a unified treatment of neutrino processes that includes one- and two-particle-hole pair states as well as mean-field (Fermi liquid) effects consistently, and to present improved rate calculations of these processes based on modern nuclear interactions beyond one-pion exchange and including many-body contributions. A convenient framework for doing this is Landau's theory of normal Fermi liquids. This work represents an extension of Ref.~\\cite{LOP}, which included two particle-hole pair states only in leading order using diagrammatic perturbation theory. Here we shall use the quasiparticle transport equation. This provides a useful framework for understanding the basic physics and for making detailed calculations. In this paper, we focus on neutral-current processes in normal (nonsuperfluid) neutron matter. We leave for future work the application to mixtures of neutron and protons, charged-current reactions, and the extension to superfluid phases.  This paper is organized as follows. Section~\\ref{nustruct} gives an introduction to neutrino processes and the dynamical structure factors. In Sect.~\\ref{FLTtrans}, we discuss Landau Fermi-liquid theory, show that it represents a useful effective theory for neutrino processes in nucleon matter, and introduce the transport equation for quasiparticles. Using a relaxation time approximation, we solve the transport equation for density and spin-density fluctuations and derive a general form for the response functions in Sect.~\\ref{relax}. The response function includes contributions from one-particle-hole pair (corresponding to elastic scattering of neutrinos from nucleons) and two-particle-hole pair states (which enter calculations of inelastic scattering, and neutrino-pair bremsstrahlung and absorption). In Sect.~\\ref{times}, we calculate the appropriate relaxation times for the one-pion exchange interaction and for a general operator representation of the quasiparticle scattering amplitude. We present results in Sect.~\\ref{results} based on modern nuclear interactions and including many-body contributions, and contrast these with rates obtained using the one-pion exchange interaction, which is typically used in supernova simulations. Finally, we assess the significance of the improved treatment of nuclear interactions for neutrino mean free paths, energy loss and energy transfer in supernovae. We summarize the improvements and conclude in Sect.~\\ref{concl}. ",
        "conclusions": "\\label{concl}  We have developed a unified treatment for neutrino processes in nucleon matter based on Landau's theory of Fermi liquids that includes one- and two-particle-hole pair states consistently. The contributions from two-particle-hole pair states are crucial for neutrino-pair bremsstrahlung and absorption, for inelastic scattering, modified Urca reactions, and axion emission. In supernovae, neutrino-pair bremsstrahlung and absorption dominate the neutrino-number changing reactions and are key to the production of muon and tau neutrinos  Neutrino rates involving two nucleons can be calculated in terms of the collision integral in the Landau transport equation for quasiparticles. Using a relaxation time approximation, we have solved the transport equation for density and spin-density fluctuations and derived a general form for the response functions. The solution includes multiple-scattering effects and effects due to non-zero wavelengths and recoil of the nucleons. We have applied our approach to neutral-current processes in neutron matter, but the generalization to isospin is straightforward. Our results for the spin response are summarized by Eqs.~(\\ref{structspin}), (\\ref{Xsigma}), (\\ref{imchi}), (\\ref{imchi0}), (\\ref{tauavC}) and the values of $C_\\sigma$ of Fig.~\\ref{fig:spin}.  We have calculated the relaxation times based on the OPE model and for a general representation of the quasiparticle scattering amplitude. For OPE, the spin relaxation rate is comparable to the quasiparticle relaxation rate, $\\tau/\\tau_\\sigma = 4/3$. This highlights the importance of noncentral contributions to nuclear interactions. We therefore performed more systematic calculations of these rates. In addition, for $|\\om| \\tau_\\sigma \\gg 1$ and in the long-wavelength limit, our result for the dynamical structure factor agrees with Raffelt {\\it et al.}~\\cite{Raffelt1,Raffelt2}.  Beyond OPE, we have calculated the relaxation times based on low-momentum interactions $\\vlowk$ and including second-order many-body contributions. The effects of three-nucleon interactions are generally weaker in neutron matter~\\cite{neutmatt}, but need to be included in future work. The OPE model significantly overestimates the strength of noncentral contributions, compared to low-momentum interactions $\\vlowk$, for all considered densities. Beyond the $\\vlowk$ results, we have found that second-order many-body contributions reduce the spin relaxation rate especially at lower densities. This provides a range in Figs.~\\ref{fig:spin} and~\\ref{fig:dens} for the effects due to many-body correlations. By using spin relaxation times that incorporate both ``in-scattering'' and ``out-scattering'' terms in the transport equation, effects corresponding to vertex corrections in the microscopic theory are automatically taken into account.  Using the spin response in the long-wavelength limit, but without further approximations or Ansaetze for the structure factor, we have estimated the significance of the improved rates for neutrino mean free paths, energy loss and energy transfer. We have found a reduction of these rates using modern nuclear forces compared to OPE and consequently weak mean-field effects. In addition, for all cases we find that the energy transfer due to inelastic scattering is not significantly larger than that due to neutrino-pair bremsstrahlung and absorption.  One may ask how good the relaxation time approximation is. Our choice of spin relaxation time is designed to agree with microscopic theory in the collisionless limit, $|\\omega| \\tau_\\sigma \\gg 1$, and at long wavelengths. For the hydrodynamic limit, $|\\omega| \\tau_\\sigma \\ll 1$, and long wavelengths, exact solutions of the transport equation have been obtained, and one finds~\\cite{exact1,exact2,BaymPethick} \\be \\frac{{\\tau_\\sigma}|_{\\rm hydro}}{\\tau} = \\frac{4}{3} \\sum_{\\nu=1,3,5,\\ldots} \\frac{2\\nu+1}{\\nu(\\nu+1)[\\nu(\\nu+1)-2+2 \\, \\tau/\\tau_\\sigma]} \\,. \\ee The relaxation time for the hydrodynamic limit is always greater than or equal to that for the collisionless limit. For $\\tau_\\sigma/ \\tau \\gg 1$, ${\\tau_\\sigma}|_{\\rm hydro}=\\tau_\\sigma$, while for $\\tau_\\sigma/\\tau = 1$, ${\\tau_\\sigma}|_{\\rm hydro} = (\\pi^2/9) \\, \\tau_\\sigma$. Since for realistic nuclear interactions, $\\tau_\\sigma$ is significantly larger than $\\tau$, this indicates that differences between spin relaxation times in the collisionless and hydrodynamic limits are expected to be of order a few per cent. Consequently, uncertainties due to the use of the relaxation time approximation are small compared with other uncertainties in the calculation.  The use of the quasiparticle transport equation with a collision term allows us to include some two-particle-hole pair states, but not all. Among contributions not included are terms that correspond to the incoherent parts of the propagator for a particle-hole pair, that is to contributions that do not correspond to an intermediate state containing a well-defined quasiparticle together with a well-defined quasihole. Moreover, there are intrinsic two-body contributions to hadronic weak currents. Further work is needed to determine how important these additional contributions are.  There are numerous directions for future work. One is to explore mixtures of neutrons and protons. A second is to extend the calculations to situations when matter is less degenerate. As one sees from our results, there is significant uncertainty in the effects of the medium on quasiparticle scattering amplitudes, since there are sizable differences between rates obtained with $\\vlowk$ and those that include many-body contributions to second order, and an important task is to reduce these uncertainties."
    },
    "0808/0808.2790_arXiv.txt": {
        "abstract": "It is generally considered, as a rule of thumb, that carbon monoxide forms very early in envelopes of AGB stars, and that it consumes most of the carbon, or most of the oxygen, depending on whether  the photosphere is oxygen-rich or carbon-rich, respectively. \\rm This work focuses on the latter case , with the purpose of quantifying the remaining fraction of gaseous carbon which is then available for forming carbonaceous grains. Since AGB stars are (probably) the main providers of cosmic carbon grains, this residual fraction is essential in establishing the validity of current grain models. Here, we use a kinetic treatment to follow the chemical evolution of circumstellar shells towards steady state. It is shown that the residual fraction depends essentially on the atomic ratio of pristine gaseous carbon and oxygen, and on the cross-section for CH formation by collision of C and H atoms. It lies between 55 and 144 C atoms per $10^{6}$ H atoms, depending on the values adopted for unknown reaction rates and cosmic C abundance. This is much larger than predicted by the rule of thumb recalled above.  The present results depend strongly on the rate of the reaction C+H$\\rightarrow$CH: far from thermodynamic equilibrium (which is the case here), CO cannot be formed if this rate is as low as generally assumed. We have, therefore, estimated this rate by chemically modelling the reaction and found it indeed much higher, and high enough to yield CO abundances compatible with observations. An accurate experimental rate determination is highly desirable. ",
        "introduction": "The spectral analysis of the CS (CircumStellar) envelopes of AGB stars has been an active field of observational astronomy for decades (see, for instance, Bujarrabal et al. \\cite{buj}, Woods et al. \\cite{woo}). As a result, the list of detected atoms, ions, radicals and molecules is continuously enriched and their characterization improves in extent and quality as instruments and telescopes are constantly upgraded (see van Dishoek \\cite{vdi} for references).  One of the most abundant molecules in both O-rich and C-rich AGBs is carbon monoxide, CO, so much so that the dominant isotopic variety, $^{12}$CO, is optically thick in its characteristic infrared bands, precluding direct determination of its column density towards the star. One way of circumventing this difficulty is to measure instead the density of the rarer isotope, $^{13}$CO, from which that of $^{12}$CO is deduced, assuming their ratio is somewhere between 1/50 and 1/100, because of (undetermined) fractionation processes which locally alter the relative cosmic abundance of 1/89 (see Burgh at al. \\cite{bur}, Liszt \\cite{lis}).  In spite of these achievements and because of observational uncertainties (such as the CO isotopic ratio) and the relative scarcity of stars amenable to in-depth analysis, it is still difficult to draw a complete and accurate observational picture of the variations of CO abundance with the type of AGB star.  For want of a better guess, one may assume an arbitrary but reasonable relative abundance, such as $[CO]/[H_{2}]=10^{-3}$ (e.g. Woods \\cite{woo})). Another simple rule of thumb has often been used: all the oxygen (respectively carbon) expelled by C-rich (respectively O-rich) stars is locked into gaseous CO. This rule is strongly suggested by the high reciprocal chemical affinities of the two elements, and the strength of the CO bond.  Precisely because of the latter, the carbon and oxygen atoms locked in CO are irreversibly lost for carbonacious and silicate grains which ultimately condense in the envelope. If strictly applied, the rule above implies that only stars with C/O$>$1 can contribute to carbon grains. Now, for many carbon grain models, the C budget is already quite tight, so the available C input deserves careful scrutiny. In particular, the carrier of the extinction bump at 2175 $\\AA{\\ }$ (see, for instance, Bless and Savage \\cite{ble}) requires large quantities of pure carbon (see Snow and Witt \\cite{sno} for a discussion). Since CS shells, together with supernovae and novae, where grain condensation probably occurs in similar ways, are the main, if not sole, genitors of grains, accurate and  generally applicable  procedures are required to estimate the amount of C available for grain condensation.  On the theoretical side, a parallel effort was made to model the chemistry which takes place in these envelopes, using long chains of known processes and reactions. Assuming local thermodynamical equilibrium (LTE) has been a powerful approach for estimating the relative abundances of the various species of interest as a function of temperature, pressure and distance from the central star (see Tsuji \\cite{tsu}, Salpeter \\cite{sal}, Duley and Williams \\cite{dul}). However, for LTE to be applicable, steady states at high densities and temperatures are required, so that detailed balance may be invoked, which is not the case in CS. Besides, it is not easy to include solid grains in the LTE scheme.  Another approach consists in selecting a number of molecular species and a number of known reactions between these species, and expressing the rate of change of each species as a function of the instantaneous densities of all species and the rates of their reactions. In general, a steady state is quickly established, which depends on local temperature and initial densities. The complexity of the problem is illustrated by the work of Prasad and Huntress \\cite{pra}, a milestone in the study of the chemistry of intersellar molecular clouds. While the physical conditions in clouds are not comparable with those of CS envelopes, that work is an inspiring example for our present purposes. Cherchneff \\cite{che06} recently used the kinetic approach to study the dependance of molecular populations in the inner parts of CS of AGB stars, as a function of the C/O ratio in the photosphere. The present research is also based on chemical kinetics rather than LTE, but the goal is much less ambitious, being limited to estimating the fraction of carbon available for grain formation.  \\rm We consider carbon grains and gaseous CO to be competitors for the available C atoms from the photosphere. A small number of chemical reactions between the most abundant gaseous chemical species are retained to define the two corresponding routes for pristine C atoms (Sec. 2). They translate into a set of differential equations which can be solved numerically for given relative abundances of C and O, to yield the relative number of C atoms that end up in grains and CO respectively (Sec. 3). The results are arranged as a function of the ratio C/O, displaying continuous trends from O-rich to C-rich stars. The fraction of carbon available for grains is thus predicted to be overall much larger than present estimates (Sec. 4).  The discussion of these results (Sec. 5) singles out the reactions which dominate the production of CO or carbonaceous grains, and should therefore be more fully documented by laboratory measurements. Finally, we discuss how the average ratio of carbon atoms available for grains to total H atoms can be predicted based on the present computations. ",
        "conclusions": "\\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig6.eps}} \\caption[]{An upper limit to the fraction of oxygen available for silicate grains, $n_{O}(f)/n_{O}(0)$.} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig7.eps}} \\caption[]{Effects of changes in the parameters, one at a time, defined by the rank, N. N=1: k1=$0.1\\,\\,10^{-10}$. N=2: k2=0; N=3: length of run multiplied by 2. N=4: $n_{OH}=0$ at all times. N=5: all rates as in Table 1. N=6: $n_{H2}(0)/n_{H}=0.1$. N=7: k1=$2\\,\\,10^{-10}$. N=8: k1=$6.6\\,\\,10^{-10}$.N=9: $a_{316}=0$. N=10: $a_{241}+a_{242}=0$. C/O=1 for all.} \\end{figure}     Back in Fig. 4 and 5, dotted lines have been drawn to illustrate the application of the rule of thumb recalled in the Introduction. As expected, the latter approximation is seen to agree with our more elaborate calculations only in the two opposite limits of $n_{C}(0)/n_{O}(0)$. The difference is considerable in between, which is the most often encountered range in the sky (see for instance Kwok et al. \\cite{kwo}).   In a further effort to gauge the relevance of our chemical scheme, we sought to reduce the number of intermediate species to the strict minimum. Only CH was found to be indispensable for the formation of both CO and grains. The corresponding reduced scheme is \\begin{eqnarray*} & &\\dot n_{C}=-a_{63}n_{C}n_{CH}-k_{1}\\,n_{C}n_{H}+a_{1}n_{H}n_{CH};\\\\ & &\\dot n_{O}=-a_{140}n_{O}n_{CH};\\\\ & &\\dot n_{Cgr}=2k_{2}n_{CH}n_{CH}+k_{63}n_{CH}n_{C};\\\\ & &\\dot n_{CH}=-a_{1}n_{H}n_{CH}-a_{63}n_{C}n_{CH}-a_{140}n_{CH}n_{O}\\\\ & &-2k_{2}n_{CH}n_{CH}+k_{1}n_{C}n_{H};\\\\ & &n_{CO}=n_{O}(0)-n_{0}.\\\\ \\end{eqnarray*}   The results are represented in Fig. 4 and 5 by open circles. These results are clearly incorrect considering their behaviour at the two ends of the X-axis. However, one notes that  $n_{CO}(f)/n_{C}(0)$ computed with this scheme is never less than one half of the value delivered by the more elaborate scheme, and the strings of squares and circles are closely similar. This is witness to the robustness of the first scheme.  This was further tested by performing other runs with the first scheme, only changing one or two parameters. Figure 7 plots $n_{CO}(f)/n_{C}(0)$ (open squares) and $n_{Cgr}(f)/n_{C}(0)$ (filled squares), for some of these cases, all with $n_{C}(0)=n_{O}(0)=3\\,\\,10^{6}$ cm$^{-3}$. The integer abscissa refers to the sentence in the legend which describes the corresponding changes applied to the parameters. The results are arranged so as to higlight the opposite effects of various changes in the chemical scheme, with respect to the ``standard\" case, corresponding to the equations of Sec. 3. Some elements of the scheme have no great impact on the final values of interest here; for instance, the inclusion of OH (N=3) and the initial fraction of molecular hydrogen (N=5). By contrast, the association of C and H to form CH (reaction k1) is essential for a high yield of carbon grains, as stated above and as evidenced by the steep reduction in carbon grain formation when k1 is reduced to $0.1\\,\\,10^{-10}$ (N=1). Other parameters have moderate impact.  Note that, in case N=1, the yield of CO has not diminished, and even increased somewhat, although CH is an important factor of CO production through reaction 140 (Table 1). This is because other quite efficient, reactions are still at play, which compensate the blocking of that route. However,  the steady state is retarded by a factor 10.  Other computations also showed that changing all initial atom densities in the same proportions by a factor 2 had only marginal effects on the final yields, but retarded or advanced the setting of steady state.    Finally, let us consider the problem of carbon grain formation. The fraction of cosmic carbon that is available for grain formation is usually equated to the difference between photospheric (cosmic) and ISM abundances of carbon, divided by the former. In our treatment, this is represented by the average of $n_{Cgr}(f)/n_{C}(0)$ over all stars. This quantity may be deduced, with the help of Fig. 5, from the measurement of the photospheric C/O ratio of a large number of stars of different types. $\\emph {If the stars were evenly distributed over the range}$  0.5$<$C/O$<$2, the fraction would be $\\sim$0.25. If, furthermore, the cosmic C/H ratio is assumed to be that recommended by Snow and Witt \\cite{sno}, $2.25\\,\\,10^{-4}$, then, the available number of C atoms for carbon grains becomes 55 for $10^{6}$ H atoms.  However, this must be considered a lower limit, as it is the result of calculations with k1=$1\\,\\,10^{-10}$ cm$^{3}$s$^{-1}$, a reasonable but arbitrary value, adopted in Table 1, in the absence of experimental data. If, on the other hand, one takes k1=$6.6\\,\\,10^{-10}$, as suggested by our chemical modeling (see Appendix), and if, moreover, one adopts the Grevesse cosmic carbon abundance for the Sun \\cite{gre}, i.e. 4$\\,\\,10^{-4}$, then  the number of available C atoms rises to 144.  Circumstances may be envisioned where this upper limit may still be raised to some extent. For instance, as the expanding gas reaches the outskirts of the envelope, it is no longer shielded from the UV radiation of the ISM. If the latter is strong enough, CO may then be dissociated, making some more free C atoms available for grain formation, as observed by Herpin et al. \\cite{her} in the case of post-AGB stars.  Another instance is a high initial ratio of molecular to atomic hydrogen. Following a suggestion by the reviewer of this paper, our program was run (after minor changes) for the case where $n_{H}(0)=n_{H2}(0)=10^{10}$cm$^{-3}$ and $n_{C}(0)=n_{O}(0)=3\\,\\,10^{6}$ cm$^{-3}$. Comparing with Fig. 1 to 3, the final value of $n_{Cgr}$ was found to increase to 1.2$\\,\\,10^{6}$ from 0.75$\\,\\,10^{6}$ cm$^{-3}$, and the final value of $n_{O}$ to decrease to $1.8\\,\\,10^{6}$ from $2.24\\,\\,10^{6}$ cm$^{-3}$. This notable increase in grain abundance in the presence of a high density of H$_{2}$ is due to reactions 48 and 237 which generate CH$_{2}$ and C$_{2}$H$_{2}$, the progenitors of grains.  \\rm The range 55 to 144 for $10^{6}$ H atoms, deduced from our model, is consistent with spectroscopic observations of a couple of stars as reported by Snow and Witt \\cite{sno}, in a compilation of measurements. While more astronomical observations could help, the present work should be an encouragement to further study simple chemical reactions between neutrals, especially reaction k1, so as to restrict the range of uncertainty of the model. Such reactions have become more amenable to measurement with the development of specialized apparata, e.g. the CRESU supersonic-flow machine (see for instance Rowe \\cite{row})."
    },
    "0808/0808.2273_arXiv.txt": {
        "abstract": "We have obtained a Gemini South T-ReCS \\bQa-band (18.3 $\\mu$m) image and a \\emph{Spitzer} MIPS SED-mode observation of \\targst, an F5/F6V member of the $\\sim$12 Myr old $\\beta$ Pictoris moving group. We resolve the disk in thermal-emission for the first time and find that the northern arm of the disk is 1.4$\\times$ brighter than the southern arm. In addition, we detect a broad peak in the combined \\emph{Spitzer} IRS and MIPS spectra at 60 -75 $\\mu$m that may be produced by emission from crystalline water ice. We model the IRS and MIPS data using a size distribution of amorphous olivine and water ice grains ($dn/da$ $\\propto$ $a^{-2.25}$ with $a_{min}$ consistent with the minimum blow out size and $a_{max}$ = 20 $\\mu$m) located at a distance of 86.3 AU from the central star, as observed in previously published scattered-light images. Since the photo-desorption lifetime for the icy particles is $\\sim$1400 yr, significantly less than the estimated $\\sim$12 Myr age of the system, we hypothesize that we have detected debris  that may be steadily replenished by collisions among icy Kuiper belt object-like parent bodies in a newly forming planetary system. ",
        "introduction": "Debris disks are dusty, gas-poor disks around main sequence stars. The estimated lifetimes for their constituent grains are typically shorter than the ages of their central stars by orders of magnitude, suggesting that the observed dust is debris, replenished from a reservoir such as by collisions among parent bodies or sublimation of comet-like bodies. To date, \\emph{Infrared Astronomical Satellite (IRAS)} and \\emph{Spitzer} studies have identified hundreds of debris disk candidates via thermal-emission from circumstellar dust (e.g. Hillenbrand et al. 2008; Rhee et al. 2007). Approximately 95\\% of these candidates possess dust with cool color temperatures ($T_{gr}$ $<$ 100 K), similar to that expected for dust in the solar system's Kuiper Belt (Meyer et al. 2007).  Broad-band near-infrared photometry and low resolution spectroscopy suggest that KBOs possess ices on their surfaces (Schaller \\& Brown 2007; Delsanti et al. 2006). The composition of dust in debris disks may be similar but has not yet been well-quantified (Chen et al. 2006).  High resolution scattered-light and thermal-emission imaging and spectroscopy may reveal disk structure and provide insight into grain properties via (1) brightness peaks in scattered-light or thermal-emission images that indicate  local dust density enhancements and/or dust temperature asymmetries and (2) spectral features at mid- and far-infrared wavelengths. Only a handful of debris disks have been imaged both in thermal-emission and scattered-light thus far: $\\beta$ Pictoris, HD 32297, HR 4796A, HD 107146, and Fomalhaut.  There may be tantalizing evidence to suggest that ices are present in these debris disks. Highly porous grains, similar to those used to model the infrared spectra of comets, have been used to model the $\\beta$ Pictoris infrared spectrum and spectral energy distribution (Chen et al. 2007; Li \\& Greenberg 1998). Excess thermal-emission in the inner region of the HD 32297 disk may indicate the presence of ice at temperatures near sublimation (Fitzgerald et al. 2007). However, moderate resolution, far-infrared spectra are needed to detect unambiguously emission from icy grains that is expected to be abundant in these systems.  \\targst\\ is an F5/F6V star with a \\emph{Hipparcos}-determined distance of 50.6 pc. It has been identified as a member of the $\\beta$ Pictoris moving group, with an estimated age of $\\sim$12 Myr (Zuckerman et al. 2001). Mannings \\& Barlow (1998) first identified HD 181327 as a possible debris disk based on the detection of \\emph{IRAS} 25 and 60$\\mu$m excesses. \\emph{Hubble Space Telescope (HST)} NICMOS and ACS coronagraphic imaging has resolved an inclined ring of dust in scattered-light at a distance of 86.3 AU with a width of 36 AU and an inclination, $i$ = 31.7$\\arcdeg$ from face-on. The scattered light is $\\sim$0.17\\% of the incident stellar light and possesses a brightness asymmetry that is well fit using an asymmetric scattering parameter, $g$ = 0.30$\\pm$0.03, at 1.1 $\\mu$m, suggesting that the grains producing the scattered-light are small, with radii smaller than 1 $\\mu$m (Schenider et al. 2006) if the grains are uniformly distributed in the disk. \\emph{Spitzer} IRS observations measure a steeply rising 5-35 $\\mu$m spectrum that is well characterized by an 80 K single-temperature black body and a fractional infrared luminosity, $L_{IR}/L_{*}$ = 3.1$\\times$10$^{-3}$ (Chen et al. 2006), commensurate with that measured toward $\\beta$ Pictoris. We have carried out a high resolution multi-wavelength imaging and spectroscopy study to infer the grain properties around \\targst. ",
        "conclusions": "We have obtained a T-ReCS \\bQa-band (18.3 $\\mu$m) image and a  55 - 90 $\\mu$m MIPS SED-mode spectrum of \\targst, an F5/F6V star in the $\\beta$ Pic moving group with a resolved scattered-light disk. Based on our analysis of the multi-wavelength imaging and spectroscopy, we conclude the following:  1. The northern arm of the \\targst\\ disk is 1.4 times brighter than the southern arm in thermal-emission, suggesting that either the density and/or the temperature of the grains in this region is higher than in the southern arm of the disk.  2. The overall properties of the unresolved spectral energy distribution (at 5.5 - 90 $\\mu$m), the resolved thermal-emission at 18.3 $\\mu$m, and the resolved scattered-light at 1.1 $\\mu$m can be reproduced by a population of 1 - 20 $\\mu$m amorphous olivine and crystalline water ice grains located at a distance of 86.3 AU from the central star.  3. Since the estimated photo-desorption lifetime of 1.5 $\\mu$m water ice grains is 1400 yr, significantly shorter than the age of \\targst, the grains must be replenished from a reservoir such as collisions among parent bodies, perhaps Kuiper Belt objects."
    },
    "0808/0808.1471_arXiv.txt": {
        "abstract": "{}{We study the X-ray morphology and dynamics of the galaxy cluster Abell 514. Also, the relation between the X-ray properties and Faraday Rotation measures of this cluster are investigated in order to study the connection of magnetic fields and the intra-cluster medium.}{We use two combined \\textit{XMM--Newton} pointings that are split into three distinct observations.}{The data allow {us} to evaluate the overall cluster properties like temperature and metallicity with high accuracy. The cluster has a temperature of 3.8$\\pm$0.2 keV and a metallicity of 0.22 $\\pm$0.07 in solar units. Additionally, a temperature map and the metallicity distribution are computed, which are used to study the dynamical state of the cluster in detail. Abell 514 represents an interesting merger cluster with many substructures visible in the X-ray image and in the temperature and abundance distributions. These results are used to investigate the connection between the ICM properties and the magnetic field of the cluster by comparing results from radio measurements. The new \\textit{XMM--Newton} data of Abell 514 confirm the relation between the X-ray brightness and the sigma of the Rotation Measure ($S_{\\rm X}$ - $\\sigma_{\\rm RM}$ relation) proposed by Dolag et al. (2001).} {} ",
        "introduction": "It is now well accepted that the intra-cluster medium (ICM)in clusters of galaxies is magnetized. The magnetic fields can be traced by diffuse cluster wide synchrotron radio emission (Giovannini et al. 1991, 1993, Feretti 1999 and Feretti \\& Giovannini 2007) or Inverse Compton hard X-ray radiation caused by relativistic electrons. Additionally, an indirect measure of the strength of magnetic fields is the rotation measure (RM), in which radiation from background radio sources is studied: according to the strength of the magnetic field inside the cluster, the polarization angle of the radio emission is rotated. The different observations lead to the conclusion that magnetic fields in clusters of galaxies have strengths of a few $\\mu$G (Carilli \\& Taylor 2002).\\\\ Dolag et al. (2001) showed that a relation exists between the X-ray {\\bf surface brightness} and the root mean square scatter ($\\sigma_{\\rm RM}$) of the Faraday Rotation Measures ($S_{\\rm X}$ - $\\sigma_{\\rm RM}$ relation) that are used to evaluate the strength of the magnetic field. This relation is an important tool to study the connection between the magnetic field and the intra-cluster gas density and temperature (Dolag et al. 2001). In particular clusters with polarized extended radio sources are of interest, because it is possible to evaluate the RM scatter well. More sources in one cluster give the possibility to get values for the magnetic field strength in different parts of the cluster, and are therefore very important observational objects to understand the relation between the magnetic field and the X-ray properties. In order to compare the magnetic field and other cluster properties at the position of each radio source an X-ray image is required. The surface brightness $S_{\\rm X}$ and the RMS can be determined at the position of each radio source.\\\\ Since Abell 514 has several radio sources that offer the possibility to study the $S_{\\rm X}$ -  $\\sigma_{\\rm RM}$ relation, it was chosen for our study. In this paper, we present results from three \\textit{XMM--Newton} observations of this cluster. \\\\ Throughout the paper, a $\\Lambda$CDM ($\\Omega_{\\Lambda}$ = 0.7 and $\\Omega_{\\rm m}$ = 0.3) cosmology with a Hubble constant of 70 km s$^{-1}$ Mpc$^{-1}$  was assumed.  \\subsection{Connection of the magnetic field and the ICM density}  The two observables $S_{\\rm X}$ and the RMS scatter ($\\sigma_{\\rm RM}$) compare the two line of sight integrals: \\begin{equation} \\label{integrals} S_{\\rm X} \\propto \\int{n^{2}_{\\rm e}\\sqrt{T}dx} \\leftrightarrow \\sigma_{\\rm RM} \\propto \\int{n_{\\rm e} B_{\\|}dx} \\end{equation} where n$_{\\rm e}$ is the electron density and $B_{\\|}$ the magnetic field component parallel to the line of sight. (Dolag et al. 2001; Clarke et al. 2001)   \\begin{figure}[h] {\\includegraphics[width=\\columnwidth]{rmvsflux.eps}} \\caption{The scatter of the root mean square of the Faraday Rotation measure (RMS) against the X-ray flux of a sample of clusters, for which both measurements are available.} \\label{rmvsflux} \\end{figure}  When $\\sigma_{\\rm RM}$ is plotted versus the X-ray flux a clear relation can be seen. This relation can be fitted by: \\begin{equation} \\label{sigma_rm_relation} \\sigma_{\\rm RM} = A\\Big(\\frac{S_{\\rm X}}{10^{-5} \\mbox{erg}/\\mbox{cm}^2/\\mbox{s}} \\Big)^{\\alpha} \\end{equation}  A simple interpretation of this relation (e.g. assuming the temperature within the ICM and the scale-length of the magnetic field to be fixed) is that the slope $\\alpha$ reflects the scaling of the magnetic field strength ($B$) with the electron density ($n_e$). An exact relation between these two scalings, $B$-$n_{\\rm e}$ and $\\sigma_{\\rm RM}$-$S_{\\rm X}$, is derived in Dolag et al. (2001) assuming a simplified model for galaxy clusters. Note that the uncertainties in the 3D position of the individual sources (which are not known) lead to significant uncertainties in the derived $\\sigma_{\\rm RM}$ and therefore imprints a substantial scatter in the scaling relation. In fact this is the largest contribution to the the error bars we calculate for $\\sigma_{\\rm RM}$ (see Dolag et al. 2001 for details).  Additionally, it seems that there is a suspected dependence on the cluster temperature: clusters with a high overall temperature also seem to have high $\\sigma_{\\rm RM}$ values (see Fig. \\ref{rmvsflux}). To study such matter in detail, clusters that contain radio sources have to be investigated very accurately in radio and X-rays. ",
        "conclusions": "\\label{discussion}  \\subsection{Candidate for a cold front or a shock?}  A prominent morphological structure of Abell 514 is a steep decline in X-ray surface brightness towards the Northwest region. This can be seen as a sharp edge in the image (see Fig. \\ref{smoIma}), as well as a quick drop of the surface brightness profile outside $\\sim$ 2 arcmin (see Fig. \\ref{NW_peak}).  Possible explanations for such a feature can be either a cold front or a shock caused by the merger process. Similar features were found by Markevitch et al. (2000) and Vikhlinin et al. (2002) in the clusters Abell 2142 and Abell 3667. Another example for a similar structure was also found in Abell 2256 by Sun et al. (2002). During a cluster merger, a cool core of a subpart of a cluster can survive the merging process. This is characterised by the fact that the temperature inside a brightness edge is lower than in the surrounding region. The other explanation for a feature like the one seen in Abell 514 would be a shock where the material is compressed.\\\\ To test if the edge in Abell 514 is caused by a cold front or a shock we study two regions, one inside and one outside the edge visible in the smoothed X-ray image (Fig. \\ref{smoIma}), with respect of their density and temperature. The regions used for this analysis are shown in Fig. \\ref{where_depr}.  \\begin{figure}[h] {\\includegraphics[width=\\columnwidth]{where_depr.eps}} \\caption{Regions considered for the deprojection analysis in order to investigate the nature of the drop in surface brightness.} \\label{where_depr} \\end{figure}  Region 1 is the region inside the \"edge\", while region 2 is the area in the outer part. We apply the deprojection method by using the Xspec model {\\ttfamily projct} to calculate densities inside and outside of this border. Also the temperatures in both regions were calculated and compared with each other. The results are shown in Table \\ref{ColdFront_Shock}.  \\begin{table}[h] \\caption{The density and temperature inside (Region 1) and outside (Region 2) the brightness edge.} \\centering \\label{ColdFront_Shock} \\begin{tabular}{|c|c|c|} \\hline & Region 1 & Region 2 \\\\ \\hline Density [10$^{-3}$cm$^{-3}$]& 0.91 $\\pm$ 0.11 & 0.51 $\\pm$ 0.06 \\\\\\hline Temperature [keV]& 4.5 $\\pm$ 0.8 & 3.6 $\\pm$ 0.5 \\\\\\hline \\end{tabular} \\end{table}  The temperature inside the border is slightly higher than outside, but no jump in temperature can be deduced from our data, especially not a jump from a cool core to a warmer surrounding. Inside the errorbars, both temperatures can be seen as the same. Therefore the discontinuity in surface brightness cannot be caused by a cold front. The density however shows a clear discontinuity. It is therefore possible that the brightness jump is due to a shock. Such a shock can be the result from an earlier merger, with the different structures not distinguishable by eye any more. \\\\ The visible interaction between the SE peak and the NW one is most likely not responsible for this feature. We see that the metallicities between the two peaks are very different. It is therefore plausible that they have not merged yet and cannot cause the feature in the surface brightness seen in the NW peak. \\\\ Another possibility could be an interaction of the main X-ray peak with the small blob from the south east part. But since this structure is still 500 kpc away from the main cluster and no connection between the two parts can be seen we do not expect to see any interaction effects yet between those parts.\\\\   \\subsection{The $S_{\\rm X}$ -  $\\sigma_{\\rm RM}$ relation and the magnetic field}  According to theory (Tribble 1993, Dolag et al. 1999), the magnetic field is amplified in a hot merger cluster. The $S_{\\rm X}$ - $\\sigma_{\\rm RM}$ relation is clearly dependent on the temperature of the cluster (see Sect. 1.1). For Abell 514, this general trend can be studied. Although Abell 514 is a merger cluster, its magnetic field is still quite low. This can be seen in good agreement with the low overall temperature of the cluster. Still, compared to other cool clusters, Abell 514 shows a slightly higher $\\sigma_{RM}$ which is most likely due to the ongoing merger that already enhanced the magnetic field.\\\\ One main aim of the \\textit{XMM--Newton} observations was to get new values for the X-ray flux in the regions where the radio sources are. It has to be mentioned that the true location of these radio sources inside the cluster is not known. This fact is taken into account in the errors given for the $\\sigma_{\\rm RM}$ value. The error bars cover the range of values between a source located in the cluster center and one behind the cluster.  \\begin{figure}[h] {\\includegraphics[width=\\columnwidth]{RMS_XrayFlux_neu.eps}} \\caption{The X-ray flux -  $\\sigma_{\\rm RM}$ relation with the new data points for Abell 514. The new values are in good agreement with the general slope of the relation} \\label{RMS_neu}. \\end{figure}  The new values are then compared to the results from measurements of the magnetic field (via the $S_{\\rm X}$ - $\\sigma_{\\rm RM}$). They fit well with other measurements from Coma, A119 etc. (see Fig. \\ref{RMS_neu}). Fig. \\ref{RMS_neu} shows the results from the new measurements, with the data point for the other clusters (Coma, A119, etc.) being converted to the same energy band. In general, the data points obtained for A514 are in good agreement with the relation found from the rest of the clusters. The lines in fig. \\ref{RMS_neu} represent the correlations for the distinct clusters. We see that the data points of Abell 514 lie above the correlations of all the other clusters. This can be seen as an indication for an amplification of the magnetic field due to the ongoing merger in Abell 514. Overall, the points from Abell 514 make the whole correlation (if we use the observational data) less steep. Without the data points from Abell 514, the slope parameter is 1.19, while it is 0.98 with them. Inside an error of 10\\% both values agree. To avoid any instrumental bias in this study in the future, we plan to obtain {\\it XMM--Newton} data for the other clusters in this sample as well.\\\\  Additionally, with the creation of a temperature map, it is possible to compare the strength of the RMS scatter $\\sigma_{\\rm RM}$ from the rotation measures with the temperature of the ICM in the area of the radio source. Table \\ref{RMT2} shows the results.  \\begin{table} \\caption{Comparison of $\\sigma_{\\rm RM}$ with the temperature. The regions are as indicated in Fig. \\ref{RegionsTmap}} \\label{RMT2} \\centering \\begin{tabular} {|p{3cm}|c|p{2.3cm}|} \\hline Region Nr. & T (keV) & $\\sigma_{RM} (rad/m^{2})$ \\\\\\hline Region 2 (includes Radio source B2) & 3.2 $\\pm$ 0.2 & 63 $^{+16}_{-41}$\\\\\\hline Region 4 (includes Radio sources D$_{north}$ and D$_{south}$) & 4.9 $\\pm$ 0.4 & 54 $^{+12}_{-21}$ (north) 38 $^{+10}_{-23}$(south) \\\\\\hline \\end{tabular} \\end{table}   Here, $\\sigma_{\\rm RM}$ is lower in the hottest region and higher in the cool, X-ray brightest part, that seems to be the most relaxed part of the cluster. However, this is not in contradiction with the above relation. Inside the cluster, more complicated effects take place additionally to the overall properties, that cannot yet be resolved with the current observations. Also, the RMS measurements are taken from a smaller area than the spectra we use to deduce the temperature. Small scale fluctuations inside these regions are therefore possible and not taken into account in table \\ref{RMT2}."
    },
    "0808/0808.1647_arXiv.txt": {
        "abstract": "The formation of the first galaxies at redshifts $z\\sim 10-15$ signaled the transition from the simple initial state of the universe to one of ever increasing complexity. We here review recent progress in understanding their assembly process with numerical simulations, starting with cosmological initial conditions and modelling the detailed physics of star formation. In this context we emphasize the importance and influence of selecting appropriate initial conditions for the star formation process. We revisit the notion of a critical metallicity resulting in the transition from primordial to present-day initial mass functions and highlight its dependence on additional cooling mechanisms and the exact initial conditions. We also review recent work on the ability of dust cooling to provide the transition to present-day low-mass star formation. In particular, we highlight the extreme conditions under which this transition mechanism occurs, with violent fragmentation in dense gas resulting in tightly packed clusters. ",
        "introduction": "One of the key goals in modern cosmology is to study the assembly process of the first galaxies, and understand how the first stars and stellar systems formed at the end of the cosmic dark ages, a few hundred million years after the Big Bang. With the formation of the first stars, the so-called Population~III (Pop~III), the universe was rapidly transformed into an increasingly complex, hierarchical system, due to the energy and heavy elements they released into the intergalactic medium (IGM; for recent reviews, see Barkana \\& Loeb 2001; Miralda-Escud{\\'e}; Bromm \\& Larson 2004; Ciardi \\& Ferrara 2005; Glover 2005). Currently, we can directly probe the state of the universe roughly a million years after the Big Bang by detecting the anisotropies in the cosmic microwave background (CMB), thus providing us with the initial conditions for subsequent structure formation. Complementary to the CMB observations, we can probe cosmic history all the way from the present-day universe to roughly a billion years after the Big Bang, using the best available ground- and space-based telescopes. In between lies the remaining frontier, and the first galaxies are the sign-posts of this early, formative epoch.  There are a number of reasons why addressing the formation of the first galaxies and understanding second-generation star formation is important. First, a rigorous connection between well-established structure formation models at high redshift and the properties of present-day galaxies is still missing. An understanding of how the first galaxies formed could be a crucial step towards undertanding the formation of more massive systems. Second, the initial burst of Pop~III star formation may have been rather brief due to the strong negative feedback effects that likely acted to self-limit this formation mode (Madau et al. 2001; Ricotti \\& Ostriker 2004; Yoshida et al. 2004; Greif \\& Bromm 2006). Second-generation star formation, therefore, might well have been cosmologically dominant compared to Pop~III stars. Despite their importance for cosmic evolution, e.g., by possibly constituting the majority of sources for the initial stages of reionization at $z>10$, we currently do not know the properties, and most importantly the typical mass scale, of the second-generation stars that formed in the wake of the very first stars. Finally, a subset of second-generation stars, those with masses below $\\simeq 1~M_{\\odot}$, would have survived to the present day. Surveys of extremely metal-poor Galactic halo stars therefore provide an indirect window into the Pop~III era by scrutinizing their chemical abundance patterns, which reflect the enrichment from a single, or at most a small multiple of, Pop~III supernovae (SNe; Christlieb et al. 2002; Beers \\& Christlieb 2005; Frebel et al. 2005). Stellar archaeology thus provides unique empirical constraints for numerical simulations, from which one can derive theoretical abundance patterns to be compared with the data.  Focusing on numerical simulations as the key driver of structure formation theory, the best strategy is to start with cosmological initial conditions, follow the evolution up to the formation of a small number ($N<10$) of Pop~III stars, and trace the ensuing expansion of SN blast waves together with the dispersal and mixing of the first heavy elements, towards the formation of second-generation stars out of enriched material (Greif et al. 2007; Wise \\& Abel 2007a). In this sense some of the most pressing questions are: How does radiative and mechanical feedback by the very first stars in minihalos affect the formation of the first galaxies? How and when does metal enrichment govern the transition to low-mass star formation? Is there a critical metallicity at which this transition occurs? How does turbulence affect the chemical mixing and fragmentation of the gas? These questions have been addressed with detailed numerical simulations as well as analytic arguments over the last few years, and we here review some of the most recent work. For consistency, all quoted distances are physical, unless noted otherwise. ",
        "conclusions": "Understanding the formation of the first galaxies marks the frontier of high-redshift structure formation. It is crucial to predict their properties in order to develop the optimal search and survey strategies for the {\\it JWST}. Whereas {\\it ab-initio} simulations of the very first stars can be carried out from first principles, and with virtually no free parameters, one faces a much more daunting challenge with the first galaxies. Now, the previous history of star formation has to be considered, leading to enhanced complexity in the assembly of the first galaxies. One by one, all the complex astrophysical processes that play a role in more recent galaxy formation appear back on the scene. Among them are external radiation fields, comprising UV and X-ray photons, as well as local radiative feedback that may alter the star formation process on small scales. Perhaps the most important issue, though, is metal enrichment in the wake of the first SN explosions, which fundamentally alters the cooling and fragmentation properties of the gas. Together with the onset of turbulence (Wise \\& Abel 2007b; Greif et al. 2008b), chemical mixing might be highly efficient and could lead to the formation of the first low-mass stars and stellar clusters (Clark et al. 2008).  In this sense a crucial question is whether the transition from Pop~III to Pop~II stars is governed by atomic fine-structure or dust cooling. Theoretical work has indicated that molecular hydrogen dominates over metal-line cooling at low densities (Jappsen et al. 2007a,b), and that fragment masses below $\\sim 1~\\rm{M}_{\\odot}$ can only be attained via dust cooling at high densities (Omukai et al. 2005; Clark et al. 2008). On the other hand, observational studies seem to be in favor of the fine-structure based model (Frebel et al. 2007), even though existing samples of extremely metal-poor stars in the Milky Way are statistically questionable (Christlieb et al. 2002; Beers \\& Christlieb 2005; Frebel et al. 2005). Moreover, these studies assume that their abundances are related to primordial star formation -- a connection that is still debated (Lucatello et al. 2005; Ryan et al. 2005; Komiya et al. 2007). In light of these uncertainties, it is essential to push numerical simulation to ever lower redshifts and include additional physics in the form of radiative feedback, metal dispersal (Greif et al. 2008a), chemisty and cooling (e.g. Glover \\& Jappsen 2007) and the effects of magnetic fields (e.g. Xu et al. 2008). We are confident that a great deal of interesting discoveries, both theoretical and observational, await us in the rapidly growing field of early galaxy formation."
    },
    "0808/0808.2534_arXiv.txt": {
        "abstract": "This paper concludes a systematic search for evidence of the Monoceros Ring and Canis Major dwarf galaxy around the Galactic Plane. Presented here are the results for the Galactic longitude range of {\\it l} = (280 - 025)$\\deg$. Testing the claim that the Monoceros Ring encircles the entire Galaxy, this survey attempts to document the position of the Monoceros Ring with increasing Galactic longitude. Additionally, with the discovery of the purported Canis Major dwarf galaxy, searching for more evidence of its interaction is imperative to tracing its path through the Galaxy and understanding its role in the evolution of the Milky Way. Two new detections of the Monoceros Ring have been found at ({\\it l, b}) = (280,$+$15)$\\deg$ and (300,$+$10)$\\deg$. Interestingly, in general there seem to be more detections above the Plane than below it; in this survey around $\\frac{2}{3}$ of the firm Monoceros Ring detections are in the North. This coincides with the Northern detections appearing to be qualitatively denser and broader than their Southern counterparts. The maximum of the Galactic Warp in the South is also probed in this survey. It is found that these fields do not resemble those in the Canis Major region suggesting that the Warp does not change the shape of the CMD as is witnessed around Canis Major. The origins and morphology of the Monoceros Ring is still elusive primarily due to its enormous extent on the sky. Continued probing of the Galactic Outer Disc is needed before a consensus can be reached on its nature. ",
        "introduction": "The Monoceros Ring (MRi), discovered in 2002 \\citep{2002ApJ...569..245N} has now been traced around the Galaxy from {\\it l} = 75 - 260$\\deg$, as shown through the Sloan Digital Sky Survey \\citep{2002ApJ...569..245N}, Two-micron All Sky Survey \\citep{2003ApJ...594L.115R}, Isaac Newton Telescope Wide Field Camera Survey \\citep{2003MNRAS.340L..21I} and the Anglo-Australian Telescope Wide Field Imager Survey \\citep{2005MNRAS.362..475C}. Continuing around the Galactic plane, this survey extends these previous results to complete the first Wide Field Imager survey of the MRi around the Galaxy that began with the INT/WFC in \\citet{2005MNRAS.362..475C}. Studies into this structure have been discussed in Paper I of this series, \\citep[][]{2007MNRAS.376..939C}, and references therein. Additional to this, an RR-Lyrae survey of the Galactic Halo using QUEST data has also revealed the presence of the MRi and investigated the overdensity in Canis Major \\citep[][]{2006AJ....132..714V,2007IAUS..241..359M}.  Residing in the Thick Disc of the Milky Way (MW), the MRi is revealed only by obtaining deep photometry of large patches of sky, typically greater than 1 square degree. In this preliminary first pass of the MW, the Thick Disc was sampled at Galactic latitudes nominally between {\\it b} = $\\pm$10$\\deg$-20$\\deg$ and about every 20 degrees in Galactic longitude. To date, the entire survey has strong detections of the MRi in 14 regions with 3 additional tentative detections out of 25 regions observed. It has been found on both sides of the Galactic plane at Galactic latitudes from 4$\\deg$ - 20$\\deg$ and its extent away from the plane is as yet undetermined although SDSS results suggest that it could be as high as +30$\\deg$ \\citep{2007ApJ...658..337B}. Numerical simulations of the MRi as a tidal stream predict it to have multiple wraps around MW, although the current dataset cannot distinguish between different aspects of the stream nor whether the different detections are part of a coherent structure. Figure~\\ref{figmonster} shows the previous detections of the MRi as reported in \\citet{2005MNRAS.362..475C} and \\citet{2007MNRAS.376..939C}. Figure 2 of Paper I shows how these CMDs can be interpreted by showing the approximate location of the Thin, Thick and Halo stars in the field. Figure 5 of Paper I illustrates which components are being referenced when discussed in the text. Each of the fields in this figure are pixelated Colour-Magnitude diagrams. The pixel values represent the square-root of the number of stars in that pixel. Ordered by increasing Galactic Longitude, Figure~\\ref{figmonster} tentatively shows the changing strength and thickness of the MRi around the Plane. Qualitatively, the strength of the MRi can be seen in comparison with the MW components. Additionally, the only apparent difference between Northern and Southern detections is perhaps that the Southern MRi features appear qualitatively narrower than their Northern counterparts. There is no clear explanation as to why this is the case.  While there is more and more evidence regarding the true extent of this structure, there is very little information concerning many of its generic properties. As such, no direct measurement of the density profile of any part of the stream around the sky has been made, nor a complete understanding of its true extent on the sky. In the region covered by the SDSS, \\citet{2008ApJ...673..864J} and \\citet{2008arXiv0804.3850I}, report on the presence of the MRi with regard to its number density, metallicity and kinematics. A clear overdensity of stars can be seen at a distance of 16 kpc \\citep{2008ApJ...673..864J} and \\citet{2008arXiv0804.3850I} reports a mean metallicity of [Fe/H] = -0.95 with a scatter of around 0.15 dex. Kinematically, they show a spread of velocities rotating consistently faster than the Local Standard of Rest and in accordance with the predictions of ~\\citet{2005ApJ...626..128P}.  The only possible candidate for the stream's progenitor is an overdensity of stars found in Canis Major but possible confusion with the Galactic Warp has created doubts on this detection. Numerical simulations created using the properties of these stars have predicted the location and extent of the MRi with good accuracy and so adds support to the dwarf galaxy scenario. This on-going debate centres on whether observations of the Canis Major overdensity conform to known Galactic structure, such as the Warp, or can be considered truly additional. Paper I of this series outlines some of the possible inconsistencies between predicted properties of the Galactic Warp and direct observations of these structures. In response to this \\citet{2007arXiv0707.4440L} have presented an explanation relying on only first order Galactic structure. Further refinement of the properties of the MRi are needed to determine whether CMa is the most likely candidate as its progenitor.   \\begin{figure} \\centerline{ \\psfig{figure=Fig01.ps}} \\caption[]{Visual summary of all the previous Monoceros Ring detections from the INT/WFC survey \\citep{2005MNRAS.362..475C} and AAT/WFI survey (Paper I) of the outer Disc. The Colour-Magnitude diagrams are of two types {\\it V,i} and {\\it g,r}. For more detailed analysis of these fields and the reported detections see the relevant articles. \\label{figmonster}} \\end{figure} ",
        "conclusions": "This paper reports on 2 new detections and 7 non-detections of the Monoceros Ring tidal stream. The results presented here conclude a survey tracing this feature around the entire Galactic plane. The previously reported detections of the survey are presented in Figure~\\ref{figmonster}. Comparing the relative strengths of the MRi and the main MW population it appears qualitatively that the stream is denser and broader above the Plane than below but as such there is no explanation why this would be the case. The part of the overall survey presented here shows no evidence of the strong Canis Major dwarf main sequence in the CMDs. The CMa sequence is historically identified as a shift in the position and shape of the strongest main sequence in the CMD. For the fields presented here, the dominant main sequence feature in the CMDs is easily attributed to the Thin and Thick Discs. In each instance where the distance has been determined to these structures it is in accordance with the Besan\\c{c}on synthetic galaxy model predictions. Therefore they can be confidentially associated with the MW. The only field which might be expected, from the \\citet{2005MNRAS.362..906M} model, to contain the CMa signature is ({\\it l,b}) = (280,-15)$\\deg$. This field does not show this CMa-style sequence in the CMD (Figure~\\ref{fig280m15}).  Comparing these new MRi detections with the two current numerical simulations of the stream and putative dwarf galaxy progenitor, has led to inconclusive results. The \\citet{2005MNRAS.362..906M} model north of the Galactic Plane roughly traces the locations of the detections. In the South the correspondence between the model and the detections is adequate with some noted exceptions. Several detections presented in this survey indicate, with reasonable certainty, the locations in which the model is incorrect. For the \\citet{2005ApJ...626..128P} model, there is less correlation between the data points and the predicted stream locations than is seen with the \\citet{2005MNRAS.362..906M} model. Although some points do seem to represent a better fit it is important to note though that a significant proportion of the \\citet{2005ApJ...626..128P} model does reside outside of a Galactic latitude of {\\it b} = $\\pm$20$\\deg$. So much of the model has not been sampled by this survey. Indeed, it is easy to see that this survey is too narrow in Galactic latitude in comparison with the data used to construct the model and the predictions it makes. Drawing a conclusion based on these results in inadvisable but there is little here to strongly support this model. Both models obviously will require reworking to include the new information available along with more observations to test their predictions.  With regard to the Besan\\c{c}on synthetic galaxy model, there is no presence of the MRi as part of natural Galactic structure. In almost all fields in this survey, the bulk Milky Way components of Thin, Thick Disc have been accurately modelled. There is no systematic discrepancy between the model and data even in regions containing the Galactic Warp. Only the regions around Canis Major, as discussed in Paper I, show a definite shift from the observational data. Given the data supports the predictions of the Besan\\c{c}on model in all but the MRi detections, it is reasonable to assume this structure is indeed additional to the usual Galactic components.  Determining the density profile of this feature around the Galaxy and indeed connecting detections is an important next step in resolving its origins. To date, targeted deep surveys, such as this, have resolved many important questions surrounding this structure. This survey sheds some light on the impact of the Galactic Warp on the Colour-Magnitude Diagrams showing it does not effect its morphology significantly and that the Besan\\c{c}on model is adequate for most fields. This has implications with regard how the fields in the Canis Major region are to be interpreted as the fields there have obviously different characteristics. While the nature of the Monoceros Ring still remains quite elusive, this is primarily due to its large extent on the sky and its location close to the Plane. For the time being, both the Galactic origin scenario and the tidal stream hypothesis are still possibilities for this structure. The completed survey, presented here, has shown that a targeted campaign of observations can provide insights on not only this structure but also generic Galactic structures as well."
    },
    "0808/0808.0641_arXiv.txt": {
        "abstract": "We report astrometric observations of H$_2$O masers around the red supergiant VY Canis Majoris (VY CMa) carried out with VLBI Exploration of Radio Astrometry (VERA). Based on astrometric monitoring  for 13 months, we successfully measured a trigonometric parallax of 0.88 $\\pm$ 0.08 mas, corresponding to a distance of 1.14 $^{+0.11} _{-0.09}$ kpc. This is the most accurate distance to VY CMa and the first one based on an annual parallax measurement. The luminosity of VY CMa has been overestimated due to a previously accepted distance. With our result, we re-estimate the luminosity of VY CMa to be (3 $\\pm$ 0.5) $\\times$ 10$^5$ L$_{\\odot}$ using the bolometric flux integrated over optical and IR wavelengths. This improved luminosity value makes location of VY CMa on the Hertzsprung-Russel (HR) diagram much closer to the theoretically allowable zone (i.e. the left side of the Hayashi track) than previous ones, though uncertainty in the effective temperature of the stellar surface still does not permit us to make a final conclusion. ",
        "introduction": "Massive stars play an important role in the evolution of the Universe particularly in the late stages of their evolution. Through their strong stellar winds and supernova explosions, they inject mechanical energy into the interstellar medium (ISM) \\citep{bib:Abbott1982}. Also, they are principal sources of heavy elements in the ISM. In spite of their importance, our understanding of late stages of massive stars' evolution is still poor. One of the reasons for this is that massive stars are extremely rare partly because of their short lifetime. Due to their small numbers, the properties of evolved massive stars are still uncertain. For instance, as discussed in \\citet{bib:Massey2003}, \\citet{bib:MasseyOlsen2003} and \\citet{bib:Levesque2005}, there was a discrepancy between observed and theoretically predicted locations of red supergiants, high-mass evolved stars, on the Hertzsprung-Russel (HR) diagram. Compared with stellar evolutionary models, observed red supergiants appear to be too cool and too luminous.  VY Canis Majoris (VY CMa) is one of the most well-studied red supergiants. Similar to other red supergiants, the location of VY CMa on the HR diagram is also uncertain. For instance, \\citet{bib:Monnier1999}, \\citet{bib:Smith2001} and \\citet{bib:Humphreys2007} obtained the luminosity of (2--5) $\\times$ 10$^5$ L$_{\\odot}$ from the spectral energy distribution (SED) assuming the distance of 1.5 kpc \\citep{bib:Lada1978} and effective temperature of 2800--3000 K based on the stellar spectral type \\citep{bib:LeSidaner1996}. However, \\citet{bib:Massey2006} suggested that the above parameters would make VY CMa cooler and more luminous than what current evolutionary models allow and would place it in the ``forbidden zone\" of the HR diagram, which is on the right-hand side of the Hayashi track in the HR diagram (see figure \\ref{fig:HRDthis}). \\citet{bib:Massey2006} reexamined the effective temperature and obtained a new value of 3650 $\\pm$ 25 K based on optical spectrophotometry combined with a stellar atmosphere model and suggested that the luminosity of VY CMa is only 6.0 $\\times$ 10$^4$ L$_{\\odot}$, which is probably a lower limit \\citep{bib:Massey2008}.  \\citet{bib:Levesque2005} determined the effective temperatures of 74 Galactic red supergiants based on optical spectrophotometry and stellar atmosphere models. They obtained the effective temperatures of the red supergiants with a precision of 50 K. Their new effective temperatures are warmer than those in the literature. These new effective temperature values seem to be consistent with the theoretical stellar evolutionary tracks. However, the luminosities of the red supergiants still had large uncertainty due to possible errors in estimated distances. Since most of red supergiants are very far, it has been difficult to apply the most reliable trigonometric parallax method to their distance measurements.  In the case of VY CMa, the currently accepted distance of 1.5 kpc \\citep{bib:Lada1978} is obtained by assuming that VY CMa is a member of the NGC 2362 star cluster and that the distance of NGC 2362 is same as that of VY CMa. The distance to the NGC 2362 was determined from the color-magnitude diagram \\citep{bib:Johnson1961} with an accuracy of no better than 30 \\%. Since the luminosity depends on the square of the distance, the accuracy of luminosity estimation is worse than 50 \\%. The proper motions measured with H$_2$O masers in previous studies (\\cite{bib:Richards1998}; \\cite{bib:Marvel1998}) have not been contradictory to the above distance value (1.5 kpc). When the H$_2$O maser region around a star is modeled with a symmetrically expanding spherical shell, the distance to the star from the Sun can be inferred by assuming that the observed mean proper motion multiplied by the distance should be equal to the observed mean radial velocity of the H$_2$O masers. The proper motion velocities in \\citet{bib:Richards1998} are well consistent with the H$_2$O maser spectral line velocities at the distance of 1.5 kpc. Also, \\citet{bib:Marvel1998} derived the distance of VY CMa to be 1.4 $\\pm$ 0.2 kpc. However, such a ``statistical parallax'' method depends on kinematical models fitted to the observed relative proper motions. In fact, if we add rotation or anisotropic expansion to the simple spherical expansion model, the inferred distance value could vary beyond quoted error ranges. An accurate distance determination without an assumption is crucial for obtaining the right location of VY CMa on the HR diagram and for determining fundamental parameters of the star.  Thanks to the recent progress in the VLBI technique, the distance measurements with trigonometric parallaxes have become possible even beyond 5 kpc (e.g., \\cite{bib:Honma2007}). Since VY CMa has strong maser emission in its circumstellar envelope, we have conducted an astrometric observations of H$_2$O masers around VY CMa to measure an accurate parallax with VLBI Exploration of Radio Astrometry (VERA) and here we present the results. ",
        "conclusions": "We have observed the H$_2$O masers around the red supergiant VY CMa with VERA during 10 epochs spread over 13 months. Simultaneous observations for both H$_2$O masers around VY CMa and the position reference source J0725--2640 were carried out. We measured a trigonometric parallax of 0.88 $\\pm$ 0.08 mas, corresponding to a distance of 1.14 $^{+0.11} _{-0.09}$ kpc from the Sun. It is the first result that the distance of VY CMa is determined with an annual parallax measurement. There had been overestimation of the luminosities in previous studies due to the previously accepted distance. Using the most accurate distance based on the trigonometric parallax measurements and the bolometric flux from the observed SED, we estimated the luminosity of VY CMa to be (3 $\\pm$ 0.5) $\\times$ 10$^5$ L$_{\\odot}$. The accurate distance measurements provided the improved luminosity. The location of VY CMa on the HR diagram became much close to the theoretically allowable region, though there is still uncertainty in the effective temperature.  ~\\\\ We are grateful to the referee, Dr. Philip Massey, for helpful comments on the manuscript. We would like to thank Prof. Dr. Karl M. Menten for his invaluable comments and for his help improving the manuscript. The authors also would like to thank all the supporting staffs at Mizusawa VERA observatory for their assistance in observations."
    },
    "0808/0808.2644_arXiv.txt": {
        "abstract": "The empirical binary properties of brown dwarfs (BDs) differ from those of normal stars suggesting BDs form a separate population. Recent work by Thies and Kroupa revealed a discontinuity of the initial mass function (IMF) in the very-low-mass star regime under the assumption of a low multiplicity of BDs of about 15 per cent. However, previous observations had suggested that the multiplicity of BDs may be significantly higher, up to 45 per cent. This contribution investigates the implication of a high BD multiplicity on the appearance of the IMF for the Orion Nebula Cluster, Taurus-Auriga, IC~348 and the Pleiades. We show that the discontinuity remains pronounced even if the observed MF appears to be continuous, even for a BD binary fraction as high as 60\\pct. We find no evidence for a variation of the BD IMF with star-forming conditions. The BD IMF has a power-law index $\\abd\\approx +0.3$ and about 2 BDs form per 10 low-mass stars assuming equal-mass pairing of BDs. ",
        "introduction": "\\label{sec:introd} The origin of brown dwarfs (BDs) remains the subject of intense discussions.  There are two broad ideas on their origin: 1.~the classical star-like formation scenario of BDs (e.g. \\citealt{AdFa96,PaNo04}), and 2.~BDs and some very-low-mass stars (VLMSs) form as a separate population (hereafter named BD-like besides the classical star-like population) with a different formation history than stars, e.g. as ejected stellar embryos \\citep{ReiCla01,KB03b} or as disrupted wide binaries \\citep{GoWi07,StHuWi07}.  Additionally, the formation of BDs from Jeans instabilities in high-density filaments near the centre of a massive star-forming cloud has been recently suggested implying such BDs to be preferentially located in clusters with strong gravitational potentials \\citep{BCB08}.  The star-like formation scenario fails to reproduce the observed different binary properties of BDs and stars. Especially the truncation of the semi-major axis distribution between 10 and 20~AU for BDs and the different mass-ratio distribution of BDs and stars \\citep{Bouyetal03,Burgetal03,Maetal03,Ketal03,Cloetal03}, as well as the BD desert \\citep{2003IAUS..211..279M,GreLin06}, are difficult to account for if BDs form indistinguishably to stars. This implies the need of treating BDs as a separate population to stars.  The assumption of two separate populations would then require two separate initial mass functions (IMFs) of the individual bodies of a star cluster. Although the observed mass function may appear approximately continuous from the lowest mass BDs up to the highest mass stars \\citep{UKIDSS2007}, a discontinuity in the IMF may be present but be masked by `hidden' (unresolved) binaries only emerging if the observed MF is corrected for unresolved multiplicity.  This issue has been discussed in greater detail in Thies \\& Kroupa (2007, hereafter TK07) for the case of a low multiplicity of 15\\pct\\ of the BD-like population.  However, a higher BD multiplicity between 20\\pct\\ and 45\\pct\\ had been reported by some authors, e.g. \\citet{JefMax05,BasRei2006}.  This raises the need, dealt with this contribution, for including higher multiplicities as well as for an analysis of the general effects of a high multiplicity on the IMF and on the BD-to-star ratio.  In Section \\ref{sec:bdpop} we review the evidence for a separate BD-like population.  Section \\ref{sec:method} briefly introduces the mathematical method of calculating the IMF including unresolved binaries. In Section \\ref{sec:results} the new results are presented and compared to those of TK07. The Summary follows in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} A discontinuity in the IMF near the hydrogen burning mass limit appears if the binary properties of BDs and VLMSs on the one hand, and of stars on the other, are taken into account carefully when inferring the true underlying single-object IMF. This implies that BDs and some VLMSs need to be viewed as arising from a somewhat different formation channel than the stellar formation channel, but this result has been obtained by TK07 under the assumption that BDs have a binary fraction of only~15\\pct. A higher binary fraction may close the gap between the stellar and the BD IMF. We refer to BDs and those VLMSs formed according to the putative BD channel as ``BD-like'' bodies, whereas stars and those BDs formed according to the stellar channel as star-like. The BD-like channel remains unknown in detail, but theoretical ideas have emerged (Sections~\\ref{sec:introd}, \\ref{ssec:rev}, and~\\ref{ssec:Rmass}).  Here we have extended the analysis of TK07 for BD-like binary fractions up to 60\\pct\\ for the Orion Nebula Cluster, the Taurus-Auriga association, IC~348 and the Pleiades by using slight modifications of the techniques introduced in TK07.  As a main result, we found that the discontinuity that comes about by treating BDs/VLMSs and stars consistently in terms of their observed multiplicity properties remains even for the highest BD binary fraction. These results suggest that the BD binary fraction, $\\fbd$, is not the dominant origin of the discontinuity in the IMF, and that, consequently, two separate IMFs need to be introduced.  It is re-emphasised that by seeking to mathematically describe the BD and stellar population in terms of the relevant mass- and binary distribution functions, it is unavoidable to mathematically separate BDs and VLMSs from stars. The two resulting mass distributions do not join at the transition mass near 0.08~\\tmsun. The physical interpretation of this logically stringent result is that BDs and VLMSs follow a different formation history or channel than stars. This result is obtained independently by theoretical consideration of star-formation processes \\citep{ReiCla01,GoWi07, StHuWi07,BCB08}.  With this contribution we have quantified how the power-law index of the BD-like IMF and the BD-to-star ratio changes with varying binary fraction of BD-like bodies. The BD-like power-law index, $\\abd\\approx0.3$, remains almost constant if equal-mass pairing of BD-like binaries is assumed, while $\\abd$ increases somewhat with increasing $\\fbd$ in the case of random pairing over the BD-like mass range. All values of $\\abd$ are between $-0.1$ and~$+1.3$. We also find that although the stellar density differs from a few stars per $\\mathrm{pc}^3$ (TA) to about 20000 stars per $\\mathrm{pc}^3$, the resulting $\\abd$ is constant within the uncertainties. Similarly, the BD-to-star ratio does not show a trend with increasing stellar density. This suggests the star-formation and BD-formation outcome to be rather universal at least within the range of densities probed here."
    },
    "0808/0808.0194_arXiv.txt": {
        "abstract": "Recent observations of UV-/optically selected, massive star forming galaxies at $z\\approx2$ indicate that the baryonic mass assembly and star formation history is dominated by continuous rapid accretion of gas and internal secular evolution, rather than by major mergers. We use the Millennium Simulation to build new halo merger trees, and extract halo merger fractions and mass accretion rates. We find that even for halos not undergoing major mergers the mass accretion rates are plausibly sufficient to account for the high star formation rates observed in $z\\approx2$ disks. On the other hand, the fraction of major mergers in the Millennium Simulation is sufficient to account for the number counts of submillimeter galaxies (SMGs), in support of observational evidence that these are major mergers. When following the fate of these two populations in the Millennium Simulation to $z=0$, we find that subsequent mergers are not frequent enough to convert all $z\\approx2$ turbulent disks into elliptical galaxies at $z=0$. Similarly, mergers cannot transform the compact SMGs/red sequence galaxies at $z\\approx2$ into observed massive cluster ellipticals at $z=0$. We argue therefore, that secular and internal evolution must play an important role in the evolution of a significant fraction of $z\\approx2$ UV-/optically and submillimeter selected galaxy populations. ",
        "introduction": "\\label{s:intro} In the cold dark matter model of hierarchical structure formation \\citep{BlumenthalG_84a,DavisM_85a,SpringelV_06a} mergers are believed to play an important role in galaxy formation and evolution \\citep{SteinmetzM_02a}. Mergers induce starbursts \\citep{HernquistL_95a} and transform galactic morphology \\citep{NaabT_03a}. Major mergers may drive the buildup of the red sequence \\citep{ToomreA_77a,HopkinsP_07b}. Dark matter models and many observations show that mergers are more frequent at high redshift \\citep{FakhouriO_07a,ConseliceC_03b}.  However, there is growing evidence that a smoother growth mode may be important for the baryonic mass assembly and star formation history at high redshift. For example, the tight correlation between star formation rate (SFR) and stellar mass in UV-/optically selected star forming galaxies is indicative of buildup by continuous gas inflow \\citep{DaddiE_07a,NoeskeK_07a}. As part of the SINS survey (\\citealp{FoersterSchreiberN_06a}; N.~M.~F\\\"orster Schreiber et al.~2008 in preparation), integral field spectroscopy of more than $50$ UV-/optically selected $z\\approx2$ star forming galaxies show a preponderance of thick gas-rich rotating disks and only a minority of major mergers \\citep{FoersterSchreiberN_06a,GenzelR_06a,GenzelR_08a,ShapiroK_08a}. In contrast, SMGs are probably short-lived maximum-starburst galaxies undergoing dissipative major mergers \\citep{TacconiL_06a,BoucheN_07a,TacconiL_08a}. Table \\ref{t:galaxies} summarises key properties of these $z\\approx2$ galaxy samples.  How do these observations fit into the concordance cosmological model? Modern simulations of dark matter structure formation are robust and fixed by the cosmological parameters. However, complicated baryonic physics makes it difficult to model the evolution of galaxies and reproduce, for example, the high SFRs of these $z\\approx2$ galaxies \\citep{DaddiE_07a}.  Galaxies at $z\\approx2$ differ significantly from local galaxies. The central mass densities of SMGs and of massive quiescent galaxies at the same redshift (\\citealp{vanDokkumP_07a} and references therein) are an order of magnitude greater than those of local spheroids and disks \\citep{TacconiL_08a}. Also, the $z\\approx2$ rotating disks are thick and turbulent, unlike local disk galaxies. These differences raise the question: what are the local Universe descendants of these $z\\approx2$ galaxies?  In this paper we use the cosmological dark matter Millennium Simulation (\\citealp{SpringelV_05a}; \\S\\ref{s:find_merger}) to investigate the possible role of major mergers in galaxy formation at $z\\approx2$ (\\S\\ref{s:z_2}), and to consider the evolution of the $z\\approx2$ galaxies to $z=0$ (\\S\\ref{s:evolution}).   \\begin{table}[tbp] \\caption{Properties of galaxy samples at $z\\approx2$} \\label{t:galaxies} \\centering \\begin{tabular}{c| c c c c} \\hline\\hline Galaxy & SFR & Halo & Comoving number & Major \\\\ sample & $[\\Msunyr]$ & mass & density & merger \\\\ && $[\\Msun]$ &[$h_{0.7}^3\\Mpc^{-3}$]& fraction \\\\ \\hline SINS & $\\approx30-300$\\tablenotemark{(a)} & $10^{11.84}v_{200}^3\\times$ & $1-2.2\\times10^{-4}$\\tablenotemark{(c,d)} & $\\approx0.3$\\tablenotemark{(e)} \\\\ &&$(\\frac{1+z}{3.2})^{-1.5}h_{0.7}^{-1}$\\tablenotemark{(b)}&&\\\\ SMGs\\tablenotemark{(f)} & $\\approx750\\pm300$\\tablenotemark{(g)} & - & $1-2\\times10^{-5}$\\tablenotemark{(c)} & $\\approx1$\\tablenotemark{(c)} \\\\ \\hline \\end{tabular} \\tablecomments{(a) N.~M.~F\\\"orster Schreiber et al.~2008 in preparation ; (b) \\citealp{FoersterSchreiberN_06a} ; (c) \\citealp{TacconiL_08a} and references therein ; (d) BX/BM \\& sBzK galaxies with $K\\leq20$ ; (e) \\citealp{ShapiroK_08a} ; (f) $S(850{\\rm \\mu m})\\geq5{\\rm mJy}$  ; (g) Genzel, priv.~comm.} \\vspace{0.5cm} \\end{table} ",
        "conclusions": "\\label{s:summary} We have constructed new halo merger trees from the \\LCDM{ }Millennium Simulation. Our trees account for merger durations, and we use them to identify halos that are undergoing mergers and to extract dark matter accretion rates. We show that the high star formation rates observed in rotating disks at $z\\approx2$ are plausibly consistent with the dark matter accretion rates expected for halos not undergoing major mergers. Given the measured star formation rates of SMGs, and the observationally supported assumption that they are undergoing major mergers, we infer their likely halo masses. Major mergers can neither lead to the complete transformation of the $z\\approx2$ disks to $z=0$ ellipticals, nor to the disappearance of the very high density SMGs from $z\\approx2$ to $z=0$. Therefore, secular/internal processes are likely important in the evolution of these high-redshift populations to present time."
    },
    "0808/0808.1067_arXiv.txt": {
        "abstract": "{Two classes of gamma-ray bursts have been identified in the BATSE catalogs characterized by durations shorter and longer than about 2 seconds. There are, however, some indications for the existence of a third class. Swift satellite detectors have different spectral sensitivity than pre-Swift ones for gamma-ray bursts. Therefore we reanalyze the durations and their distribution and also the classification of GRBs. } {We analyze the bursts duration distribution, published in The First BAT Catalog, whether it contains two, three or more groups. } { Using The First BAT Catalog the maximum likelihood estimation was used to analyze the duration distribution of GRBs.} { The three log-normal fit is significantly (99.54\\% probability) better than the two for the duration distribution. Monte-Carlo simulations also confirm this probability (99.2\\%). Similarly, in previous results we found that the fourth component is not needed. The relative frequencies of the distribution of the groups are 7\\% short 35\\% intermediate and 58\\% long.  } { Similarly to the BATSE data, three components are needed to explain the BAT GRBs' duration distribution. Although the relative frequencies of the groups are different than in the BATSE GRB sample, the difference in the instrument spectral sensitivities can explain this bias. This means theoretical models may be needed to explain three different type of gamma-ray bursts.} ",
        "introduction": "It has been a great challenge to classify  gamma-ray bursts (GRBs). \\cite{maz81} and \\cite{nor84} suggested there might be a separation in their duration distribution. Using The First BATSE Catalog, Kouveliotou et al. (1993) found a bimodality in the distribution of the logarithms of the durations. In that paper they used the parameter $T_{90}$ (the time in which 90\\% of the fluence is accumulated \\citep{kou93}) to characterize the duration of GRBs \\citep{mcb94,kos96,belli97,pen97}. Today it is widely accepted that the physics of these two groups (also called \"subclasses\" or simply \"classes\") are different, and these two kinds of GRBs are different phenomena \\citep{nor01,bal03,fox05}. In the Swift database the measured redshift distribution for the two groups are also different, for short bursts the median is 0.4 \\citep{os08} and for the long ones it is 2.4 \\citep{bag06}.  The bimodal distribution was further quantified in another paper \\citep{kou96}, where a two-log-normal fit was made; the best parameters of the fit were published in \\cite{mcb94} and \\cite{kos96}.  In a previous paper using the Third BATSE Catalog \\citep{m6} \\cite{hor98} showed that the  duration ($T_{90}$) distribution of GRBs observed by BATSE could be well fitted by a sum of three log-normal distributions. We find it statistically unlikely (with a probability $\\sim 10^{-4}$) that there are only two groups.  Simultaneously, Mukherjee et al. (1998) report the finding (in a multidimensional parameter space) of a very similar group structure of GRBs.  Somewhat later, several authors \\citep{hak00,bala01,rm02,hak03,bor04,hak04,chat07} included more physical parameters in the analysis of the bursts (e.g. peak-fluxes, fluences, hardness ratios, etc.). A cluster analysis in this multidimensional parameter space suggests the existence of the third (\"intermediate\") group as well \\citep{muk98,hak00,bala01,rm02,chat07}. The physical existence of the third group is, however, still not convincingly proven. However, the celestial distribution of the third group is anisotropic \\citep{mesz00,li01,mgc03}.  All these results  mean that the existence of the third intermediate group in the BATSE sample is acceptable, but its physical meaning, importance and origin is less clear than those of the other groups. Hence, it is worth studying new samples if their size is large enough for statistical analysis. In the HETE-II database \\citep{van04} there are only 104 GRBs and in the Swift first BAT database \\citep{sak08} there are 237 GRBs.  Therefore, in this paper we use the Swift data because of its better statistics.   \\begin{figure} \\centering \\resizebox{\\hsize}{!} {\\includegraphics[angle=0,width=7cm]{t90.eps}} \\caption{Duration distribution of the observed BAT bursts.  } \\end{figure}  In Sect. 2 we discuss the  method used in the paper. In Sect.  3 uni-, bi-, tri- and tetra-modal log-normal fits made by using the maximum likelihood method are discussed. In Sect. 4 one thousand Monte-Carlo simulations are shown investigating the significance of the fits. In Sect. 5 we discuss some further details. The conclusions are given in Sect. 6. ",
        "conclusions": "\\begin{enumerate}  \\item Assuming that the $T_{90}$ distribution of the short and long  GRBs is log-normal,  the probability that the third group is a chance occurance is about 0.5-0.8 \\%.  \\item Although the statistics indicate that a third component is present, the physical existence of the third group is still debatable. The sky distribution of the third component is anisotropic as proven by \\cite{mesz00} and \\cite{li01}. Alternatively \\cite{hak00} believe the third statistically proven subgroup is only a deviation caused by complicated instrumental effects, which can reduce the duration of some faint long bursts. This paper does not deal with this particular effect, however the previously studied BATSE sample shows a similar group structure. This agreement suggests that the third component is possibly real, not an instrumental effect (the BATSE detectors and the Swift BAT are different kinds of instruments).  \\item The observed  frequencies in the three classes are different for BATSE and BAT.  Both samples are dominated, however, by the long bursts. The short bursts are less populated in BAT than in BATSE but the intermediate group is more numerous. This is understandable, since BAT is less sensitive in high energy than BATSE was and more sensitive in low energy and short bursts are the hardest group and intermediate ones are the softest. Therefore BAT can observe more intermediate bursts and much fewer short ones than BATSE did.  \\item The existence and physical properties of the intermediate group need further discussion to elucidate the reality and properties of this class of GRBs.   \\end{enumerate}"
    },
    "0808/0808.2841_arXiv.txt": {
        "abstract": "Data obtained in the very high energy $\\gamma$-ray band with the new generation of imaging telescopes, in particular through the galactic plane survey undertaken by H.E.S.S., low threshold observations with MAGIC and more recently by operation of VERITAS, have revealed dozens of galactic and extragalactic sources, providing a wealth of information on a variety of high energy acceleration sites in our universe. Also, the water Cherenkov instrument Milagro has provided several extended sources after seven years of data integration. An overview of these results with focus on some of the most recent highlights is given. ",
        "introduction": "\\label{intro} Almost two decades after the establishment of the Crab nebula as the first TeV emitting source, thanks to the pioneering work at the Whipple Observatory in Arizona~\\cite{Weekes89}, the field of very high energy (VHE) $\\gamma$-ray astronomy has entered the maturity age. Although limited in aperture (3-6$\\times10^{-3}$ sr, currently) and duty cycle ($\\sim 10\\%$), the imaging atmospheric Cherenkov telescopes (IACTs) have proved to be the most sensitive, thanks to precise angular reconstruction of the shower origin and to powerful background rejection capabilities. The new generation of these instruments has allowed the discovery of more than 70 galactic and extragalactic sources. The major contributions have come from H.E.S.S. in Namibia (system of four 13~m diameter telescopes, 2003), MAGIC on La Palma (single 17 m diameter dish, 2004), CANGAROO-III in Woomera (three 10 m diameter dishes, one more pending, 2004), and  VERITAS (four 12 m diameter telescopes, 2006) in Arizona. Large aperture and duty cycle instruments, the water Cherenkov detector Milagro (1999-2007) and the particle counter array Tibet As-$\\gamma$ (1992), although having a much lower sensitivity, have also interestingly contributed to the field through large exposures obtained after few-years scale integration times. A brief overview of the field, as of early 2008, is given below. ",
        "conclusions": "The field of VHE $\\gamma$-ray astrophysics has gone through a dramatic evolution since 2004, thanks to the high sensitivity of the new generation IACTs. The HESS GPS represents a major step in that it has revealed, beyond the large number of sources, diverse classes of $\\gamma$-ray emitting galactic objects and acceleration sites: young shell-type SNRs, SNRs interacting with molecular clouds, middle-aged offset PWNe, very young composite PWNe and $\\gamma$-ray binaries. Given the large number of still unidentified sources, other potential classes of sources could emerge, including the promising case of massive star clusters. The increasing number of blazar sources in the extragalactic domain allows now for population studies, and one non-blazar source, M~87 is under scrutiny, in particular  by VERITAS. Also, while the early attempts to constrain the intergalactic radiation field suffered from the very limited number of sources and a reduced range in redshift, the growing number of objects, and especially the detection of 3C~279 obtained at a low energy threshold by MAGIC, have definitely opened the path towards the cosmological application of $\\gamma$-ray astrophysics. There is no doubt that VHE $\\gamma$-ray  astronomy is now a genuine branch of astronomy with multiple connections to cosmology and fundamental physics."
    },
    "0808/0808.0769_arXiv.txt": {
        "abstract": "We investigate the correlation between the mass of the supermassive black holes (SMBHs) and metal abundance, using existing data sets. The SMBH mass  $M_{bh}$ is well correlated with integrated stellar feature of Mgb. For 28 galaxies, the best-fit $M_{bh}$-Mgb relation has a small scatter, which is an equivalent level with other well-known relation, such as a correlation between the stellar velocity dispersion and the mass. An averaged iron index $\\langle $Fe$\\rangle $ also positively correlates with $M_{bh}$, but the  best-fit $M_{bh}$-$\\langle $Fe$\\rangle $ relation has a larger scatter. The difference comes from the synthesis and evolution mechanisms, and may be important for the SMBH and star formation history in the host spheroid. ",
        "introduction": "Recent observations of nearby massive spheroids (ellipticals, lenticular and spiral bulges) have established that supermassive black holes (SMBHs) are present in the nuclei of galaxies. The mass $M_{bh}$ of SMBHs ranges from $10^{6}$ to $10^{9}M_{\\odot }$. Several statistical correlations with characteristic parameters of the host spheroids are explored: relation with the luminosity $L$\\citep{KR95,MH03,G07}, the mass $M_{s}$ of the spheroids\\citep{Ma98,HR04}, the stellar velocity dispersion $\\sigma $\\citep{Ge00,FM00,Tr02}, the light concentration or S\\'{e}rsic index $n$ \\citep{Gr01,GD07}, gravitational binding energy $E_{g}$ \\citep{AR07} and so on. The recent updated survey for various galaxy parameters is given in \\citet{AR07}. If the fitting models are good, the correlations may provide some clues to the coevolution of SMBHs and the host spheroids. Some tentative theories have been proposed to clarify the origin of these correlations, e.g., \\citet{SR98,U01}.  The galaxies have some aspects as an assemble of stars, gases and dark matters. The quantities $M_{s}$, $\\sigma $, $E_{g}$ represent the dynamical aspect, while $L$, $n$ represent the photometric one, although they are closely related each other. Very little attention is paid to the chemical one. The observation is not so easy and therefore the correlation between SMBH mass and spheroid chemical property is not well studied so far. Both the metal abundance and black hole mass increase with time and we therefore expect some correlations between them. It is important to clarify the correlation, and the extent, if any. The chemical parameter using Lick/IDS absorption line indices, e.g., \\citet{Wo94}, especially Mg and Fe is systematically studied in literatures\\citep{Tr00,De05,Ku06} for the galaxies in which presence of SMBH is reliable. Using the published data, we investigate the correlation with SMBH mass in this paper.  As a sample we select only 28 galaxies, for which SMBH masses are measured by dynamical methods. In section 2, we describe the sample of galaxies and criterion of our choice.  In section 3, we derive the correlation between $M_{bh}$ and $\\sigma $, and that between $M_{bh}$ and the B-band magnitude $M_{B}$ from our sample. These correlations are known to be good, and are used in order to check no serious bias involved in the sample. Our results for $M_{bh}$-$\\sigma $ and $M_{bh}$-$M_{B}$ relations are not so different from those by \\citet{AR07,G07}, although the source sample is not the same.  We also show the correlation between black hole mass and metal abundance index using our sample. Finally, a discussion of our results is given in section 4. ",
        "conclusions": ""
    },
    "0808/0808.3847_arXiv.txt": {
        "abstract": "The dynamics of stars in the inner regions of nearby globular clusters (GCs) such as G1 indicate the presence of central concentrated dark masses, and one would like to know whether these are indeed intermediate mass black holes (IMBHs). As the number of surrounding stars, and their motions, are roughly known, the capture rate can be estimated; the question then arises of whether the apparent quiescence of the nuclei of these GCs is compatible with the IMBH's presence. The role of debris from disrupted stars in {\\it activating} quiescent nuclei of GCs is studied here employing three-dimensional hydrodynamics simulations. It is argued that when individual stars are disrupted, the bulk of the debris would be swallowed or expelled rapidly compared with the interval between successive disruptions.  A transient (predominantly of soft X-ray emission) signal could persist steadily with $L\\sim L_{\\rm Edd}=10^{41}(M_{\\rm h}/10^3M_\\odot)$ erg/s for at most tens of years; thereafter the flare would rapidly fade. While the infall rate declines as $t^{-5/3}$, some material may be stored for a longer time in an accretion disk. The IMBH luminosity could then remain as high as $L_X\\ge 10^{39}$ erg/s for several hundreds of years after disruption.  In a given object, this ultraluminous X-ray activity would have a duty cycle of order $10^{-4}$.  Quiescent GCs, those with IMBHs now starved of fuel, should greatly outnumber active ones; observational constraints would not then be stringent until we had observed enough candidates to constitute a proper ensemble average. ",
        "introduction": "Suggestive evidence has accumulated that intermediate mass black holes exist in some globular clusters. There is some dynamical evidence for a mass concentration within the central regions of some globular clusters.  The dynamics of stars in the inner regions of nearby clusters such as M15, G1 and $\\omega$ Centauri suggest the presence of black holes with masses of about $10^{3}$, $10^{4}$ and $4\\times 10^4\\;M_\\odot$, respectively \\citep{ge2002,ge2003,geb2002,geb2005,po2006,noyola08}. The peculiarities of an unresolved radio source in G1 indicate some unique object at the center \\citep{ul2007}. Recently, somewhat arguable evidence has arisen for the presence of IMBHs in young star clusters, where ultraluminous, compact X-ray sources (ULXs) have been preferentially found to occur \\citep{ze2002}. Their high luminosities suggest that they are IMBHs rather than binaries containing a normal stellar mass black hole \\citep{po2004}. But before accepting this conclusion (and dismissing alternative ideas) one would like some independent corroboration of the IMBH hypothesis, or, conversely, some way of ruling it out.  We are used to the idea that black holes are implicated in the most powerful sources in the universe, and can (when accreting) be ultra-efficient radiators. But there is, for example, no sign of such activity in M15 - the X-ray upper limit is no more than $6 \\times 10^{32}$ erg/s \\citep{ho2003}. So could a black hole be so completely starved of fuel that it does not reveal its presence? We do not directly know how much gas there is near the IMBH \\citep{p00,po2006}, and there is no a priori reason why this region should be {\\it swept clean} of gas.  Gas that is lost from nearby stars \\citep{ba2006} as well as mass transfer from potentially bound companions \\citep{bl2006,pa2006,hopman04} can produce observable signatures. However, it is uncertain whether gas can be adequately supplied to explain ULX activity. The star density, on the other hand, is much better known -- after all, if the stars were not closely packed near the center of the cluster we would not have evidence for the black hole at all.  As stellar orbits diffuse in phase space, it therefore seems inevitable that some may wander sufficiently close to the hole that they suffer tidal disruption.  When a star is disrupted, there is bound to be some radiation from the sudden release of gas. The flares resulting from a disrupted star could be the clearest diagnostic of a IMBH's presence.  It is an intricate although tractable problem in stellar dynamics to calculate the chance that a star passes within the IMBH's tidal radius \\citep{fr1976}. In a simple case when the velocities are isotropic, the frequency with which a star would enter the zone of vulnerability is $\\xi \\sim 10^{-7} M_{\\rm h,3}^{4/3}(n_*/10^{6}\\;{\\rm pc^{-3}})(\\sigma/ {\\rm 10 \\;km\\; s^{-1}})^{-1}\\;{\\rm yr^{-1}}$, where $n_*$ is the stellar number density in the cluster nucleus and $M_{\\rm h,3}$ is the black hole's mass in units of $10^3 M_\\odot$. This estimate although simplified, agrees well with the fiducial rates derived from detailed $N-$body simulations of multi-mass star clusters containing IMBHs \\citep{ba2004}.  It may seem, however, that even the modest rate of stellar disruptions given above could have conspicuous consequences. The debris from a disrupted one solar-mass star per hundred million years, swallowed steadily with 10 per cent radiative efficiency, would yield a luminosity of $L_{\\rm steady} \\sim 6\\times 10^{37}$ erg/s -- higher than is observed for X-ray binaries in quiescence.  But in reality, one expects brighter flares with short duty cycles, as the time it takes to digest or expel the debris from one star is much shorter than the mean interval between one stellar disruption and the next.  This {\\it Letter} is concerned with the observational manifestations of such phenomena, with particular reference to globular clusters where the masses of the black holes (if indeed present) are perhaps of order $10^3-10^4M_\\odot$. ",
        "conclusions": "\\subsection{Observability of  $L_{\\rm Edd}$ Flares } A distinctive consequence of a $10^3-10^4\\;M_\\odot$ IMBH's presence in the centers of globular clusters would be the transient flares produced as the bound debris from the disrupted solar-type stars is swallowed, the luminosities being as high as $L_{\\rm Edd}=10^{41}M_{\\rm h,3}\\;{\\rm erg\\;s^{-1}}$. The rise and the peak bolometric luminosity can be predicted with some confidence. However, the effective surface temperature (and thus the fraction of luminosity that emerges predominantly in the soft X-ray band) is harder to predict, as it depends on the size of the effective photosphere that shrouds the hole. For $t\\lesssim t_{\\rm Edd}$, the effective surface temperature should be $T \\lesssim T_{\\rm Edd} = (L_{\\rm Edd}/4\\pi \\sigma R_{\\rm g}^2) ^{1/4}\\sim 1 M_{\\rm h.3}^{-1/4}$ keV. In a given globular cluster, these flares would have a duty cycle of order $\\xi t_{\\rm Edd} \\sim 10^{-6} (\\xi /10^{-7}\\;{\\rm yr^{-1}})(t_{\\rm Edd}/10 {\\rm yr})$, and as a result, quiescent GCs, those with IMBHs now starved of fuel, should greatly outnumber active ones.  Given a globular cluster space density of $n_{\\rm gc} \\sim 4\\;{\\rm Mpc}^{-3}$ \\citep{br2006}, we expect the density rate of flares with $L=L_{\\rm Edd}$ to be at most $\\sim 4000 {\\rm yr^{-1}}\\;{\\rm Gpc}^{-3}$, if $10^3-10^4\\;M_\\odot$ black holes were prevalent in globular clusters. Therefore we would not yet expect to have detected such a flare. Observational constraints on the presence of IMBH would not then be stringent until a sufficiently large sample of galaxies should be sampled in order to reveal some of their GC members in a flaring state. Such objects should be searched for out to large distances. However, assuming $L \\sim 10^{41}$ erg s$^{-1}$ and $T \\sim T_{\\rm Edd}$, we find that EXIST \\citep{grindlay04} would only be able to detect such a flaring event out to a distance of about 1 Mpc.  \\subsection{ULX activity} A further question is how fast the luminosity fades after the flare. This is important because we want to know whether the IMBH has faded below ULX levels before the next stellar disruption occurs. The answer to this question depends on how long it takes the last residues of the star to be ingested. One expects the infall rate $\\dot{M}$ to decline as $t^{-5/3}$ for late times. Some material may, however, be stored for a longer time in an accretion disk. This is mainly because the specific angular momentum for Keplerian orbits grows $\\propto r^{1/2}$, so angular momentum transport via disk viscosity requires that 10\\% of the debris goes out to $10^3\\;R_{\\rm T}$ and 1\\% to $10^5\\;R_{\\rm T}$ before being swallowed.  The evolution of the bound debris onto a black hole for $t\\geq t_{\\rm a}$ has been studied by \\citet{ca1990} using a time dependent $\\alpha$-disk model. Cannizzo and collaborators concluded that the black hole's luminosity fades more slowly: $L \\propto t^{-1.2}$ and depends only weakly on the disk viscosity. The IMBH luminosity could then remain as high as $L_X\\ge 10^{39}\\;{\\rm erg\\;s^{-1}}$ for several hundreds of years after disruption.  In a given object, this ULX activity would have a duty cycle of order $10^{-4}$ unless the rate of disruptions is much larger than $10^{-7}$ or the viscosity were so low that the bulk of the mass could be stored for many thousands of years at $R\\sim R_{\\rm T}$ (where the dynamical timescale is only a few hours).  \\subsection{Relevance to GCs and  IMBH Growth} The mass fraction that is ejected rather than ingested, although less spectacular than the accretion-powered flares could nonetheless have an important influence on the energy balance within a GC - similar than a supernova exploding in the same volume. When the star is disrupted in a single flyby, about half the debris is ejected in a fast moving spray of gas; the kinetic energy output being $\\sim 10^{50} M_{\\rm h,3}^{1/3}\\,{\\rm erg}$. The ejecta would be braked and their kinetic energy thermalized, as they ran into the diffuse gas (probably originating from stellar mass loss). If this material is effectively shock heated, it could contribute to the X-ray and radio luminosity of the GC long after the star has been disrupted. Stellar disruption may have other implications -- for instance, to the dynamics of line emitting regions within GCs \\citep{ch2008}.  The most tightly bound debris, on the other hand, would traverse an elliptical orbit before returning to $R_{\\rm T}$.  There would therefore be no conspicuous flare until, as discussed above, the bound debris fall back into the hole. Electron scattering opacity almost certainly dominates the radiative transfer and the photons will be trapped out to a radius $R_\\tau =\\dot{M}\\kappa_\\tau/4\\pi c$. Energy will be dissipated at a supercritical rate as the material swirls closer to the hole, some gas may then be ejected in a radiation-driven wind. In principle it is possible for only a fraction $\\epsilon^{-1} (R_{\\rm g}/R_{\\rm T}) \\sim 0.004 \\epsilon_{0.1}(R_\\ast/R_\\odot)^{-1}(M_\\ast/M_\\odot)^{1/3}M_{\\rm h,3}^{2/3}$ to be actually swallowed, all the remainder being ejected. If a fraction $f$ of the mass from disrupted stars is accreted onto the black hole, the average rate of tidal disruptions required to form an IMBH out of a $\\sim 50 M_\\odot$ progenitor would need to be $\\ge 10^{-6} M_{\\rm h,3} (f/0.1)^{-1} (\\tau_{\\rm GC}/10^{10}{\\rm yr})^{-1} {\\rm yr}^{-1}$, where $\\tau_{\\rm GC}$ is the globular cluster age. In such cases, the density would, however, be high enough that runaway merging of high-mass main-sequence stars could lead directly to the formation of an IMBH \\citep{ba2006,po2004}."
    },
    "0808/0808.2466_arXiv.txt": {
        "abstract": "{ The BL\\,Lac object \\rgb\\ ($z=0.080$) was predicted to be a very high-energy (VHE; $> 100$\\,GeV) $\\gamma$-ray source, due to its high X-ray and radio fluxes.  We report recent observations of this source made in late October and November 2007 with the H.E.S.S. array consisting of four imaging atmospheric \\v{C}erenkov telescopes. Contemporaneous observations were made in X-rays with the {\\textit{Swift\\/}} and {\\textit{RXTE\\/}} satellites, in the optical band with the ATOM telescope, and in the radio band with the Nan\\c{c}ay Radio Telescope.  As a result, \\rgb\\ is discovered as a source of VHE $\\gamma$-rays by H.E.S.S. A signal of 173 $\\gamma$-ray photons corresponding to a statistical significance of 6.6\\,$\\sigma$ was found in the data. The energy spectrum of the source can be described by a powerlaw with a spectral index of \\spectralindex. The integral flux above 300\\,GeV corresponds to $\\sim$2\\% of the flux of the Crab nebula. The source spectral energy distribution (SED) can be described using a two-component (extended jet and blob in jet) non-thermal synchrotron self-Compton (SSC) leptonic model, plus a thermal host galaxy component. The parameters that are found are very close to those found for TeV blazars in similar SSC studies.  The location of its synchrotron peak, as derived from the SED in {\\textit{Swift\\/}} data, allows clear classification as a high-frequency-peaked BL\\,Lac (HBL). }  \\FullConference{Workshop on Blazar Variability across the Electromagnetic Spectrum\\\\ April 22-25 2008\\\\ Palaiseau, France}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "The BL\\,Lac \\rgb\\ was detected in VHE at energies $> 300$\\,GeV with H.E.S.S. The contemporaneous multi-wavelength observations allow the SED of \\rgb\\ to be derived for the first time, clearly confirming its high-frequency-peaked nature at the time of H.E.S.S. observations."
    },
    "0808/0808.0655_arXiv.txt": {
        "abstract": "{ Recently, the MID-infrared Interferometric instrument (MIDI) at the VLTI has shown that dust tori in the two nearby Seyfert galaxies NGC~1068 and the Circinus galaxy are geometrically thick and can be well described by a thin, warm central disk, surrounded by a colder and fluffy torus component. By carrying out hydrodynamical simulations with the help of the TRAMP code \\citep{Klahr_99}, we follow the evolution of a young nuclear star cluster in terms of discrete mass-loss and energy injection from stellar processes. This naturally leads to a filamentary large scale torus component, where cold gas is able to flow radially inwards. The filaments open out into a dense and very turbulent disk structure. In a post-processing step, we calculate observable quantities like spectral energy distributions or images with the help of the 3D radiative transfer code MC3D \\citep{Wolf_03}. Good agreement is found in comparisons with data due to the existence of almost dust-free lines of sight through the large scale component and the large column densities caused by the dense disk. ",
        "introduction": "Within the current model of Active Galactic Nuclei (AGN), two observed classes of galaxies can be explained by one intrinsically unique AGN model, the so-called {\\it Unified Scheme of Active Galactic Nuclei}: a hot accretion disk surrounds the central black hole and crosses over into a geometrically thick torus. Nowadays, it is thought that the gap between these two is filled by the so-called Broad Line Region, see \\citet{Gaskell_08}. The optically thick torus blocks the sight towards the centre, when viewed edge-on and enables the viewing of the characteristic emission peak of the accretion disk and the Broad Line Region clouds when viewed face-on. Numerous observations back this model and also hint towards a clumpy nature of the gas and dust torus. ",
        "conclusions": "A physical model for the built-up and evolution of gas tori in Seyfert galaxies is presented. We follow the evolution of a young nuclear star cluster, during the phase, where the ejection of planetary nebulae is the dominant process for mass input. Together with supernova energy input and the onset of cooling instability yield a large scale, filamentary torus component with cavities of hot gas in-between. Within the filaments, gas cools and gets accreted in radial direction. Close to the angular momentum barrier, it opens out into a dense disk component with turbulent character due to the inflow of filaments and clumps of gas from random directions. This two component structure enables a good comparison with observations."
    },
    "0808/0808.0713_arXiv.txt": {
        "abstract": "We compare waveforms and orbital dynamics from the first long-term, fully non-linear, numerical simulations of a generic black-hole binary configuration with post-Newtonian predictions. The binary has mass ratio $q\\sim0.8$ with arbitrarily oriented spins of magnitude $S_1/m_1^2\\sim0.6$ and $S_2/m_2^2\\sim0.4$ and orbits 9 times prior to merger. The numerical simulation starts with an initial separation of $r\\approx11M$, with orbital parameters determined by initial 2.5PN and 3.5PN post-Newtonian evolutions of a quasi-circular binary with an initial separation of $r=50M$. The resulting binaries have very little eccentricity according to the 2.5PN and 3.5PN systems, but show significant eccentricities of $e\\sim0.01-0.02$ and $e\\sim0.002-0.005$ in the respective numerical simulations, thus demonstrating that 3.5PN significantly reduces the eccentricity of the binary compared to 2.5PN. We perform three numerical evolutions from $r\\approx11M$ with maximum resolutions of $h=M/48,M/53.3,M/59.3$, to verify numerical convergence. We observe a reasonably good agreement between the PN and numerical waveforms, with an overlap of nearly 99\\% for the first six cycles of the $(\\ell=2,m=\\pm2)$ modes, 91\\% for the $(\\ell=2,m=\\pm1)$ modes, and nearly 91\\% for the $(\\ell=3,m=\\pm3)$ modes. The phase differences between numerical and post-Newtonian approximations appear to be independent of the $(\\ell,m)$ modes considered and relatively small for the first 3-4 orbits. An advantage of the 3.5 PN model over the 2.5 PN one seems to be observed, which indicates that still higher PN order (perhaps even 4.0PN) may yield significantly better waveforms. In addition, we identify features in the waveforms likely related to precession and precession-induced eccentricity. ",
        "introduction": "\\label{sec:intro}  The discoveries of quasars, AGN, and other black-hole driven astrophysical phenomena in the 1960's demonstrated that the most energetic astrophysical phenomena are powered by gravity in the strong-field regime. This, in turn, spurred a renewed interest in classical General Relativity. The second major milestone in the revival of the theory was the realization that when astrophysical black holes merge, they release incredible amounts of energy in the form of gravitational radiation, making them the brightest objects in the universe.  During their last few orbits, merging black-hole binaries release energy with a peak luminosity of about $10^{-3}c^5/G$, $10^{23}$ times the power output of the Sun.  There are currently major experimental and theoretical efforts underway to measure these gravitational wave signals. On the experimental side, these efforts required the construction of kilometers long interferometers, such as LIGO~\\cite{LIGO3} and VIRGO~\\cite{Acernese:2004ru}, sensitive enough to measure arm length distance changes smaller than the radius of a proton. While on the theoretical side, these efforts required major advancements in signal extraction techniques and the theoretical modeling of the gravitational wave sources. Modeling the gravitational radiation from compact object sources has been particularly difficult, as they require solving the fully non-linear Einstein Equations of General Relativity on powerful supercomputers. However, even with the rapid advancements in computer power, solving the two-body problem in General Relativity proved to be remarkably difficult, requiring over thirty years of research for the field to mature. Then in 2005, two complementary and independent methods were discovered that allowed numerical relativists to finally solve the black-hole binary problem in full strong-field gravity~\\cite{Pretorius:2005gq, Campanelli:2005dd, Baker:2005vv}.  The rapid progress and the number of new theoretical insights that followed these breakthroughs have transformed the field of numerical relativity (NR); turning it into a very valuable tool with significant impact on astrophysics~\\cite{Campanelli:2004zw, Herrmann:2006ks, Baker:2006vn, Sopuerta:2006wj, Gonzalez:2006md, Sopuerta:2006et, Herrmann:2006cd, Herrmann:2007zz, Herrmann:2007ac, Campanelli:2007ew, Koppitz:2007ev, Choi:2007eu, Gonzalez:2007hi, Baker:2007gi, Campanelli:2007cga, Berti:2007fi, Tichy:2007hk, Herrmann:2007ex, Brugmann:2007zj, Schnittman:2007ij, Krishnan:2007pu, HolleyBockelmann:2007eh, Pollney:2007ss, Dain:2008ck, Redmount:1989, Merritt:2004xa, Campanelli:2007ew, Gualandris:2007nm, HolleyBockelmann:2007eh, Kapoor76, Bogdanovic:2007hp, Loeb:2007wz, Bonning:2007vt, HolleyBockelmann:2007eh, Komossa:2008qd, Komossa:2008ye}, gravitational wave detection~\\cite{Campanelli:2006gf, Baker:2006yw, Campanelli:2006uy, Campanelli:2006fg, Campanelli:2006fy, Pretorius:2006tp, Pretorius:2007jn, Baker:2006ha, Bruegmann:2006at, Buonanno:2006ui, Baker:2006kr, Scheel:2006gg, Baker:2007fb, Marronetti:2007ya, Pfeiffer:2007yz}, and on our theoretical understanding of black-binary spacetimes~\\cite{Sperhake:2008ga, Hannam:2006vv, Hannam:2008sg, Brown:2007tb, Garfinkle:2007yt, Dain:2008ck, Krishnan:2007pu, Campanelli:2006fy, Campanelli:2008dv}.  One of the breakthrough methods, the `moving puncture' approach~\\cite{Campanelli:2005dd, Baker:2005vv}, was adopted by a majority of the NR groups and has proven to be accurate for the neutron-star binary and mixed neutron-star---black-hole binary problems~\\cite{Etienne:2007jg, Yamamoto:2008js}, as well as for black-hole configurations with more than two black holes~\\cite{Campanelli:2007ea, Lousto:2007rj}.  On the subject of black-hole binaries, the NR community is in very good agreement concerning a variety of results. Black-hole binaries will radiate between $2\\%$ and $8\\%$ of their total mass and up to $40\\%$ of their angular momenta, depending on the magnitude and direction of the  spin components, during the last few orbits and merger~\\cite{Campanelli:2006uy, Campanelli:2006fg, Campanelli:2006fy, Dain:2008ck}. In general, these binaries will radiate net linear momentum, causing the final remnant black hole to recoil~\\cite{Campanelli:2004zw,  Herrmann:2006ks, Baker:2006vn, Sopuerta:2006wj, Gonzalez:2006md, Sopuerta:2006et, Herrmann:2006cd, Herrmann:2007zz, Herrmann:2007ac, Campanelli:2007ew, Koppitz:2007ev, Choi:2007eu, Gonzalez:2007hi, Baker:2007gi, Campanelli:2007cga, Berti:2007fi, Tichy:2007hk, Herrmann:2007ex, Brugmann:2007zj, Schnittman:2007ij, Krishnan:2007pu, HolleyBockelmann:2007eh, Pollney:2007ss, Dain:2008ck}. These recoils can be very large when the black holes in the binary have significant spin components in the orbital plane~\\cite{Campanelli:2007ew, Gonzalez:2007hi, Campanelli:2007cga, Healy:2008js} (up to $4000\\ \\KMS$ for astrophysical binaries~\\cite{Campanelli:2007cga} and even $10000\\ \\KMS$ for extremely close hyperbolic encounters~\\cite{Healy:2008js}), which has astrophysically important effects~\\cite{Redmount:1989, Merritt:2004xa, Campanelli:2007ew, Gualandris:2007nm, HolleyBockelmann:2007eh, Kapoor76}.  The observational consequences of these large recoil velocities is an active area of current research~\\cite{Bogdanovic:2007hp, Loeb:2007wz, Bonning:2007vt, HolleyBockelmann:2007eh, Komossa:2008qd, Komossa:2008ye}.   Currently, one of the most important tasks of NR is to assist LIGO, VIRGO, and other interferometric observatories, in detecting gravitational radiation and extracting the  physical parameters of the sources. Given the demanding resources required to generate these black-hole-binary simulations, and the sheer volume of the seven-dimensional space of intrinsic parameters of black-hole binaries, we need to develop techniques to model arbitrary binary configuration based on numerical simulations in a carefully chosen sample of the parameters space, in combination with post-Newtonian and perturbative calculations. One of the most promising of these approaches involves determining the region of common validity of the numerical simulations and post-Newtonian expansions, with the goal of modeling the full waveform using post-Newtonian waveforms for the initial inspiral and numerical waveforms for the late-inspiral and merger. This method was pioneered with the use of the Lazarus waveforms~\\cite{Baker:2006ha} and has readily been pursued after the breakthroughs in NR.  Comparisons of numerical simulations with post-Newtonian ones have several benefits aside from the theoretical verification of PN.  From a practical point of view, one can try to parametrize deviations of the current 3.5PN expansions to fit the numerical results~\\cite{Pan:2007nw, Buonanno:2007pf, Damour:2007vq, Damour:2008te, Boyle:2008ge}, or directly propose a phenomenological description~\\cite{Ajith:2007kx}, and thus make predictions in regions of the parameter space still not explored by numerical simulations. Another important application, from the theoretical point of view, is to have a calibration of the post-Newtonian error in the last stages of the binary merger.  The first results of comparisons for equal mass, non-spinning binaries are encouraging~\\cite{Buonanno:2006ui, Baker:2006kr, Hannam:2007ik, Hinder:2008kv, Gualtieri:2008ux}.  Recently this analysis was applied to equal-mass, equal-spin binaries with the spins aligned with the orbital angular momentum (and thus non-precessing)~\\cite{Hannam:2007wf, Shoemaker:2008pe, Vaishnav:2007nm}.  In this paper we compare the numerical and post-Newtonian waveforms for the challenging problem of a generic black-hole binary, i.e.\\ a binary with unequal masses and  unequal, non-aligned, and precessing spins.  The goal here is to evaluate accuracy of the current order of post-Newtonian expansions when including spins effects, as well as to develop new criteria for testing both numerical and post-Newtonian developments.  The paper is organized as follows, in Sec.~\\ref{sec:techniques} we review the numerical techniques used for the evolution of the black-hole binaries, in Sec.~\\ref{sec:FN} we present results from the numerical evolution of two similar generic black-hole binaries, and in~\\ref{sec:PN-intro} we analyze and compare different waveform modes as computed numerically and with the highest available post-Newtonian approximation. Finally in Sec.~\\ref{sec:discussion} we present our conclusions. ",
        "conclusions": "\\label{sec:discussion}  We analyzed the first long-term generic waveform produced by the merger of unequal mass, unequal spins, precessing black-hole binaries (a shorter simulation of this kind, which led to the discovery of the very large recoil configuration, was reported in~\\cite{Campanelli:2007ew}).  We demonstrated eighth-order convergence of the waveform phase and fourth-order convergence of the amplitude (consistent with the order of accuracy of the extraction routine) in the numerical results. These waveforms clearly show the effects of eccentricity and precession on the amplitude in the sub-leading $(\\ell=2,m=1)$  and $(\\ell=3,m=3)$ modes. In particular, analyzing the $(\\ell=2,m=1)$ mode provides a way of detecting precessional effects in the observed waveforms. We have also found that there are two sources of eccentricity for a generic binary. Residual eccentricity, due to a non-ideal choice of initial data parameters that tends to damp out as the binary separation decreases, and precession-induced eccentricity that grows as the orbital separation falls below $\\sim 15M$ (this increase in eccentricity at later times is apparent in the $(\\ell=3,m=3)$ mode of $h$ in both the PN and numerical waveforms).  We have compared these waveforms with the truncated 2.5 post-Newtonian waveforms, as well as the waveforms with the non-spinning 3.0 PN conservative and 3.5 PN radiative corrections.  We find a good initial agreement of waveforms for the first six cycles, with overlaps of over $97\\%$ for the $(\\ell=2, m=\\pm2)$ modes, $90\\%$-$98\\%$ for the $(\\ell=2,m=\\pm1)$, and over $90\\%$ for the $(\\ell=3,m=\\pm3)$ modes. This provides a natural way to match numerical waveforms to the post-Newtonian ones with a time translation (the same for all modes) motivated by the physical location of the observer (See Fig.~\\ref{fig:G3.5_re_h22}, for instance).  The agreement degrades as we approach the more dynamical region of the late merger and plunge.  The disagreement begins in a region where the numerical waveform is still very accurate. Thus it appears that the disagreement is mainly due to errors introduced by truncating the PN series. Hence the overlap should be improved significantly by including 3.0 PN and higher-order conservative and radiative corrections, including spin terms~\\cite{Steinhoff:2007mb, Porto:2007tt, Porto:2008tb, Porto:2008jj, Steinhoff:2008zr}.  In fact, our results indicate that higher-order PN corrections to the orbital motion may further increase the accuracy of the PN waveforms. Although, the PN expansion has not yet been shown to converge, we do find remarkably better agreement in $\\psi_4$ between the PN and numerical waveforms when moving from a 2.5PN EOM to a 3.5 EOM. This would appear to underscore the need for higher-order post-Newtonian calculations of both spin-orbit and spin-spin terms (especially in the EOM). Spin effects first appear at 1.5PN order, producing the spin-orbit hangup effect~\\cite{Campanelli:2006uy, Dain:2008ck}. Other spin effects, such as those due to precession, generate more subtle effects in the waveforms, and require higher-order PN corrections to accurately model (while subtle, these effects are also responsible for the very large kicks seen in spinning binaries with the spins oriented in the orbital plane). Our results seem to indicate that calculating these higher-order correction may prove to be invaluable for generating waveform templates from generic black-binary configurations."
    },
    "0808/0808.4130_arXiv.txt": {
        "abstract": "We report a multi-week sequence of B-band photometric measurements of the dwarf planet Eris using the {\\it Swift} satellite.  The use of an observatory in low-Earth orbit provides better temporal sampling than is available with a ground-based telescope.  We find no compelling evidence for an unusually slow rotation period of multiple days, as has been suggested previously.  A $\\sim$1.08 day rotation period is marginally detected at a modest level of statistical confidence ($\\sim$97\\%).  Analysis of the combination of the $Swift$ data with the ground-based B-band measurements of \\citet{2007AJ....133...26R} returns the same period ($\\sim$1.08~day) at a slightly higher statistical confidence ($\\sim$99\\%). ",
        "introduction": "The recently discovered Kuiper belt object Eris is 1.05$\\pm$0.4\\% the size of Pluto \\citep{2006ApJ...643L..61B}.  The near-infrared spectrum of Eris is dominated by methane \\citep{2005ApJ...635L..97B}, suggesting that its surface, like Pluto, is covered in significant deposits of methane frost.  The surface of Pluto is variegated, with regions of low and high albedo  \\citep{cruikshank_plutobook}. The heterogeneity of Pluto's surface is revealed in its light curve, which has a large amplitude of 0.33 magnitude \\citep{1997Icar..125..233B}.  The V band geometric albedo of Eris (85$\\pm$7\\%; \\citet{2006ApJ...643L..61B}) is approximately equal to the albedo of the brightest patches on Pluto's surface \\citep{1997AJ....113..827S}.  This led to the suggestion that the surface of Eris is likely homogeneous and of a composition similar to the brightest patches of Pluto \\citep{2006ApJ...643L..61B}.  Pluto's rotation period (6.4 days) is set by tidal interactions with its large moon Charon.  Dysnomia, the one known moon of Eris, is far too small to have significantly altered the rotation rate of Eris. Several previous observers have not conclusively identified the rotation period of Eris \\citep{2006A&A...460L..39C, 2007AdSpR..40..238L, 2007AJ....134..787S,2007AJ....133...26R,2008A&A...479..877D}. The goal of the observations reported here was to measure the rotation period of Eris. ",
        "conclusions": "The Lomb-Scargle periodogram \\citep{numrecipes} for these data is shown in Fig.~\\ref{swift_periodogram}A.  The significance levels are calculated following the technique suggested by \\citet{numrecipes}, which is to shuffle the data, randomly reassigning observed magnitudes to observation times, and recalculate the periodogram.  This operation is performed repeatedly and the peak power in each shuffled periodogram is recorded.  The significance levels are determined from $1-f$, where $f$ is the fraction of the samples for which the highest peak is greater than the power level in question.  Two peaks (at $\\sim$1.1 and $\\sim$15 days) are at suggestive levels of significance.  The data are shown phased to these periods in Fig.~\\ref{swift_periodogram}B and \\ref{swift_periodogram}C.  Periodograms are  powerful tools, but have serious limitations.  These data and the $\\sim$15 day peak are a case study in some of the hazards of  periodograms if incorrectly interpreted. One should be extremely wary of strong peaks at periods longer than $\\sim$1/3-1/2 of the total time period covered. The phased data of Fig.~\\ref{swift_periodogram}B show significant gaps in data coverage, which is an additional warning sign in periodogram analysis that results may not be trustworthy. From Fig.~\\ref{swift_periodogram}A it would be very tempting to conclude that we detected a $\\sim$15 day period in Eris, but the data incompletely span only approximately twice this length of time. Examination of the statistics of the variability of the daily mean does not support the detection of a 15 day period.  Although the human eye is drawn to the three consecutive daily mean magnitudes in Fig.~\\ref{swift_fluxes} that lie $>1\\sigma$ fainter than the overall mean of the dataset (JD 2454100-2454102), these appear to be a statistical fluke. A simple Monte Carlo simulation reveals that in a Gaussian noise dataset of 23 points (the number of daily means in Fig.~\\ref{swift_fluxes}) about 36\\% of the time there will be three consecutive data points that all sit 1-3$\\sigma$ above the mean or all sit 1-3$\\sigma$ below the mean. Due to these issues we strongly discount the significance of the peak at $\\sim$15 days.  More interesting is the peak at $\\sim$1.1 days, which has a significance level of 97$\\%$.  There are no obvious issues in the phased data of Fig.~\\ref{swift_periodogram}C. To probe the validity of this possible $\\sim$1.1 day period we ran a variety of tests on our data. We split the dataset in half and found a similar result from each half, although with the expected lower confidence level due to fewer data.  We ran a sequence of tests in which we randomly selected half of the data points to analyze and again found similar results.  We searched the data for any possible correlations, e.g.\\ flux of Eris with position on the detector, but identified none that could explain the signal seen in Fig.~\\ref{swift_periodogram}C or that could not be ruled out by examination of the ensemble of comparison stars.  We examined the Lomb-Scargle periodogram for each of the 26 comparison stars.  As expected, given that we had eliminated potential comparison stars with obvious periodicities, none of the 26 comparison stars displayed peaks in the periodogram above a significance level of 90\\%.  To further test the validity of the possible $\\sim$1.1 day period we combined the $Swift$ dataset with the B-band measurements of \\citet{2007AJ....133...26R}, which is the one other published dataset with a significant number of high-quality B-band measurements. After scaling the $Swift$ measurements to the reduced magnitude system of \\citet{2007AJ....133...26R} the periodogram of the combined dataset is shown in Fig.~\\ref{combinedfig}A.  The peak of the periodogram remains at approximately the same period (1.08~days) as with the $Swift$ dataset alone, although the significance level of the peak increases slightly to nearly 99\\%.  The combined data phased to this 1.08~day period are shown in Fig.~\\ref{combinedfig}B.  The results from the combined dataset look very similar to the results from the $Swift$ dataset alone.  We can find no reason to discount the validity of the $\\sim$1.1 day period. Using the 50$\\%$ confidence levels of the periodogram as a guide for estimating the uncertainty in the period, we report that Eris appears to be rotating once every 1.08$\\pm$0.02 days with a peak-to-valley amplitude of nearly 0.1 mag.  At the modest level of confidence available from these data the periodic signal appears less like a sinusoid and more like what would be expected from a large dark patch on the partially hidden hemisphere of an otherwise homogeneous body.  At rotational phases where the patch is hidden from the observer the light curve is constant.  A dip in the light curve is observed during the rotational phases that the dark patch is visible to the observer. Additional observations will be required to confirm this result and more precisely determine the rotation period of Eris.  Of the previously published datasets the long-term sequence of \\citet{2007AJ....133...26R} and the several nights of precision photometry of \\citet{2007AJ....134..787S} are the most likely to have been able to identify the apparent periodicity seen in these {\\it Swift} data.  \\citet{2007AJ....134..787S} did precision $R$-band photometry of Eris over several hours on each of three contiguous nights in October 2005 and three contiguous nights in December 2005.  A careful evaluation of the approximately daily sampling of their data reveals that a periodic signal such as that suggested by the {\\it Swift} data in Fig.~\\ref{swift_periodogram}C could have been missed.  The signal in Fig.~\\ref{swift_periodogram}C appears nearly constant over $\\sim$50\\% of the rotation period, while phased to a period of 1.08 days the \\citet{2007AJ....134..787S} data cover less than half of this rotation period.  The magnitude of variation in the light curve of Eris is likely a function of wavelength.  Most of the surface of Eris must be uniformly bright to achieve the high $V$-band geometric albedo \\citep[86\\%$\\pm$7\\%;][]{2006ApJ...643L..61B}.  Any darker patches are likely to be photochemically processed hydrocarbons, which will have a red color.  At longer wavelengths (e.g.\\ $R$-band and $I$-band) the albedo of these red patches will be closer to the albedo of the uniformly bright dominant surface material.  At shorter wavelengths (e.g.\\ the $B$-band used here by $Swift$) there will be greater contrast between these red patches and the rest of the surface of Eris.  This is directly analogous to what is seen on Pluto, where the light curve at $B$-band has an amplitude of $\\sim$0.1~mag greater than the amplitude at $R$-band \\citep{2003Icar..162..171B}.  Thus, the light curve of Eris is likely to be less pronounced at longer wavelengths (e.g.\\ R-band) than the shorter wavelength B-band observed with $Swift$.  \\citet{2007AJ....133...26R} reported observations in several filters, taken once per night over many months.  The analysis of the \\citet{2007AJ....133...26R} data is thus dependent upon the simultaneous fitting of color terms and phase function coefficients. We experimented with inserting fake signals into the \\citet{2007AJ....133...26R} data with a range of periods and amplitudes consistent with the periodicity detected in the {\\it Swift} data.  We then used the same analysis tools as before to search for periodicity.  In many cases the inserted periodicity is recovered, however whether the fake signal is detected and the period at which it is detected is strongly dependent on the exact amplitude and period of the fake signal.  This is primarily because the test periods ($\\sim$1.08 days) are near the sampling interval ($\\sim$1 day) of the \\citet{2007AJ....133...26R} data.  In some example cases varying the inserted signal's period by only 0.005 days made the difference between a strong detection and a non-detection.  However, as described above, the combination of the B-band measurements of \\citet{2007AJ....133...26R} with the $Swift$ measurements reported here improves the significance of the retrieved periodicity somewhat from that of the $Swift$ data alone.   A coherent picture of Eris is emerging.  The surface is primarily covered in bright methane frost, much like the brightest patches of Pluto's surface.  However, our results suggest that the surface of Eris is not perfectly homogeneous. Under thermal equilibrium the vapor pressure of methane on Eris is negligible at its current near-perihelion distance of 97.5 AU.  If the entire surface of Eris is uniformly covered in very high albedo methane frost, even at aphelion (38.2 AU) the equilibrium surface temperature of Eris is not warm enough to generate significant methane evaporation from the surface. This presents a problem as methane frost in the outer solar system is expected to redden and darken due to photochemical processing, but Eris appears bright and shows no sign of redness.  This suggests the methane frost deposits on its surface must be fresh and a replenishment mechanism is required.  Our detection of variability is consistent with regional darker areas on the surface.  At aphelion the widespread high albedo regions will not warm enough to sublimate methane into the atmosphere, however the small darker areas will be heated by the increased insolation to kickoff feedback effects that lead to dramatic global surface change, generating a temporary atmosphere and replenishing the methane surface each Eris year."
    },
    "0808/0808.1465_arXiv.txt": {
        "abstract": "The Large Volume Detector (LVD) in the Gran Sasso Laboratory of INFN is an observatory mainly devoted to search for neutrinos from core collapse supernovae. It consists of 1000 tons of liquid scintillator divided in 840 stainless steel tanks 1.5m$^3$ each. In this letter  we present the possibility for LVD to work both as a passive shield and moderator for the low energy gamma and neutron background and as an active veto for muons and higher energy neutrons.\\\\ An inner region inside the LVD structure (\"LVD Core Facility\") can be identified, with a volume of about 30m$^3$, with the neutron background typical of an underground laboratory placed at a much deeper site. This region can be realized with a negligible impact on the LVD operation and sensitive mass. The LVD Core Facility could be effectively exploited by a compact experiment for the search of rare events, such as double beta decay or dark matter. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.1186_arXiv.txt": {
        "abstract": "{We present $^{12}$CO(1-0) and $^{12}$CO(2-1) observations of the SA(rs)bc Seyfert 2 galaxy NGC\\,3147, obtained with the IRAM interferometer at 1\\farcs9\\,$\\times$1\\farcs6 and 1\\farcs6\\,$\\times$1\\farcs4 resolutions, respectively. We have also observed the central region of NGC\\,3147 with the IRAM 30\\,m telescope (at resolutions of 22\\arcsec\\ and 12\\arcsec\\ for $^{12}$CO(1-0) and $^{12}$CO(2-1), respectively), in order to obtain complete sampling at low spatial frequencies. These observations have been made in the context of the NUclei of GAlaxies (NUGA) project, aimed at the study of the different mechanisms for gas fueling of active galactic nuclei (AGN). A central peak seen mainly in $^{12}$CO(2-1) and a ring-like structure at $r \\simeq 10$\\arcsec\\,$\\sim$2\\,kpc dominate the $^{12}$CO maps. In $^{12}$CO(1-0) an outer spiral at $r \\simeq 20$\\arcsec\\,$\\sim$4\\,kpc is also detected, not visible in $^{12}$CO(2-1) emission because it falls outside the field-of-view of the primary beam. The average $I_{21}/I_{10}$ line ratio is $\\sim0.7$ in temperature units over the region mapped in both lines, consistent with the optically thick emission expected in the nuclei of spiral galaxies.  The kinematics of the molecular structures are quite regular, although there is evidence for local non-circular or streaming motions. We show that the molecular gas distribution is similar but not exactly identical to those of star formation tracers, i.e., infrared ({\\it Spitzer}) and ultraviolet ({\\it GALEX}) emission. This agreement is consistent with a scenario of steady, ongoing formation of stars from the molecular clouds at a rate of $\\sim 1\\,M_\\odot\\,{\\rm yr}^{-1}$ within the innermost 4\\,kpc in radius.\\\\ Using a near-infrared (NIR) image obtained with adaptive optics at the Canada-France-Hawaii Telescope (CFHT), we identify a weak bar in NGC\\,3147, which is classified as non-barred galaxy in the optical.  We then compute the gravity torques exerted by this stellar bar on the gas.  The torque is obtained first at each point in the map, and then azimuthally averaged with a weighting determined by the gas surface density traced by the CO emission. We find that the gas inside the inner CO ring is subject to a net negative torque and loses angular momentum.  This is expected for gas at the ultra-harmonic resonance (UHR), just inside the corotation resonance of the stellar bar. In contrast, the gas outside corotation, in the spiral arms comprising the outer spiral structure, suffers positive torques and is driven outwards. We conclude that some molecular gas is presently flowing into the central region, since we find negative torques down to the resolution limit of our images. ",
        "introduction": "The NUclei of GAlaxies (NUGA) project \\citep[][]{garcia03} is an IRAM Plateau de Bure Interferometer (PdBI) survey of nearby active galaxies to map the distribution and dynamics of molecular gas in the inner 1\\,kpc at high spatial resolution ($\\sim\\,0\\,\\farcs5 - 1^{\\prime\\prime}$, corresponding to $\\sim50-100\\,{\\rm pc}$), and to study the mechanisms for gas fueling of low-luminosity active galactic nuclei (AGNs).  Most galaxies possess central supermassive black holes (SMBHs), and gas accretion onto these black holes is the phenomenon usually invoked to explain nuclear activity in galaxies.  However, even if most galaxies host black holes, the existence of nuclear activity is far from universal. It is not clear whether the main limiting factor is the global gas mass available for fueling the AGN or the mechanisms for efficiently removing the angular momentum of the gas.  The main and non-trivial problem linked to the fueling of AGN is the removal of angular momentum from the disk gas \\citep[e.g.,][and references therein]{jogee06}, a process which can be accomplished through non-asymmetric perturbations. These can be perturbations of external origin, such as galaxy collisions, mergers, and mass accretion \\citep{heckman86}, or of internal origin due to density waves, such as spirals or bars, and their gravity torques \\citep[e.g.][]{sakamoto,combes01}. In addition to primary bars, fueling processes can be associated with more localized phenomena, such as nested nuclear bars \\citep[e.g.][]{friedli}, lopsidedness or $m = 1$ perturbations \\citep[e.g.][]{shu,garcia00}, or warped nuclear disks \\citep[e.g.][]{schinnerer00a,schinnerer00b}.  Since molecular gas is the predominant phase of the interstellar medium (ISM) in the inner kiloparsec of spiral galaxies, the study of its morphology and dynamics represents an optimal possibility for investigating AGN fueling mechanisms and their link with circumnuclear star formation.  In order to do this, high-resolution maps of molecular gas are needed.  Previous CO single dish \\citep[][]{heckman89,young,braine93,casoli96,vila98} and interferometric \\citep[][]{sakamoto,regan01,helfer03} surveys have mapped the gas in galaxies with relatively low spatial resolution ($4^{\\prime\\prime} -7^{\\prime\\prime}$). Moreover, the majority of these surveys only included small numbers of AGN in their samples. This paper, dedicated to the galaxy NGC\\,3147, is the tenth of a series where results obtained for the galaxies in the NUGA sample are described on a case-by-case basis.  \\begin{table} \\caption[]{Fundamental parameters for NGC\\,3147.} \\begin{center} \\begin{tabular}{lll} \\hline \\hline Parameter  & Value & References$^{\\mathrm{a}}$ \\\\ \\hline RA (J2000) & 10$^h$16$^m$53.6$^s$ & (1) \\\\ DEC (J2000) & 73$^{\\circ}$24$^{\\prime}$03$^{\\prime\\prime}$ & (1) \\\\ $V_{\\rm hel}$ & 2813 km\\,s$^{-1}$ & (1) \\\\ RC3 Type & SA(rs)bc & (1) \\\\ T Hubble Type & 3.88 & (2) \\\\ Inclination & 29.5$^{\\circ}$ & (2) \\\\ Position Angle & 150$^{\\circ}$ & (2) \\\\ Distance & 40.9\\,Mpc ($1^{\\prime\\prime} = 198\\,{\\rm pc}$) & (1) \\\\ $L_{B}$ & $4.3 \\times 10^{10}\\,L_{\\odot}$ & (3) \\\\ $M_{\\rm H\\,I}$ & $8.3 \\times 10^{8}\\,M_{\\odot}$ & (3) \\\\ $M_{\\rm dust}$(60 and 100\\,$\\mu$m)& $1.7 \\times 10^{6}\\,M_{\\odot}$ & (3) \\\\ $L_{\\rm FIR}$ & $3.4 \\times 10^{10}\\,L_{\\odot}$ & (4) \\\\ \\hline \\hline \\end{tabular} \\label{table1} \\end{center} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] (1) NASA/IPAC Extragalactic Database (NED); (2) Lyon Extragalactic Database (LEDA); (3) \\citet{bettoni}; (4) {\\it IRAS} Catalog. \\end{list} \\end{table}  NGC\\,3147 ($D = 40.9\\,{\\rm Mpc}$, $1^{\\prime\\prime} = 198\\,{\\rm pc}$, $H_0 = 75\\,{\\rm km\\,s^{-1}\\,Mpc}^{-1}$) is an isolated \\citep{bettoni} Seyfert 2 galaxy \\citep{ho} of Hubble morphological type SA(rs)bc.  Both \\textit{ROSAT} \\citep{roberts} and \\textit{Chandra} \\citep{terashima03} observations have shown that this galaxy possesses a pointlike nuclear X-ray source, which is presumed to be the location of the AGN. The 0.3--10\\,keV \\textit{Chandra} image reveals a bright, compact source surrounded by very faint, soft, and diffuse emission. The 2--10\\,keV core is clearly detected and confined to a region of $2^{\\prime\\prime}$ in radius. Interferometric observations of the continuum emission at cm wavelengths with MERLIN and the VLBA have also shown a pointlike nonthermal continuum source at the position of the nucleus of NGC\\,3147, coincident with the X-ray source \\citep[][]{ulvestad,krips06,krips07a}. Very recent optical and X-ray observations of the nuclear region of this galaxy suggest that NGC\\,3147 is the first ``true'' Seyfert 2 in the sense that it intrinsically lacks a broad-line region \\citep{bianchi08}. Table \\ref{table1} reports the fundamental characteristics of NGC\\,3147.  The structure of this paper is as follows.  In Sect. 2, we describe our new observations of NGC\\,3147 and the literature data with which we compare them. In Sects. 3 and 4, we present the observational results, both single dish and interferometric, describing the principal properties of the molecular gas including its morphology, its kinematics, and its excitation. A comparison between the CO observations and those obtained at other wavelengths is given in Sect. 5.  In Sect. 6, we describe the computation of the gravity torques exerted on the gas by a weak stellar bar that we have identified using a near-infrared (NIR) image. In Sect. 7, we discuss our principal observational and theoretical results, which are summarized in Sect. 8.  \\begin{figure*} \\centering \\includegraphics[width=0.9\\textwidth,angle=-90]{9545fig1n.ps} \\caption{$^{12}$CO(1-0) velocity channel maps observed with the IRAM PdBI+30\\,m in the nucleus of NGC\\,3147, with a spatial resolution of 1\\farcs9 $\\times$ 1\\farcs6 (HPBW). The center of the maps, given in Table \\ref{table1}, is $\\alpha_{2000}$ = 10$^h$16$^m$53.6$^s$, $\\delta_{2000}$ = 73$^{\\circ}$24$^{\\prime}$03$^{\\prime\\prime}$.  Velocity channels range from $\\Delta V = -230\\,{\\rm km\\,s^{-1}}$ to $+220\\,{\\rm km\\,s^{-1}}$ in steps of $10\\,{\\rm km\\,s^{-1}}$ relative to $V_{\\rm hel} = 2813\\,{\\rm km\\,s^{-1}}$. The contours begin at $5\\,{\\rm mJy\\,beam^{-1}}$, their spacing is $5\\,{\\rm mJy\\,beam^{-1}}$, and their maximum is $50\\,{\\rm mJy\\,beam^{-1}}$.} \\label{channel} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[width=0.9\\textwidth,angle=-90]{9545fig2n.ps} \\caption{Same as Fig. \\ref{channel} but for the $^{12}$CO(2-1) line, with a spatial resolution of 1\\farcs6 $\\times$ 1\\farcs4. The velocity range is the same as for Fig. \\ref{channel}. The contours begin at $10\\,{\\rm mJy\\,beam^{-1}}$, their spacing is $5\\,{\\rm mJy\\,beam^{-1}}$, and their maximum is $45\\,{\\rm mJy\\,beam^{-1}}$.} \\label{channel21} \\end{figure*} ",
        "conclusions": "\\label{sec:Conclusions} The molecular gas in the Seyfert 2 galaxy NGC~3147 has been mapped with high resolution (1\\farcs9 $\\times$ 1\\farcs6 Gaussian beam for the $^{12}$CO(1-0) line) inside a radius of 25$^{\\prime\\prime}$ ($\\sim$5 kpc). The CO emission shows a central peak (mainly in $^{12}$CO(2-1)) and a ring at distance of 2\\,kpc. In $^{12}$CO(1-0) also an outer spiral at distance of 4\\,kpc is detected. The observed CO has a mean line intensity ratio $I_{21}/ I_{10}$  $\\sim$0.7, consistent with the optically thick emission expected in the nuclei of spiral galaxies, and broadly regular kinematics with some evidence for local non-circular motions.  Comparing the molecular gas distribution with tracers of star formation, we find that central emission and the inner CO ring coincide well with the FUV (\\textit{GALEX}) and 8\\,$\\mu$m (\\textit{Spitzer}) emission, while the outer CO spiral structure/ring is instead located within an interarm interval detected by \\textit{GALEX} and \\textit{Spitzer}. We interpret this partial disagreement as due to the transient nature of the outer CO ring. It would be a structure in re-condensation that traces a future stellar generation, but is decoupled from those that have heated the dust observable at 8\\,$\\mu$m or from whose emission is detected in the FUV.   With a NIR image obtained with the Canada-France-Hawaii Telescope, we have identified the presence of a weak bar in NGC~3147, a galaxy classified as non-barred in the optical. This stellar bar acting on the gas produces gravity  torques. A portion of the gas present in the inner $^{12}$CO(1-0) ring is trailing the bar, while part of the gas leads it. The gravity torques are negative inside a radius of $\\sim$2\\,kpc, while they become zero or positive outside. In the outer $^{12}$CO(1-0) spiral the torques almost cancel out on average, but are positive. In the more clumpy inner $^{12}$CO(2-1) ring, the predominant torques are negative. We interpret these results by identifying the resonances with the bar. The inner CO ring would correspond to the ultra-harmonic resonance (r = 2 kpc) just inside corotation, and the outer CO spiral to the spiral structure emerging from the bar, beyond its corotation, which we estimate is located at $\\sim$3\\,kpc. In the inner CO ring, the gas is inside corotation due to the negative torques there; it loses angular momentum and is flowing inwards.  NGC~3147 is not the first case in the NUGA sample where inflowing gas has been found: NGC 2782 \\citep{hunt}, and NGC 6574 \\citep{lindt} also show this feature.   However, only NGC 2782 has a strong and significant inflow, due to a nuclear bar embedded in the primary one and to gas aligned with the nuclear bar. The amount of inflowing gas in NGC~3147, quantified from the gravity torques computation, is modest but could still be sufficient to feed the Seyfert 2 nucleus at a rate of about 0.5 M$_\\odot$\\,yr$^{-1}$. Although higher spatial resolution is needed to determine whether the gas is actually feeding the AGN, we are observing the gas inflow to the $\\sim$100--200 pc central region."
    },
    "0808/0808.2129_arXiv.txt": {
        "abstract": "We present new optical and infrared photometry for a statistically complete sample of seven 1.1 mm selected sources with accurate coordinates.  We determine photometric redshifts for four of the seven sources of 4.64, 4.54, 1.49 and 0.18.  Of the other three sources two are undetected at optical wavelengths down to the limits of very deep Subaru and Canada-France-Hawaii Telescope images ($\\sim$27 mag AB, i band) and the remaining source is obscured by a bright nearby galaxy. The sources with the highest redshifts are at higher redshifts than all but one of the $\\sim$200 sources taken from the largest recent 850 $\\mu$m surveys, which may indicate that 1.1 mm surveys are more efficient at finding sources at very high redshifts than 850 $\\mu$m surveys.  We investigate the evolution of the number density with redshift of our sample using a banded $V_{e}/V_{a}$ analysis and find no evidence for a redshift cutoff, although the number of sources is very small. We also perform the same analysis on a statistically complete sample of 38 galaxies selected at 850$\\mu$m from the GOODS-N field and find evidence for a drop-off in the number density beyond $z\\sim1$ and 2, confirming the earlier conclusion of Wall, Pope \\& Scott (2008).  We also find evidence for the existence of two differently evolving sub-populations separated in luminosity, with the drop-off in density for the low-luminosity sources occurring at a lower redshift. ",
        "introduction": "Submillimetre (submm) galaxies (SMGs), first detected at 850 $\\mu$m with the Submillimetre Common User Bolometer Array (SCUBA) \\citep{holland:1999}, are a significant population of high redshift star forming galaxies \\citep[e.g.,][]{hughes:1998,blain:2002}.  They are believed to be dust enshrouded galaxies undergoing prodigious levels of star formation \\citep[e.g.,][]{hughes:1998,eales:1999} in which the optical/UV radiation emitted by the stars is absorbed by the dust and then re-emitted in the submm.  Star formation rates in excess of 1000 M$_{\\odot}$yr$^{-1}$ have been inferred \\citep{scott:2002}, much higher than locally.  The galaxies in these samples have been found to account for up to one tenth of the total far-infrared/submm energetic background \\citep[e.g.,][]{dye:2007} and many authors have argued that these galaxies are the progenitors for the elliptical galaxies we see in the local Universe \\citep{eales:1999,scott:2002,dunne:2003}.  Thus understanding the nature of these sources is of great importance for the understanding of galaxy formation and evolution as a whole.  Observations of SMGs at $\\sim$1 mm benefit from a negative K-correction out to high redshifts due the shape of their spectral energy distribution (SED).  As the redshift of an SMG increases, the peak of its rest-frame SED moves toward the observed waveband, offsetting the dimming caused by the increasing luminosity distance. This fact accounts for the surprising ability of SCUBA to find large numbers of high-redshift galaxies.  The large amount of dust responsible for the strong submm emission gives rise to high levels of attenuation in the optical.  This in conjunction with the poor angular resolution of single dish submm facilities makes the cross identification of SMGs at different wavelengths difficult.  Moreover, even when an optical counterpart can be identified, the high levels of dust attenuation makes the determination of a spectroscopic redshift difficult.  As such we are currently unable to determine spectroscopic redshifts for the majority of SMGs.  The strong correlation between dust emission and radio emission which appears to hold true in both the low-redshift and high-redshift universe \\citep{vlahakis:2007} has been useful for both identifying the counterparts and estimating redshifts.  Due to the low surface density of radio sources on the sky, the probability of the radio counterpart being coincidental with the submm source by chance is small.  Due to the high positional accuracy of radio observations, it is then possible to identify the optical counterpart and retrieve a spectroscopic redshift.  It is also possible to estimate the redshift using the ratio of radio to submm flux \\citep[e.g.,][]{hughes:1998, carilli:1999, carilli:2000, smail:2000}.  \\citet{chapman:2005}, using the Low Resolution Imaging Spectrograph (LRIS) \\citep{oke:1995} on the Keck I telescope, managed to obtain spectroscopic redshifts for a total of 73 radio-identified SMGs with a median 850 $\\mu$m flux of 5.3 mJy.  The galaxies in this sample were found to lie at a median redshift of $z=2.2$ out to a maximum value of $z_{max}=3.6$.  However, the K-correction which allows us to detect high-redshift SMGs does not similarly benefit their radio fluxes and so radio identified SMGs are subjected to a radio selection effect which limits redshifts to $z\\simless3$.  \\citet{pope:2006}, produced the first complete (i.e. not requiring radio IDs) sample of 850 $\\mu$m selected SMGs that has close to 100\\% redshifts.  The sample consists of 35 galaxies, 21 with secure optical counterparts and 12 with tentative optical counterparts, and its completeness means that unlike previous surveys it is not biased towards low-z sources.  The median redshift determined for this sample is $z\\sim 2.2$.  Using this sample, \\citet{wall:2008} examined the epoch dependency of the number density of SMGs.  They found an apparent redshift cutoff at $z>3$ with further evidence for two separately evolving populations, divided by luminosity.  However this result was based on calculations using a single model galaxy SED. Since the predicted relationship between submm flux-density and redshift depends strongly on the assumed SED, one of the aims of this paper is to re-examine their conclusion using a range of empirical SEDs.  There have been a number of explanations for the lack of high-redshift SMGs.  Given that dust is thought to form in the atmospheres of highly evolved stars, it is possible that at high redshifts simply not enough time has passed for dust to form \\citep{morgan:2003}.  Observations of high-redshift quasars have however detected high levels of dust \\citep[e.g.,][]{priddey:2001}, suggesting that this is not the explanation.  Another possible explanation is that there are fewer large star-forming galaxies at high redshifts.  \\citet{eales:2003} presented evidence that SMGs typically have low values for the ratio of the 850 $\\mu$m to 1200 $\\mu$m fluxes compared to that expected from a low redshift galaxy.  One possible explanation is that these sources are at very high redshifts.  If this is true, then observations at 1.1 mm would be better at detecting SMGs at the highest redshifts than observations at 850 $\\mu$m.  A new complete sample of 1.1 mm selected SMGs located in the COSMOS field \\citep{scoville:2007} has been compiled by \\citet{younger:2007}.  The sources were selected initially at 1.1 mm with the AzTEC camera \\citep{scott:2008, wilson:2008} on the JCMT.  The resultant catalogue consists of 44 sources with S/N$>3.5\\sigma$, 10 of which are robust with S/N$>5\\sigma$.  Follow up observations by \\citet{younger:2007} were then made with the Submillimetre Array (SMA) at 890 $\\mu$m for the 7 highest significance AzTEC sources, allowing their positions to be determined with an accuracy of $\\sim 0.2$''.  The COSMOS field offers a wealth of data over a great number of wavebands including the optical and infrared.  Thus the high positional accuracy allows for the identification of optical counterparts and hence the determination of photometric redshifts.  Of the seven AzTEC sources imaged with the SMA six have IRAC counterparts, and one source is obscured by a nearby bright galaxy.  Using deep Hubble Space Telescope (HST) imaging acquired with the Advanced Camera for Surveys (ACS), \\citet{koekemoer:2007} found optical counterpart candidates for only three of these sources.  The main aim of this paper is to carry out a deeper search for the optical counterparts for the AzTEC sources.  We give photometry from deep Subaru and Canada-France-Hawaii-Telescope (CFHT) imaging and find one new possible optical ID.  We estimate photometric redshifts for the AzTEC sources using the {\\scshape HyperZ} photometric redshift package \\citep{bolzonella:2000}.  Throughout this work we employ a concordance cosmological model with $\\Omega_{total}=1$, $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_{0} = 75$ kms$^{-1}$Mpc$^{-1}$.  All magnitudes quoted are AB. ",
        "conclusions": "We give new Subaru, CFHT and IRAC photometry for a number of sources in the AzTEC / COSMOS survey with accurate coordinates from SMA imaging.  We have estimated photometric redshifts for four of the seven galaxies in the sample.  We find a median redshift of $z_{mean}\\sim2.57$ and a maximum of $z_{max}=4.50$.  Of the sources in the combined 850 $\\mu$m surveys presented in \\citet{chapman:2005}, \\citet{pope:2006}, \\citet{dye:2008} and \\citet{clements:2008}, consisting of $\\sim$200 sources, only one is at a redshift greater than our two highest redshift sources.  This in addition to the fact that we are unable to detect two of our sources in the optical bands down to very faint magnitudes may indicate that 1.1~mm surveys are more efficient at detecting very high-redshift sources than 850 $\\mu$m surveys.  Re-investigating the space density evolution of a sample of 38 GOODS-N sources \\citep{pope:2006,wall:2008} with more realistic SEDs we find a redshift cutoff at $z\\sim1$ if we assume a 'hot' SED and marginal evidence for a cutoff at $z\\sim2$ if we assume a 'cold' SED (in reasonable agreement with Wall et al.).  Similar to \\citet{wall:2008} we also found evidence for two differently evolving sub-populations of SMGs, separated in luminosity, with high luminosity sources showing a less negative evolution.  We performed a similar test on the AzTEC sources but were unable to draw any reliable conclusions as the sample is too small.  The GOODS-N sample is also relatively small, and therefore any evidence for redshift cutoffs and differently evolving sub-populations must be treated with caution.  In order to harden our conclusions in general we require larger surveys with accurate redshifts.  We would also need surveys taken over larger areas of sky in order to take into account the effects of cosmic variance.  Future, larger surveys (e.g. with Herschel, SCUBA2) therefore will enable us to more robustly determine the nature of the number density evolution of SMGs in the Universe.  \\begin{flushleft} {\\bf Acknowledgements} \\end{flushleft} G. Raymond, S. Eales and S. Dye acknowledge support from the Science and Technologies Facilities Council.  Based on observations obtained with MegaPrime/MegaCam, a joint project of CFHT and CEA/DAPNIA, at the Canada-France-Hawaii Telescope (CFHT) which is operated by the National Research Council (NRC) of Canada, the Institut National des Science de l'Univers of the Centre National de la Recherche Scientifique (CNRS) of France, and the University of Hawaii. This work is based in part on data products produced at TERAPIX and the Canadian Astronomy Data Centre as part of the Canada-France-Hawaii Telescope Legacy Survey, a collaborative project of NRC and CNRS."
    },
    "0808/0808.2403_arXiv.txt": {
        "abstract": "{Cosmic microwave background (CMB) polarimetry has the potential to provide revolutionary advances in cosmology.  Future experiments to detect the very weak B mode signal in CMB polarization maps will require unprecedented sensitivity and control of systematic errors.  Bolometric interferometry may provide a way to achieve these goals. In a bolometric interferometer (or other adding interferometer), phase shift sequences are applied to the inputs in order to recover the visibilities.  Noise is minimized when the phase shift sequences corresponding to all visibilities are orthogonal.  We present a systematic method for finding sequences that produce this orthogonality, approximately minimizing both the length of the time sequence and the number of discrete phase shift values required. When some baselines are geometrically equivalent, we can choose sequences that read out those baselines simultaneously, which has been shown to improve signal to noise ratio. } ",
        "introduction": "The field of observational cosmology has been advancing quickly in recent years. Observations of the cosmic microwave background (CMB) radiation have been leading the way, as evidenced by WMAP's highly successful mapping of CMB anisotropy \\citep{wmap5yrbasic}, DASI's detection of the polarized component of the CMB \\citep{dasi}, and the 2006 Nobel Prize in Physics awarded to John Mather and George Smoot. Momentum is building for experiments that characterize the CMB polarization in detail \\citep{weissreport}.  A linear polarization map can always be expressed as the sum of two component maps, denoted E and B \\citep{selzal,kks}. CMB experiments to date have detected only the ``curl-free'' E component, which is produced primarily by (scalar) density perturbations. The ``divergence-free'' B component is not produced by scalar perturbations at linear order, and is therefore a clean probe of other, smaller effects. In particular, inflationary models predict a B-mode signal produced by gravitational wave (tensor) perturbations in the early universe. These B modes promise to hold key information about the process of inflation and particle physics above the Grand Unification scale. The challenge of finding the B modes is no small task, however: the B component is expected to be at least an order of magnitude weaker than the E component (which is itself small compared to the temperature anisotropy) over all angular scales. Experiments to search for B modes will require unprecedented sensitivity and control of systematic errors.  Bolometric interferometry is one proposed method for achieving these goals \\citep{brain1,brain2,mbi1,mbi2,tucker:70201M}. A bolometric interferometer is a marriage between highly sensitive, incoherent bolometric detectors and the phase-sensitive, systematic-error-reducing observing technique of interferometry. \\citet{hamilton} have shown that a bolometric interferometer can achieve sensitivities comparable to traditional technologies.  The question of whether bolometric interferometry is useful for CMB polarimetry will thus depend on the method's ability to control systematic errors. Systematic errors in interferometers are certainly different from those in imaging experiments \\citep{bunn}; it can be argued that interferometers are superior in this regard, although this question requires further research.  In a bolometric interferometer, the signals from a set of input feedhorns are combined with either a Butler combiner or a quasi-optical (Fizeau) combiner.  In either case, bolometers measure the total power in the combined beam -- that is, each bolometer is illuminated by signals from all of the input horns.  Since the signal in each detector is the sum of all the inputs, a bolometric interferometer is an example of an ``adding'' interferometer (as opposed to traditional radio interferometers, which are ``multiplying'' interferometers). One of the keys to making this method work is to arrange for the phase information to be encoded in the bolometer signals, so that individual pairwise visibilities can be extracted. To achieve this goal, a sequence of phase shifts can be applied to each of the input horns, in such a way that each visibility is phase-shifted in an independent fashion.  The resulting time series can be solved for the individual visibilities. The phase shift sequences should be chosen so that this inversion can be done with minimal noise.  We would like the length of the phase shift sequence to be as short as possible, to avoid error due to $1/f$ noise in the detectors. Clearly the number of phase shifts must be at least as large as the number of visibilities to be recovered.  In the most general case, an $N$-horn interferometer has $N(N-1)/2$ distinct visibilities, requiring long phase shift sequences for interferometers with many inputs.  On the other hand, if the input horns are arranged in a regular pattern, such as a square array, then many antenna pairs correspond to identical visibilities.  These can be given identical phase shifts and read out together.  This coherent treatment of equivalent (or redundant) baselines has two advantages. First, it allows for shorter phase shift sequences.  Second, by coadding equivalent signals, the signal-to-noise ratio is improved \\citep[hereinafter C08]{charlassier}.  C08 gave an excellent overview of how a bolometric interferometer works and considered the choice of phase shift sequences in detail.  For the case of a square array of horns, the paper presented a method of phase modulating the inputs that gives equivalent baselines identical phase shift sequences. In this paper we present independently-developed work on phase modulation and coherent addition of baselines that complements the methods of C08. We consider general horn arrangements as well as a regular square lattice.  In the general case, we present a method for finding the optimal phase shift sequence assuming no baselines are redundant.  In the case of a square array, we present a refinement of the method of C08.  Unlike the original method, which achieves optimal noise performance only in the limit as the number of time steps tends to infinity, our method is optimal for sequences of nearly or exactly the minimum possible length.  In Sect.~\\ref{sec:Form} we present our formalism for denoting sequences of phase shifts and consider the criterion for an optimal phase shift sequence.  Section \\ref{sec:Meth} introduces a shorthand notation and applies this to a method for constructing bases for phase sequences and selecting optimal sequences for interferometers without equivalent baselines. In Sect.~\\ref{sec:Array} we consider a square array of horns, accounting for redundant baselines. For simplicity, we consider only one linear polarization state in these sections, but in Sect.~\\ref{sec:Pol} we generalize to an array with two polarizations.  Section \\ref{sec:Conclusions} contains a brief concluding discussion, and Appendix A contains a useful mathematical result. ",
        "conclusions": "\\label{sec:Conclusions} Optimal recovery of visibilities in a bolometric interferometer depends on the proper choice of phase shift sequences. We have laid out a method for finding such sequences that lead to fully orthogonal masks for all visibilities and introduced a compact notation for describing such phase shift sequences.  In the case of an array with a regular lattice structure, equivalent baselines can be read out simultaneously, reducing the length of the required phase shift sequence and improving the signal-to-noise ratio.  This method refines that of C08.  The method applies to arrays that are based on replication of any parallelogram-shaped fundamental cell, including for example hexagonal arrays.  For the case of rectangular arrays, the method described herein is very similar to that of C08, although our method imposes strict orthogonality on the masks for distinct baselines, rather than relying on approximate orthogonality resulting from random sequences.  As a result, our method leads to optimal recovery of visibilities for Stokes' $I,U,V$, with shorter time sequences than that of C08.  The ability to shorten the sequence of phase shifts is likely to be important in instrument design, because it reduces the degree to which $1/f$ noise must be controlled.  For example, suppose that we can shift phase states at a rate of one state per 10 ms (either because of the design of the phase shifters or the bolometer time constants).  As we saw in the previous section, an $8\\times 8$ array requires $\\sim 1000$ phase shifts, which would take 10 seconds.  We therefore require the $1/f$ noise knee to be below $\\sim 0.1$ Hz. An alternative scheme involving a longer phase shift sequence would require correspondingly tighter control of the $1/f$ knee."
    },
    "0808/0808.2824_arXiv.txt": {
        "abstract": "{We use HAWK-I, the recently-commissioned near-IR imager on Yepun (VLT-UT4), to obtain wide-field, high-resolution images of the X-ray luminous galaxy cluster \\object{XMMU\\,J2235.3-2557} in the J and Ks bands, and we use these images to build a colour-magnitude diagram of cluster galaxies. Galaxies in the core of the cluster form a tight red sequence with a mean $\\mathrm{J-Ks}$ colour of 1.9 (Vega system). The intrinsic scatter in the colour of galaxies that lie on the red sequence is similar to that measured for galaxies on the red sequence of the Coma cluster.  The slope and location of the red sequence can be modelled by passively evolving the red sequence of the Coma cluster backwards in time. Using simple stellar population (SSP) models, we find that galaxies in the core of \\object{XMMU\\,J2235.3-2557} are, even at $z=1.39$, already 3\\,Gyr old, corresponding to a formation redshift of $z_{\\rm f} \\sim 4$. Outside the core, the intrinsic scatter and the fraction of galaxies actively forming stars increase substantially. Using SSP models, we find that most of these galaxies will join the red sequence within 1.5\\,Gyr. The contrast between galaxies in the cluster core and galaxies in the cluster outskirts indicates that the red sequence of \\object{XMMU\\,J2235.3-2557} is being built from the dense cluster core outwards.} ",
        "introduction": "The most notable feature in the colour-magnitude (C-M) diagram of a rich galaxy cluster is the tight red sequence of passively-evolving galaxies. First observed in nearby galaxy clusters \\citep[e.g.,][]{Vaucouleurs61,Visvanathan77}, it has also been observed in some of the most distant galaxy clusters currently known \\citep[e.g.,][]{vanDokkum01,Blakeslee03,Lidman04,Holden04,Mei06}.  The location and tightness of the red sequence are used to infer that the bulk of the stars in red sequence galaxies formed over a relatively short space of time at much higher redshifts \\citep{Bower92,Stanford98}. In \\object{RDCS~J1252-2927} at $z=1.24$, for example, \\citet{Blakeslee03} find that most of the stars in the galaxies that lie on the red sequence formed 11.9 Gyr ago.  The red sequence is tilted in the sense that brighter galaxies are also redder. The tilt is understood as a relationship between mass and metallicity \\citep{Kodama99,Gallazzi06}. Brighter, more massive galaxies are more metal rich and are hence redder. The tilt is observed at all redshifts, and there is little evidence that it evolves by more than that expected from passive evolution of an old stellar population \\citep{Stanford98,Blakeslee03}. However, there are hints that it flattens at the bright end for some $z > 1$ clusters \\citep{vanDokkum01,Demarco07} and for S0 galaxies in the cluster \\object{RDCS~J0910+5422} at $z=1.106$ \\citep{Mei06}.  When sufficiently precise measurements have been made, the sequence has a small, but non-zero, scatter in colour \\citep{Bower92,Eisenhardt07}.  For the most massive clusters, the scatter appears to be independent of redshift \\citep{Stanford98,vanDokkum01,Blakeslee03,Mei06,Homeier06}; however, in some less massive clusters \\citep{Holden04}, the scatter is larger and the red sequence is less regular. While it is generally accepted that the tilt in the colour-magnitude relation is primarily the result of a relationship between mass and metallicity, the reasons for the scatter are less clear. For early-type galaxies in the field, the scatter is caused by differences in age and metallicity, with variations in metallicity playing an increasingly important role as the mass of the galaxy increases \\citep{Gallazzi06}. Dust could also play a role.  Within the framework of hierarchical structure formation, one expects the morphology of the red sequence to evolve as one approaches the epoch of cluster formation. The major processes are thought to be: passive evolution of old stellar populations, conversion of galaxies that are in the ``blue cloud'' into red ones through the quenching of star formation, and dry mergers \\citep{Bell06,Faber07}. The time at which these processes initiate and the speed at which they run depend on environment, with the result that the red sequences of the most massive clusters form first \\citep{Tanaka05,Romeo08,Menci08}.  Recent observations of the most distant clusters have started to note changes in the morphology of the red sequence. A deficit of galaxies at the faint end of the red sequence has been noted in the large scale structure surrounding the high redshift cluster \\object{RDCS~J1252-2927} at $z=1.24$ \\citep{Tanaka07}. A similar deficit in the cluster itself was not seen, suggesting that the rate at which the red sequence develops depends on environment, as one expects in hierarchical models and as seen in some clusters at lower redshifts \\citep{Tanaka05}. The faint end deficit has been noted in clusters from the ESO distant cluster survey \\citep{deLucia04b,deLucia07}; however, no such deficit was found in the cluster sample that was analysed by \\citet{Andreon08}, perhaps indicating that the clusters analysed by Andreon were, on average, richer than those in ESO distant cluster survey. A similar richness dependence was noted in clusters from the Red Cluster Survey \\citep{Gilbank08}.  In a recent study of several proto-clusters with redshifts ranging from $z=2.2$ to $z=3.1$, \\citet{Kodama07} found that the bright end of the red sequence becomes progressively less well defined as the redshift of the proto-cluster increases. In one of the most well-studied proto-clusters, \\object{PKS~1138-262} at $z=2.16$, \\citet{Zirm08} found that the red sequence of \\object{PKS~1138-262} is considerably less well defined than the red sequence of rich clusters at lower redshifts.  Hence, between the redshift of most well studied clusters at $z \\sim 1.3$ and the most well studied proto-clusters at $z \\sim 2$, there is evidence for a dramatic change in the morphology of the red sequence. Massive X-ray luminous clusters beyond a redshift of 1.3 are rare, and currently, only 2 are known, \\object{XMMU~J2235.3-2557} \\citep{Mullis05} at $z=1.39$ and \\object{XMMXCS~J2215.9-1738} at $z=1.45$ \\citep{Stanford06}.  In this paper, we present the near-IR colour-magnitude (C-M) diagram of \\object{XMMU~J2235.3-2557}, one of the most distant X-ray luminous clusters currently known. In section 2 of the paper, we describe the near-IR observations and the methods used to process the data. In section 3, we present the C-M diagram and derive the intrinsic scatter in the colour of galaxies that lie on the red sequence. In section 4, we discuss the results. Throughout this paper, we assume a flat, $\\Lambda$-dominated universe with $\\Omega_{\\mathrm M}=0.27$ and $H_{\\mathrm 0} = 71\\,\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$. In this cosmology, 1\\arcsec\\ on the sky corresponds to 8.5\\,kpc at $z=1.39$. Unless specified otherwise, all magnitudes are Vega magnitudes and are on the 2MASS system. In a forthcoming paper, we will present an analysis of multi-wavelength observations of \\object{XMMU~J2235.3-2557}, including spectroscopy from VLT/FORS2 and optical imaging from ACS/HST. ",
        "conclusions": "We have presented near-IR observations of the X-ray luminous cluster {XMMU\\,J2235.3-2557} at $z=1.39$ and we have built a C-M diagram of objects inside and outside of the core of the cluster. While galaxies inside the cluster core form a well defined red sequence with no evidence of ongoing star formation, cluster members outside the core are much more diverse.  The colour of galaxies inside the core can be matched with SSP models that are $\\sim 3$ Gyr old, corresponding to a redshift of formation of $z_{\\rm f}\\sim 4$. These galaxies are already very old, especially when we consider that the cluster was observed at a time when the age of Universe was 4.6\\,Gyr.  Cluster members outside the core do not form a well defined red sequence. Over half these galaxies are forming stars and some of these are either considerably redder than the red sequence, perhaps indicating the presence of dust, or considerably bluer. The other half do not appear to be forming stars, but are, on average, displaced towards bluer colours, perhaps indicating that they either stopped forming stars recently or are younger than galaxies on the red sequence.  The contrast between the cluster core, which consists of a population of evolved galaxies with uniform colours, and the cluster outskirts, which consists of a population of active galaxies with diverse colours, suggests that the red sequence of this cluster is being built from the inside to the outside or, alternatively, from the dense core to the relatively sparse outskirts."
    },
    "0808/0808.0010_arXiv.txt": {
        "abstract": "We examine the UV emission from luminous early-type galaxies as a function of redshift.  We perform a stacking analysis using Galaxy Evolution Explorer (GALEX) images of galaxies in the NOAO Deep Wide Field Survey (NDWFS) Bo\\\"otes field and examine the evolution in the UV colors of the average galaxy. Our sample, selected to have minimal ongoing star formation based on the optical to mid-IR SEDs of the galaxies, includes 1843 galaxies spanning the redshift range $0.05\\leq~z\\leq0.65$.  We find evidence that the strength of the UV excess decreases, on average, with redshift, and our measurements also show moderate disagreement with previous models of the UV excess.  Our results show little evolution in the shape of the UV continuum with redshift, consistent either with the binary model for the formation of Extreme Horizontal Branch (EHB) stars or with no evolution in EHB morphology with look-back time.  However, the binary formation model predicts that the strength of the UV excess should also be relatively constant, in contradiction with our measured results. Finally, we see no significant influence of a galaxy's environment on the strength of its UV excess. ",
        "introduction": "The UV excess is defined as the presence of more UV flux than predicted for a simple, old stellar population and was first reported by \\citet{code79}.  It is sometimes also called the UV upturn, because $F_{\\lambda}$ is seen to rise shortward of ${\\rm 2500\\AA}$ \\citep{brow04}.  \\citet{dona95}, for example, found that the average early-type galaxy in the Coma cluster is more than a magnitude bluer in $m_{UV}-B$ than predicted by the population synthesis models of \\citet{bruz93}.  They explaind the observed emission by invoking residual star formation (RSF), but it is unlikely that all of the early-type galaxies in the cluster experienced a recent burst of star formation at around the same time; this indicates the need for a source of UV emission not associated with star formation. Population synthesis models have since suggested a number of potential sources for this emission, usually involving significant mass loss by stars leaving the main sequence.  The proposed sources include post-red giant stars, hot horizontal branch stars and post-AGB stars (e.g. \\citealt{bres94}). A combination of population synthesis models and high resolution spectra obtained with the Far UV Spectroscopic Explorer (FUSE; \\citealt{brow02}) and the Hopkins Ultraviolet Telescope (HUT; e.g. \\citealt{ferg93}) suggest that extreme horizontal branch (EHB) stars, also called hot subdwarfs (sdB), are the objects most likely to give rise to the UV emission, as the observed spectra closely match predictions from a population of EHB stars with various surface temperatures \\citep{brow04}.  Thus, conventional wisdom indicates that the stars giving rise to the UV emission are produced by significant mass loss from stellar envelopes on the red giant branch (RGB).  If EHB stars are formed from stars with massive winds on the RGB, then the fraction of stars that find themselves in this unusual stage of stellar evolution, and thus the strength of the UV excess, should depend on the average metallicity of the host galaxy.  Evidence for this picture can be found from several sources.  First is the direct correlation between the strength of the UV excess and the Lick ${\\rm Mg_{2}}$ spectral index, which measures a galaxy's average metallicity \\citep{burs88}. Also in keeping with such a metallicity dependence is the correlation between the UV excess and the mass (luminosity) of a galaxy suggested by \\citet{ocon99}. However, this correlation has recently become controversial. For example, \\citet{rich05} found no apparent correlation between the ${\\rm Mg}_{2}$ index and the UV excess in a sample of 172 early-type galaxies from the Sloan Digital Sky Survey (SDSS). By contrast, \\citet{dona06} found a weak but significant correlation among elliptical ($-5.5\\leq~T<-3.5$) galaxies, but they found no such correlation in lenticular ($-3.5\\leq~T<-1.5$) galaxies. They attribute this difference to the presence of residual star formation in lenticulars.  Recently, Ree et al. (2007; R07) used Galaxy Evolution Explorer (GALEX) photometry of galaxy clusters below redshift $z=0.2$ to measure the evolution in the UV excess of Brightest Cluster Galaxies (BCGs) of rich clusters with redshift.  The galaxies they measured showed no significant evolution, but by expanding their galaxy sample with objects studied using HST by Brown et al. (2000, 2003), they found that the $FUV-V$ colors of early-type galaxies in massive clusters become redder at higher redshift. They compare their expanded sample to two evolutionary models, one favoring sdB formation in metal-rich populations and the other favoring metal-poor populations, attempting to determine which model agrees better with the measured colors.  These models are developed by picking a pair of populations to bracket the $z=0$ galaxies and then passively ``evolving'' them backwards in time. They find that both models agree reasonably well with the measurements, and marginally favor the metal-poor model.  Han, Podsiadlowski \\& Lynas-Gray (2007; HPL) suggested that sdB stars might form primarily via close binary interactions rather than forming via wind-driven mass loss on the RGB.  Stars in close binary systems may eject much of their hydrogen envelope after they evolve off the main sequence via angular momentum exchange between the envelope and the binary companion.  Their model makes several specific predictions, including that the correlation of the UV excess with metallicity should be weak, as should the evolution of the strength and shape of the UV excess with redshift. The HPL model is rather appealing, as it can explain the limited strength of the metallicity correlation found in the SDSS galaxies \\citep{rich05} and the large fraction of Galactic field sdB stars found in binary systems compared to those found in Galactic globular clusters \\citep{cate07}.  The alternative hypotheses for the formation of sdB stars can be tested by examining the evolution of the UV colors of early-type galaxies with redshift.  In this work we measure the evolution of the average early type galaxy by stacking GALEX images of galaxies in the Bo\\\"otes field of the NOAO Deep Wide Field Survey (NDWFS). The galaxies are selected based on their optical to mid-IR spectral energy distribution (SEDs), following \\citet{asse07}.  We study the UV emission from elliptical galaxies out to $z=0.65$, where the number of sample galaxies in each redshift bin begins to diminish and the risk of AGN contamination increases. In \\S\\ref{secTarget} we describe our galaxy selection procedures. We describe our stacking algorithm and the analysis of the resulting images in \\S\\ref{secGalaxy}, refine our galaxy selection criteria in \\S\\ref{secRefine}, and we examine the evolution of the UV excess in \\S\\ref{secEvolution}.  Finally, in \\S\\ref{secConclusion} we consider the consequences of our measurements for models of sdB formation. ",
        "conclusions": "We have measured the evolution of the UV emission from the average luminous, early-type galaxy with redshift by performing a stacking analysis with photometrically selected galaxies from the Bo\\\"otes field. To the extent possible with the available data, we have eliminated AGN from our sample, although the stacked AGES spectra suggest that weak LINERs are ubiquitous.  LINERs show very weak UV continua compared to broad-line AGN of similar luminosity, so the contamination of our results due to nuclear activity is minimal. We find that the observed $FUV-V$ colors of our stacked galaxies are bluer than the colors of the BCGs studied by \\citet{ree07}, \\citet{brow00} and \\citet{brow03} and show a pronounced tendency to become bluer with redshift. The presence of a small excess in 8$\\mu m$ emission indicates the need to correct for residual star formation.  The necessity of such a correction, even among a sample of galaxies that has been carefully selected to have as little star formation as possible, suggests that contributions from star formation should always be considered when measuring the UV excess photometrically.  This is consistent with the results of other authors (e.g. \\citealt{kavi06}).  After correcting for star formation, we find that the intrinsic strength of the UV excess, as measured by the $FUV-V$ colors, is inconsistent with a non-evolving model.  The measured evolution shows reasonable agreement with previous studies.  However, the $FUV-V$ colors of the averaged galaxies remain modestly bluer than the individual cluster galaxies that have been studied previously. Our results agree with the R07 results in the first two redshift bins, but their preferred models are inconsistent with our higher redshift data, suggesting that one or more of the model parameters needs to be adjusted.  At least some of the evolution may be due to an increase in $\\langle L_{bol}\\rangle$ from $5\\times10^{10}{\\rm L_{\\odot}}$ to $8\\times10^{10} {\\rm L_{\\odot}}$ as we go from $z=0.1$ to $z=0.5$, since more luminous galaxies tend to show redder $FUV-V$. The relatively good agreement between our results and the individual galaxies measured by other authors indicates that our stacking analysis is able to probe the evolution of the UV excess as well as detailed studies using small galaxy samples.  We also measured the evolution of the intrinsic shape of the UV excess, finding little evidence for evolution in the rest-frame $FUV-NUV$ colors of early-type galaxies in the 6 Gyr since $z=0.6$. Our two highest-redshift data points are slightly redder than the others, possibly due to systematic effects in our K-corrections, but the differences are not significant. The UV colors of our stacked galaxies are also consistent with the slow evolution predicted by HPL, but the change in the intrinsic $FUV-V$ color with redshift is inconsistent with their predictions for the evolution in the UV-optical color.  While HPL do not show any predictions for evolution in $FUV-V$, they do predict the change in $FUV-r$, which should be quite similar to $FUV-V$, to be only $\\sim0.1$ mag.  (See HPL, Figure 3.)  We found evidence that the UV excess might weaken in more luminous, and therefore more massive, galaxies.  The evidence for this trend is marginal at best, however, and this result should be explored further. If true, it would contradict the suggestion that the UV excess is stronger in more metal-rich galaxies, since the most massive galaxies tend also to have the highest metallicities.  We also divided our galaxies into subsamples by density to explore the impact of environment on the UV excess, and found that the trend for bluer $FUV-V$ colors at higher redshift occurs in all environments.  There is no identifiable trend in color with density.  The lack of such trends is surprising, as effects such as ram pressure stripping, felt by galaxies in rich clusters, could lead to changes in the age, mean metallicity or residual star formation rates of the constituent galaxies.  While several open questions remain, including the possible variation in the UV excess with stellar mass, it appears that the average early-type galaxy evolves in roughly the same way as the individual cluster galaxies measured thus far. The measured evolution, along with the lack of significant trends in color with galaxy environment, should provide interesting constraints for future models of stellar evolution."
    },
    "0808/0808.1961_arXiv.txt": {
        "abstract": "{ We report on the environmental dependence of properties of galaxies around the RDCSJ0910$+54$ cluster at $z=1.1$. We have obtained multi-band wide-field images of the cluster with Suprime-Cam and MOIRCS on Subaru and WFCAM on UKIRT. Also, an intensive spectroscopic campaign has been carried out using LRIS on Keck and FOCAS on Subaru. We have collected 161 spectra with secure redshifts, with which we calibrate a larger sample of photometric redshifts. We discover a possible large-scale structure around the cluster in the form of three clumps of galaxies. Two out of the three newly discovered clumps of galaxies are detected as extended X-ray sources, suggesting that they are bound systems. There seem to be filaments of galaxies in between the clumps. This is potentially one of the largest structures found so far in the $z>1$ Universe. We then examine stellar populations of galaxies in the structure. First, we quantify the color-radius relation. Red galaxies have already become the dominant population in the cores of rich clusters at $z\\sim1$, and the fraction of red galaxies has not strongly changed since then. The red fraction depends on richness of clusters in the sense that it is higher in rich clusters than in poor groups. We confirm that this trend is not due to possible biases in photometric redshifts. Next, we examine red sequence galaxies. The luminosity function of red galaxies in rich clusters is consistent with that in local clusters. On the other hand, luminosity function of red galaxies in poor groups shows a deficit of faint red galaxies. This confirms our earlier findings that galaxies follow an environment-dependent down-sizing evolution. Interestingly, the truncation magnitude of the red sequence appears brighter than that found in the RDCS~J1252$-$29 field at $z=1.24$. This suggests that there is a large variation in the evolutionary phases of galaxies in groups with similar masses. Further studies of high redshift clusters will be a promising way of addressing the role of nature and nurture effects in shaping the environmental dependence of galaxy properties observed in the local Universe. }{}{}{}{} ",
        "introduction": "While galaxy clusters at $z>1$ are now within our reach, larger scale structures are still difficult to locate and characterize. There are a handful of $z > 1$ clusters known to date and detailed analyses of their cluster galaxies have been made (e.g., \\citealt{blakeslee03,lidman04,rosati04,nakata05,stanford05,stanford06,mei06a,mei06b,demarco07,tanaka07a}). There is clear evidence that red early-type galaxies have become a dominant population by $z=1$ in rich clusters. A tight red sequence has already been in place (e.g., \\citealt{blakeslee03,lidman04}) and spectra of $z\\sim1$ cluster galaxies do not typically show signs of active star formation \\citep{mullis05,demarco07}. The majority of cluster galaxies seem to have become red and dead by $z=1$.  While rich clusters have been intensively studied, poor groups at $z>1$ remain largely unexplored. This has been hampered by the fact that high redshift poor groups are difficult to find.  But the recent advent of large telescopes with wide-field capabilities along with deep X-ray surveys have pushed studies of poor groups beyond a redshift of unity. These high redshift poor groups are particularly interesting since they should be the precursors of present-day clusters with moderate masses. Furthermore, there is some observational evidence that galaxy properties depend on cluster richness in the sense that cluster galaxies are more evolved than group galaxies \\citep{zabludoff98,tran01,merchan02,balogh02,martinez02,tanaka05,tanaka07a,koyama07,gilbank08}. Galaxies in rich clusters and in poor groups may follow different evolutionary paths. Therefore, it is essential to observe the full range of local galaxy density at various redshifts up to and beyond $z \\sim 1$.  We are undertaking a high redshift cluster survey with the Subaru telescope \\citep{kodama05}, taking advantage of its wide-field imaging capabilities. Our aim is to observe a range of environments at various redshifts and improve our understanding of galaxy evolution as a function of environment. As part of this survey, we have observed several high redshift clusters \\citep{kodama05,nakata05,umetsu05,tanaka05,tanaka06,tanaka07a,tanaka07b,koyama07}. In this paper, we focus on one of the $z>1$ clusters from our sample, RDCSJ0910+54, which was discovered in the ROSAT Deep Cluster Survey \\citep{rosati98}. \\citet{stanford02} carried out spectroscopic follow-up observations of this cluster, which was confirmed to lie at $z=1.1$, and presented the results of a Chandra observation which measured a relatively high X-ray temperature. \\citet{mei06a} presented a detailed analysis of the red sequence galaxies using high resolution ACS images. Following these studies, we have carried out deep and wide-field observations of this cluster and we present combined analyses of the photometric and spectroscopic properties of galaxies over a wider field around the central cluster at $z=1.1$.   The layout of this paper is as follows. We summarize our observations in Section 2. Section 3 describes photometric catalogs and photometric redshifts. Then, we present large-scale structures around the RDCSJ0910 cluster in Section 4. We carry out detailed analyses of the photometric and spectroscopic properties of galaxies in the structures in Sections 5 and 6. Section 7 discusses implications of our results for galaxy evolution, and the paper is summarized in Section 8.  Unless otherwise stated, we adopt H$_0=70\\rm km\\ s^{-1}\\ Mpc^{-1}$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_\\Lambda =0.7$. Magnitudes are on the AB system. We use the following abbreviations: CMD for color-magnitude diagram, and LF for luminosity function. ",
        "conclusions": "We have carried out the intensive multi-band wide-field observations of the cluster RDCSJ0910 at $z=1.1$ with Suprime-Cam, WFCAM, MOIRCS, FOCAS, and LRIS. The potential large-scale structure is discovered based on the photometric redshifts. Two of the three newly discovered systems show the extended X-ray emissions, suggesting that these are physically bound systems. There seems to be filaments in between the clumps. If confirmed to be real, this will be one of the most prominent structure ever found in the $z>1$ Universe.   We then look into properties of galaxies in the structure. The fraction of red galaxies in rich clusters has not changed since $z\\sim1$, suggesting that the environmental dependence of galaxy colors is already in place. Poor groups at $z\\sim1$ show a lower red fraction than rich clusters.   We find that the red sequence of galaxies in relatively poor systems appears to be truncated at $K_s=21$, while that in richer systems exhibit a clear sequence down to fainter magnitudes. This confirms our earlier claim that the red sequence grows from bright magnitudes to faint magnitudes, and the build-up of the red sequence happens earlier in richer systems. Therefore, galaxies follow the environment-dependent down-sizing --- massive galaxies in higher density environments are more evolved than less massive galaxies in less dense environments. Interestingly, the red sequence is truncated at brighter magnitudes than that found in the groups in the RDCSJ1252 field at $z=1.24$. Although the groups have similar masses, the truncation magnitude is $\\sim1$ mag brighter in the RDCSJ0910 field at $z=1.1$.   Both the initial conditions of galaxy formation and the environmental effects determine the fate of a galaxy. It is difficult to disentangle these two effects, but $z>1$ might be an era when the nature effects can be seen. Further detailed studies on $z>1$ clusters would be a promising way of addressing this fundamental issue of galaxy evolution."
    },
    "0808/0808.3007_arXiv.txt": {
        "abstract": "In this introductory review presented at the IAU Symposium 253 ``Transiting Planets'', I summarize the path from the initial 1995 radial-velocity discovery of hot Jupiters to the current rich panoply of investigations that are afforded when such objects are observed to transit their parent stars. Forty transiting exoplanets are now known, and the time for that population to double has dropped below one year.  It is only for these objects that we have direct estimates of their masses and radii, and for which (at the current time) we can undertake direct studies of the chemistries and dynamics of their atmospheres. Informed by the successes of hot Jupiter studies, I outline a path for the spectroscopic study of certain habitable exoplanets that obviates the need for direct imaging. ",
        "introduction": "Although the possibility of detecting planets that transit their parent stars had been described decades before (Struve 1952; Rosenblatt 1971; Borucki \\& Summers 1984), it was not until the discovery of hot Jupiters by the radial-velocity method (Mayor \\& Queloz 1995; Butler et al.\\ 1997) that dedicated transit surveys were begun in earnest.  The justification was straightforward:  Whereas Jupiter analogs would present a signal only once every 12 years and have a geometric transit probability of 0.1\\%, hot Jupiters orbiting at 0.05~AU from their stars have orbital periods of roughly 3 days and geometric transit probabilities of 10\\%. The first success came in 1999 with HD~209458b (Charbonneau et al.\\ 2000; Henry et al.\\ 2000; Mazeh et al.\\ 2000), but this object was first identified by radial-velocity monitoring and subsequently determined to transit.  A barrage of ground and space based transit surveys ensued (see review by Charbonneau et al.\\ 2007) yet it was not until 2003$-$2004 that the OGLE survey delivered the first handful of exoplanets identified by transits (e.g.\\ Udalski et al. 2002; Konacki et al.\\ 2003; Bouchy et al.\\ 2004).  Although the relative importance of the OGLE planets may arguably diminish in the future (since they orbit much more distant stars and hence follow-up observations are difficult or precluded), they played an absolutely pivotal role in quenching the thirst of the community after the 2000-2002 drought. The first wide-field survey (geared toward identifying transiting planets of nearby stars) to deliver a discovery was TrES, with the 2004 discovery of the aptly named TrES-1 (Alonso et al.\\ 2004).  By 2006, three additional surveys (XO, McCullough et al.\\ 2006; HAT, Bakos et al.\\ 2007; and SuperWASP, Collier Cameron et al.\\ 2007) produced their first detections and all 3 are now operating in production-line mode with numerous new planets in the past 18 months.  Fig.\\ 1 makes clear the very recent rise in the rate of discovery.  Thus, this conference convened at a very special moment in history, when the doubling time for new discoveries of transiting planets (40 as of 31 July 2008) first dipped below one year. A little over two years ago I co-authored a review (Charbonneau et al.\\ 2007) of the 9 transiting exoplanets known at the time, and we concluded that our paper would surely be eclipsed by the coming rapid pace of new discoveries. It has been a rare joy to see a publication rendered obsolete so quickly.  My personal excitement over transiting planets stems from three distinct opportunities that they present:  First, transiting planets permit us to determine their masses and radii in a manner that is nearly free of astrophysical assumption.  Those estimates, in turn, bear upon our knowledge of the bulk composition and physical structure of these worlds, and surely constrain models of their formation (just imagine how our understanding of the Solar system would differ if he had no knowledge of the bulk composition or true masses of its planets!).  Second, by observing the modulation of the combined light of the planet and star as the two undergo mutual eclipses, astronomers have devised cunning techniques to study the chemistries and dynamics of their atmospheres \\textit{without} the need to image the planet directly. Third, I feel that it is these methods, which I will call the dynamics-based approach, rather than high-contrast-ratio imaging, that will permit the first studies of the compositions of habitable planets and their atmospheres.  A particularly attractive opportunity exists for such planets orbiting low-mass M-dwarf stars.  In the sections below, I will consider each of these three motivations in modestly more detail and provide a sampling of recent results, leaving a thorough and detailed review to my Vulcan colleagues in their respective chapters of the book that follows.  \\begin{figure}[htb] \\begin{center} \\includegraphics[width=5.4in]{charbonneau_253_f1.eps} \\caption{Histogram of the number of transiting exoplanets (as of 31 July 2008) with confirming radial velocity measurements as a function of the year that the discovery paper that had been accepted for publication in a refereed journal was released to the general public.  An incomplete list of notable events is given on the right hand side.  Ground-based transit surveys were not pursued in earnest until after the radial-velocity discovery of hot Jupiter exoplanets in 1995$-$1996, since hot Jupiters present a geometric probability of a transit that is 100$\\times$ greater than that of a planet orbiting at the Jupiter-Sun separation.  As this plot shows only those planets meeting the publication criterion as of mid-2008, the 2007$-$2008 bar is incomplete.  At the current epoch, the time for the number of transiting exoplanets to double is less than 1 year.} \\label{fig1} \\end{center} \\end{figure} \\newpage ",
        "conclusions": ""
    },
    "0808/0808.3141_arXiv.txt": {
        "abstract": "We report on Suzaku observations of four unidentified sources from the INTEGRAL and {\\it Swift} BAT Galactic plane surveys. All the sources have a large neutral hydrogen column density and are likely members of an emerging class of heavily absorbed high mass X-ray binary (HMXB) first identified in INTEGRAL observations. Two of the sources in our sample are approximately constant flux sources, one source shows periodic variation and one source exhibited a short, bright X-ray outburst. The periodicity is transient, suggesting it is produced by a neutron star in an elliptical orbit around a stellar wind source. We analyze the flaring source in several segments to look for spectral variation and discuss the implications of the findings for the nature of the source. We conclude that all four sources in our sample can be identified with the emerging class of highly absorbed HMXBs, that one is a newly identified transient X-ray pulsar and that at least one is a newly identified supergiant fast X-ray transient (SFXT). ",
        "introduction": "With its first observation, INTEGRAL \\citep{Winkler03} began a collection of highly absorbed sources, some of which may represent a new class of binary (see \\cite{Kuulkers05} for a review; K05 hereafter). IGRJ16318-4848 was a bright X-ray source with extreme absorption that would later be shown to have bright X-ray emission lines. Since then, the INTEGRAL Galactic plane survey has discovered several more similar sources, suggesting that these objects form a new class of X-ray binary, characterized by high absorption columns, slow spin periods and occasional bright X-ray outbursts. The first INTEGRAL catalog \\citep{Bird04} notes 28 unidentified sources from the INTEGRAL Galactic plane survey, ten of which have been followed up at X-ray and optical wavelengths and have been shown to belong to this new source class.  Nearly simultaneously to the discovery of this potentially new class of source by INTEGRAL, the {\\it Swift} BAT Galactic plane survey began to uncover many new sources from its own survey. It was speculated, due to their absence from previous soft X-ray surveys, that some of these sources may be highly absorbed binaries, possibly members of the new class of INTEGRAL binaries.  Several objects identified in the first INTEGRAL catalog, as of Fall 2006, remained without significant follow-up in the X-ray or at other wavelengths. During the Suzaku cycle 1 call for proposals, we proposed to follow-up three INTEGRAL detected sources that were suspected to be members of the new class but had yet to be studied in detail. We proposed, furthermore, to observe two newly identified sources from the BAT Galactic plane survey to determine whether they were also members of this new source class.  In this paper, we detail the analysis of the Suzaku data collected on these sources, discuss the likelihood that these sources are similar to the originally discovered IGR sources and briefly discuss the possible nature of the new class. The paper is organized as follows: in \\S2 we describe the observations and data analysis; in \\S3 we present our results; in \\S4 we discuss the similarity of each source and the propriety of calling each a member of the IGR source class and discuss the potential nature of these sources; in \\S5 we summarize our results and conclusions. ",
        "conclusions": "Our primary goal in this study is to determine whether these four sources exhibit similar characteristics to the emerging new class of highly absorbed IGR sgHMXB or SFXT sources. The defining characteristics of this new class are i) high absorption column ($>$1$\\times$10$^{23}$~cm$^{-2}$), ii) periodicity on timescales of a few to $\\sim$100 minutes, generally interpreted as a neutron star spin period, iii) periodicity on timescales of $\\sim$10 days, generally interpreted as a binary orbital period and iv) occasionally strong X-ray emission lines.  \\noindent {\\bf i)} In all four sources the total absorbing column is $\\gtrsim$1$\\times$10$^{23}$~cm$^{-2}$. Generally, we do not see significant changes in the absorption over the duration of the observations. While this is not surprising during a single day-long observation, which probably samples a small segment of the orbital phase, it is somewhat surprising that two observations of SWJ20006.+3210, separated by more than six months (and furthermore apparently at different parts of the orbital phase) also show no significant difference in \\NH. A dramatic increase in absorbing column is seen in IGR16195-4945 for a duration of $\\sim$30~ks, during a bright outburst (when all parameters are left free to vary). Since this dramatic increase in \\NH is associated with a dramatic increase in flux, one might interpret this as an increase in emission due to interaction of the neutron star with a density enhancement. Such a scenario might be expected due to the passage of the neutron star through a thick disk associated with the donor star. Since an interpretation in which \\NH is approximately constant and only $\\Gamma$ varies also produces adequate spectral fits, an alternative model in which the increased emission is due to variable accretion is also possible.  \\noindent {\\bf ii)} In SWJ2000.6+3210, we see a period of 1056~s which is only observed during one of two observations of the source. This is a newly identified transient X-ray pulsar. Given the low magnetic field implied by the relatively slow rotation period of 1056~s, it is not surprising that the same mechanism that produces periodic observations near periastron is too weak to produce an observable period near apoastron where the wind density will be lower.  \\noindent {\\bf iii)} Due to the short nature of the observations (generally $\\sim$1 day), we do not expect to directly measure orbital periods, previously reported to be on the order of ten days in other sgHMXBs and, indeed, we do not. During observations of the one source for which we have well separated observations we also do not see evidence of an orbital period. We have, however, used our observations to calculate a predicted orbital period for IGRJ16195-4945. Assuming that IGRJ16195-4945 is a SFXT and that the bright flare we see signals the interaction of the compact object with an equatorial wind of the donor star, we have combined our Suzaku observations with long-term {\\it Swift}-BAT observations to predict the orbital period for this source as P$\\sim$16 days. This is unusually short compared to previously measured SFXT orbital periods, but is consistent with orbital periods previously measured for the highly absorbed IGR sources. More consistent monitoring, perhaps using the {\\it Swift} satellite, would be very useful in confirming or refuting this predicted period and in characterizing the orbital periods of these highly absorbed binaries in general.  \\noindent {\\bf iv)} We find only weak evidence of Fe fluorescence emission in one source, and only upper limits to Fe lines in the other three sources. This is similar to the behavior reported in the original INTEGRAL highly absorbed X-ray binary sources \\citep{Kuulkers05} in which only 1 of 10 sources showed strong emission lines while the others showed only weak lines or upper limits. The Fe line measurements that we report here are either the first for the source in the literature (IGRJ16465-4507, IGRJ16493-4348, SwiftJ2000.6+3210) or several times more restrictive than previous measurements (IGRJ16195-4945). Moreover, our measurements are similarly or more restrictive than those reported on similar sources in the literature (see, e.g., \\cite{Sidoli05, Patel04, Rodriguez03}).  While none of these sources is ruled out from the IGR class based on these observations, further observations would be helpful in refining their nature. Of particular value would be periodic monitoring of all four systems."
    },
    "0808/0808.3377_arXiv.txt": {
        "abstract": "We use {\\em Swift}/BAT Earth occultation data at different geomagnetic latitudes to derive a sensitive measurement of the Cosmic X-ray background (CXB) and of the Earth albedo emission in the 15--200\\,keV band. We compare our CXB spectrum with recent ({\\it INTEGRAL}, BeppoSAX) and past results (HEAO-1) and find good agreement. Using an independent measurement of the CXB spectrum we are able to confirm our results. This study shows that the BAT CXB spectrum has a normalization $\\sim8\\pm3$\\,\\% larger than the HEAO-1 measurement. The BAT accurate Earth albedo spectrum can be used to predict the level of photon background for satellites in low Earth and mid inclination orbits. ",
        "introduction": "There is a general consensus that the cosmic X-ray background (CXB), discovered more than 40 years ago \\citep{giacconi62}, is produced by integrated emission of extra-galactic point sources. The deepest X-ray surveys to date \\citep{giacconi02,alexander03,hasinger04} have shown that up to virtually 100\\% of the $<$ 2\\,keV CXB radiation  is accounted for by Active Galactic Nuclei (AGN) hosting accreting super-massive black holes (SMBHs). However, the fraction of CXB emission resolved into AGNs declines with energy   being $<50$\\% above 6\\,keV \\citep{worsley05}. The unresolved  component  may be attributed to the emission of a yet undetected population of highly-absorbed  AGN. These AGN should be characterized by having column densities $\\sim10^{24}$ H-atoms cm$^{-2}$ and a space density peaking at redshift below 1 \\citep{worsley05}. Such a population of Compton-thick AGN is invoked by population synthesis models \\citep[e.g.][]{comastri95,treister05,gilli07} to reproduce the peak of the CXB emission at 30\\,keV \\citep{marshall80}.  Thus, an accurate measurement of the CXB spectrum in the 15--200\\,keV energy range is important to assess and constrain the number density of Compton-thick AGNs. Such measurements are complicated by the fact that instruments sensitive in this energy range are dominated by internal detector background and are not designed to measure the CXB spectrum directly (excluding HEAO-1 A2). The typical approach is to produce an $ON-OFF$ measurement, where taking the difference between the $ON$ and the $OFF$ pointings eliminates the internal background component.  There are different methods to obtain a suitable $OFF$ observation; the HEAO-1 measurement of the CXB spectrum in the 13--180\\,keV range was obtained by blocking the aperture with a movable CsI  crystal. Also the Earth disk can be used to modulate the CXB emission. This approach is  the one used in recent CXB intensity measurements performed by {\\it INTEGRAL} and BeppoSAX \\citep{churazov07,frontera07}.    Here we report on two independent measurements  of the CXB emission using  {\\em Swift}/BAT. For the first method, we use the Earth occultation technique similarly to the {\\it INTEGRAL} and BeppoSAX analyses while for the second one we make use of the spatial distribution of the BAT background.  The structure of the paper is as follows. In $\\S$~\\ref{sec:obs} we present the details of the observations and describe the BAT background components. We also derive a rate-rigidity relation which is fundamental for suppressing the background variability due to Cosmic Rays (CRs). In $\\S$~\\ref{sec:analysis} we present the details of the Earth's occultation episodes undergone by BAT and the analysis method for the occultation measurement. In $\\S$~\\ref{sec:err}, we discuss all sources of uncertainties which affect our occultation measurement which is then presented in $\\S$~\\ref{sec:results}. The alternative measurement used to verify the results of the occultation analysis is reported in $\\S$~\\ref{sec:quad}. We discuss the broad band properties of the CXB and Earth spectra in $\\S$~\\ref{sec:disc}. Finally, the last section summarizes our findings. ",
        "conclusions": ""
    },
    "0808/0808.2976_arXiv.txt": {
        "abstract": "Following the traditional naming of ``eruptive flare\" and ``confined flares\" but not implying a causal relationship between flare and coronal mass ejection (CME), we refer to the two kinds of large energetic phenomena occurring in the solar atmosphere as eruptive event and confined event, respectively: the former type refers to flares with associated CMEs, while the later type refers to flares without associated CMEs. We find that about 90\\% of X-class flares, the highest class in flare intensity size, are eruptive, but the rest 10\\% confined. To probe the question why the largest energy release in the solar corona could be either eruptive or confined, we have made a comparative study by carefully investigating 4 X-class events in each of the two types with a focus on the differences in their magnetic properties. Both sets of events are selected to have very similar intensity (X1.0 to X3.6) and duration (rise time less than 13 minutes and decaying time less than 9 minutes) in soft X-ray observations, in order to reduce the bias of flare size on CME occurrence. We find no difference in the total magnetic flux of the photospheric source regions for the two sets of events. However, we find that the occurrence of eruption (or confinement) is sensitive to the displacement of the location of the energy release, which is defined as the distance between the flare site and the flux-weighted magnetic center of the source active region. The displacement is 6 to 17 Mm for confined events, but is as large as 22 to 37 Mm for eruptive events, compared to the typical size of about 70 Mm for active regions studied. In other words, confined events occur closer to the magnetic center while the eruptive events tend to occur closer to the edge of active regions. Further, we have used potential-field source-surface model (PFSS) to infer the 3-D coronal magnetic field above source active regions. For each event, we calculate the coronal flux ratio of low corona ($<$ 1.1 $R_\\odot$) to high corona ($\\geq$ 1.1 $R_\\odot$). We find that the confined events have a lower coronal flux ratio ($< 5.7$), while the eruptive events have a higher flux ratio($> 7.1$). These results imply that a stronger overlying arcade field may prevent energy release in the low corona from being eruptive, resulting in flares without CMEs. A CME is more likely to occur if the overlying arcade field is weaker. ",
        "introduction": "Coronal mass ejections (CMEs) and flares are known to be the two most energetic phenomena that occur in the atmosphere of the Sun, and have profound effects on the physical environment in the geo-space environment and human technological systems. In this paper, we intend to understand the physical origins of CMEs and flares by comparatively studying two different kinds of energetic phenomena, both of which have almost identical flares, but one is associated with CMEs and the other not. In the past when lacking direct CME observations, these two kinds of phenomena have been called as eruptive flares and confined flares, respectively~\\citep[e.g.,][]{Svestka_Cliver_1992}. An eruptive flare (also called a dynamic flare) is usually manifested as having two-ribbons and post-flare loops in H$\\alpha$ imaging observations and long duration (e.g., tens of minutes or hours) in soft X-ray, while a confined flare occurs in a compact region and last for only a short period (e.g., minutes). Following this convention but trying not to imply a causal relation between flare and CME, hereafter we refer the flare event associated with an observed CME as an ``eruptive event\", and the flare event not associated with a CME as a ``confined event\".  It has been suggested based on observations that CMEs and flares are the two different phenomena of the same energy release process in the corona~\\citep[e.g.,][]{Harrison_1995, Zhang_etal_2001a, Harrison_2003}. They do not drive one another but are closely related. In particular,~\\citet{Zhang_etal_2001a} and~\\citet{Zhang_etal_2004} showed that the fast acceleration of CMEs in the inner corona coincides very well in time with the rise phase (or energy release phase) of the corresponding soft X-ray flares, strongly implying that both phenomena are driven by the same process at the same time, possibly by magnetic field reconnections. However, this implication raises another important question of under what circumstances the energy release process in the corona leads to an eruption, and under what circumstances it remains to be confined during the process. An answer to this question shall shed the light on the origin of flares as well as CMEs.  The occurrence rate of eruptive events depends on the intensity and duration of flares. A statistical study performed by~\\citet{Kahler_etal_1989} showed that the flares with longer duration tend to be eruptive, while more impulsive flares tend to be confined. By using the CME data from Solar Maximum Mission (SMM) and the flare data from GOES satellites during 1986 -- 1987,~\\citet{Harrison_1995} found that the association ratio of flares with CMEs increases from about 7\\% to 100\\% as the flare class increases from B-class to X-class, and from about 6\\% to 50\\% as the duration of flares increases from about 1 to 6 hours.~\\citet{Andrews_2003} examined 229 M and X-class X-ray flares during 1996 -- 1999, and found that the CME-association rate, or eruptive rate is 55\\% for M-class flares and 100\\% for X-class flares. With a much lager sample of 1301 X-ray flares,~\\citet{Yashiro_etal_2005} obtained a similar result that the eruptive rates of C, M, and X-class flares are $16-25$\\%, $42-55$\\%, and $90-92$\\%, respectively.~\\citet{Yashiro_etal_2005} work showed that a flare is not necessarily associated with a CME even if it is as intense as an X-class flare. Such kind of confined but extremely energetic events were also reported by~\\citet{Feynman_Hundhausen_1994} and~\\citet{Green_etal_2002}.  The studies mentioned above indicate the probability of CME occurrence for a given flare. On the other hand, there is also a probability of flare occurrence for a given CME. There are CMEs that may not be necessarily associated with any noticeable X-ray flares.~\\citet{Zhang_etal_2004} reported an extremely gradually accelerated slow CME without flare association, implying that the non-flare eruptive event tends to be slowly driven. By Combining the corona data from SMM and 6-hr soft X-ray data from GOES satellite,~\\citet{St_Cyr_Webb_1991} reported that about 48\\% frontside CMEs were associated with X-ray events near the solar minimum of solar cycle 21. Based on SOHO observations,~\\citet{Wang_etal_2002a} studied 132 frontside halo CMEs, and found that the association rate of CMEs with X-ray flares greater than C-class increased from 55\\% at solar minimum to 80\\% near solar maximum. With 197 halo CMEs identified during 1997 -- 2001,~\\citet{Zhou_etal_2003} concluded that 88\\% CMEs were associated with EUV brightenings.  Some attempts have been made to explain the confinement or the eruptiveness of solar energetic events in the context of the configuration of coronal magnetic field.~\\citet{Green_etal_2001} analyzed the 2000 September 30 confined event utilizing multi-wavelength data, and suggested that the event involve magnetic reconnection of two closed loops to form two newly closed loops without the opening of the involved magnetic structure.~\\citet{Nindos_Andrews_2004} made a statistical study of the role of magnetic helicity in eruption rate. They found that the coronal helicity of active regions producing confined events tends to be smaller than the coronal helicity of those producing eruptive events.  In this paper, we address the eruptive-confinement issue of solar energetic events with the approach of a comparative study. We focus on the most energetic confined events that produce X-class soft X-ray flares but without CMEs. While the majority of X-class flares are eruptive, a small percentage (about 10\\%) of them are confined. The magnetic properties of these confined events shall be more outstanding than those less-energetic confined events. For making an effective comparison, we select eruptive events with X-ray properties, in terms of intensity and duration, very similar to those selected confined events. The differences on magnetic configuration between these two sets of events shall most likely reveal the true causes of eruption or confinement. How to select events and the basic properties of these events are described in section 2. Detailed comparative analyses of the two sets of events are given in section~\\ref{sec_photosphere} and~\\ref{sec_corona}, which are on the properties of the photospheric magnetic field distribution and the extrapolated coronal magnetic field distribution, respectively. In section~\\ref{sec_summary}, we summarize the paper. ",
        "conclusions": "\\label{sec_summary} In summary, among the 104 X-class flares occurred during 1996 -- 2004, we found a total of 11 ($\\sim10\\%$) are confined flares without associated CMEs, and all the others ($\\sim90\\%$) are eruptive flares associated with CMEs. Four suitable confined flares are selected to make a comparative study with four eruptive flares, which are similar in X-ray intensity and duration as those confined events. We have carefully studied the magnetic properties of these events both in the photosphere and in the corona. The following results are obtained:  (1) In the photosphere, we can not find a difference of the total magnetic fluxes in the surface source regions between the two sets of events. However, there is an apparent difference in the displacement parameter, which is defined as the surface distance between the flare site and the center of magnetic flux distribution. For the confined events, the displacement is from 6 to 17 Mm, while for those eruptive events it is from 22 to 37 Mm. This result implies that the energy release occurring in the center of an active region is more difficult to have a complete open eruption, resulting in a flare without CME. On the other hand, the energy release occurring away from the magnetic center has a higher probability to have an eruption, resulting in both flares and CMEs. Whether an eruption could occur or not may be strongly constrained by the overlying large scale coronal magnetic field. The overlying coronal magnetic field shall be strongest and also longest along the vertical direction over the center of an active region. On the other hand, the overlying constraining field shall be weaker if the source is away from the center. This scenario is further supported by our study of coronal magnetic field.  (2) Calculation of coronal magnetic field shows that the flux ratio of the magnetic flux in the low corona to that in the high corona is systematically larger for the eruptive events than that for the confined events. The magnetic flux ratio for the confined events varies from 1.6 to 5.7, while the ratio for the eruptive events from 7.1 to 10.2. However, there is no evident difference between the two sets of events in the total magnetic flux straddling over neutral lines, and there is only a weak trend indicating a systematic difference in the low- and high-corona magnetic fluxes. This low-to-high corona magnetic flux ratio serves as a proxy of the strength of the inner core magnetic field, which may play an erupting role, relative to the strength of the overlying large scale coronal magnetic field, which may play a constraining role to prevent eruption. The lower this ratio, the more difficult the energy release in the low corona can be eruptive.  There is variety of theoretical models on the initiation mechanism of CMEs and the energy release of flares ~\\citep[e.g.,][]{Sturrock_1989, Chen_1989, vanBallegooijen_Martens_1989, Forbes_Isenberg_1991, Moore_Roumeliotis_1992, Low_Smith_1993, Mikic_Linker_1994, Antiochos_etal_1999, Lin_Forbes_2000}. These models differ in pre-eruption magnetic configurations, trigger processes, or where magnetic reconnection occurs. Nevertheless, in almost all these models, the magnetic configuration involves two magnetic regimes, one is the core field in the inner corona close to the neutral line, the other is the large scale overlying field or background field. The core field is treated as highly sheared or as a fully-fledged flux rope; in either case, the core field stores free energy for release. On the other hand, the overlying field is regarded as potential and considered to be main constraining force to prevent the underlying core field from eruption or escaping.~\\citet{Torok_Kliem_2005} and~\\citet{Kliem_Torok_2006} recently pointed out that the decrease of the overlying field with height is a main factor in deciding whether the kink-instability (in their twist flux rope model) leads to a confined event or a CME. On the other hand,~\\citet{Mandrini_etal_2005} reported the smallest CME event ever observed by 2005, in which the CME originated from the smallest source region, a tiny dipole, and developed into the smallest magnetic cloud. They suggested that the ejections of tiny flux ropes are possible. Therefore, it is reasonable to argue that, whether an energy release in the corona is eruptive or confined, is sensitive to the balance between the inner core field and the outer overlying field. Our observational results seem to be consistent with this scenario.  This study is only a preliminary step to investigate the confinement and/or eruptiveness of solar flares, or coronal energy releases in general. However, it demonstrates that the distribution of magnetic field both in the photosphere and in the corona may effectively provide the clue of the possible nature of an energetic event: whether a flare, a CME or both. To further evaluate the effectiveness of this methodology, a more robust study involving more events is needed."
    },
    "0808/0808.1414_arXiv.txt": {
        "abstract": "{} {We present abundance analysis based on high resolution spectra of 105 isolated red giant branch (RGB) stars in the Galactic Globular Cluster NGC~6121 (M4). Our aim is to study its star population in the context of the multi-population phenomenon recently discovered to affect some Globular Clusters.} {The data have been collected with FLAMES+UVES, the multi-fiber high resolution facility at the ESO/VLT@UT2 telescope. Analysis was performed under LTE approximation for the following elements: O, Na, Mg, Al, Si, Ca, Ti, Cr, Fe, Ni, Ba, and NLTE corrections were applied to those (Na, Mg) strongly affected by departure from LTE. Spectroscopic data were coupled with high-precision wide-field UBVI$_{\\rm C}$ photometry from WFI@2.2m telescope  and infrared JHK photometry from 2MASS.} {We derived an average $\\rm {[Fe/H]}=-1.07\\pm0.01$ (internal error), and an $\\alpha$ enhancement of $\\rm {[\\alpha/Fe]}=+0.39\\pm0.05$ dex (internal error). We confirm the presence of an extended Na-O anticorrelation, and find two distinct groups of stars with significantly different Na and O content. We find no evidence of a Mg-Al anticorrelation. By coupling our results with previous studies on the CN band strength, we find that the CN strong stars have higher Na and Al content and are more O depleted than the CN weak ones. The two groups of Na-rich, CN-strong and Na-poor, CN-weak stars populate two different regions along the RGB. The Na-rich group defines a narrow sequence on the red side of the RGB, while the Na-poor sample populate a bluer, more spread portion of the RGB. In the $\\rm {U}$ vs.\\ $\\rm {U-B}$ color magnitude diagram the RGB spread is present from the base of the RGB to the RGB-tip. Apparently, both spectroscopic and photometric results imply the presence of two stellar populations in M4. We briefly discuss the possible origin of these populations.} {} ",
        "introduction": "Observational evidence for variations in the chemical composition of light elements in Globular Cluster (GC) stars were known since Cohen (1978), who noted a scatter in Na among stars in M3 and M13. During the last few decades, high resolution spectroscopic studies have definitely confirmed that a GC stellar population is not chemically homogeneous. Even if GC stars are generally homogeneous in their Fe-peak element content, they show large star-to-star abundance variations in the light elements such as C, N, O, Na, Mg, Al, and others (see Gratton et al.\\ 2004 for a review).  During the last twenty years, it has become clear that in red giant stars the abundances of some of these elements follow a well defined pattern. In particular, there are clear anticorrelations between the Na and O content, and between Mg and Al. Variations in the molecular CH, CN and NH band strengths, due to a spread in the abundances of carbon and nitrogen have been observed, as well as anticorrelations between CH and CN strengths, and, in some cases, a clear bimodality in the CN content.  Despite the spectroscopic observational evidence collected in more than thirty years, the pattern in the light elements is not yet well understood. Two scenarios have been proposed to explain this observed chemical heterogeneity: the evolutionary scenario and the primordial one, both apparently supported by observations.  The observed decline of C content (e.g. in M13, as shown by Smith \\& Briley 2006) and the decreasing ratio $\\phantom{}^{12}\\rm {C}/\\phantom{}^{13}\\rm {C}$ (observed in M4 and NGC~6528 by Shetrone 2003) along the RGB phase, support the evolutionary scenario. According to this theory, the origin of the observed star-to-star scatter in some elements is due to the mixing of CNO-cycle-processed material transported, in a way not well understood yet, to the stellar surface. In this way the observed anticorrelations would be present in the evolved stages of the life of stars, after the RGB bump.  At odds with this scenario, in the last few years, spectroscopic studies have revealed light element abundance variations in unevolved main sequence stars and less-evolved RGB stars, fainter than the RGB bump. The Na-O anticorrelation was found at the level of the main sequence turn-off (TO) and sub giant branch (SGB) in M13 (Cohen \\& Mel\\'endez 2005), NGC~6397 and NGC~6752 (Carretta et al.\\ 2005; Gratton et al.\\ 2001), NGC~6838 (Ram\\'irez \\& Cohen 2002) and 47~Tuc (Carretta et al.\\ 2004). In 47~Tuc, a bimodal distribution in the CN strengths, similar to that found among RGB stars (Norris \\& Freeman 1979), was found also in the main sequence (MS) by Cannon et al.\\ (1998). Moreover, Grundahl et al.\\ (2002) have shown that in NGC~6752 the observed scatter in the Str$\\rm {\\ddot {o}}$mgren index $c_{1}$ is due to the abundance variations in NH bands in stars both brighter and fainter than the RGB bump. This result is consistent with a primordial scenario since theory does not predict significant mixing below the luminosity of the first dredge-up, observationally corresponding to the magnitude of the RGB bump. These observations suggest that the light element variations should be primordial, i.e. they are derived from the chemical composition of the primordial site where the GC stars have been generated, or alternatively, that a second generation of stars has been formed from a medium enriched in some elements (e.g., the self enrichment model by Ventura et al.\\ 2001). A primordial scenario in which such GCs have experienced multiple episodes of star formation challenges the paradigm that GCs host a single stellar population, i.e. that stars of a given cluster are coeval and chemically homogeneous.  Very recently, a spectacular, and somehow unexpected confirmation that, at least in some GCs, the origin of the chemical anomalies must be primordial came from high precision photometry from Hubble Space Telescope observations. The first object challenging the paradigm of GCs hosting simple stellar populations was $\\omega$~Cen. As shown by Bedin et al.\\ (2004), the MS of $\\omega$~Cen is split into two sequences. But the most exciting discovery came from the spectroscopic investigation by Piotto et al.\\ (2005), who found that the bluest MS is more metal rich than the redder one. The only way to account for the spectroscopic and photometric properties of the two main sequences is to assume that the bluest sequence is also strongly He enhanced, to an astonishingly high Y=0.38.  More recently, Villanova et al.\\ (2007) showed that the two main sequences split in at least four sub giant branches (SGB) which must be connected in some way to the multiplicity of RGBs identified by Lee et al.\\ (1999) and Pancino et al.\\ (2000). The metal content of the different SGBs measured by Villanova et al.\\ (2007) also implies a large age difference among $\\omega$~Cen stellar populations, larger than 1 Gyr. The exact age dispersion depends on the assumed abundances (including the He content) of the different stellar populations, and it is still controversial (Sollima et al.\\ 2007). Villanova et al.\\ (2007) demonstrated that there is also a third MS, running on the red side of the two main MSs, and likely connected with the anomalous RGB-a of Pancino et al.\\ (2000).  Omega Centauri was well-known since the seventies because of its peculiar metallicity distribution. It is the only GC showing iron-peak element dispersion (Freeman \\& Rodgers 1975 and, more recently, Norris et al.\\ 1996, Suntzeff \\& Kraft 1996). In a sense, the findings by Pancino et al.\\ (2000), Bedin et al.\\ (2004), Piotto et al.\\ (2005), Villanova et al.\\ (2007) and Sollima et al.\\ (2007) could simply be considered additional evidence that $\\omega$~Cen is so peculiar that it might not be a GC. Perhaps, as suggested by many authors (Freeman 1993, Hughes \\& Wallerstein 2000), it might simply be the nucleus of a much larger system, likely disrupted by the tidal field of our Galaxy.\\\\ In this sense the most recent discovery by Piotto et al.\\ (2007) that the MS of NGC~2808 is split into three, distinct sequences came as a sort of surprise, shaking at its foundation our understanding of GC stellar populations. NGC~2808 has been always considered a GC, with many peculiar properties regarding its metal content and its color-magnitude diagram, but a genuine, massive GC. Still, it hosts multiple stellar populations.  Moreover, also in this case, in view of the negligible dispersion in iron-peak elements in NGC~2808, the only way so far available to reproduce the three MSs is to assume that there are three populations, characterized by three different helium contents, up to an (again) astonishingly high Y=0.40. Interestingly enough, D'Antona et al.\\ (2005) already made the hypothesis of three groups of stars, with three different helium abundances in order to explain the peculiar, multi-modal Horizontal Branch (HB) of NGC~2808 (Sosin et al.\\ 1997, Bedin et al.\\ 2000). Piotto et al.\\ (2007) simply found them in the form of a MS split. Piotto et al.\\ (2007) also noticed that the different stellar populations in NGC~2808 are consistent with the spectroscopic observations by Carretta et al.\\ (2006), who identified three groups of stars with different Oxygen abundances. The fraction of stars in the three abundance groups is in rough agreement with the fraction of stars in the three MSs.  NGC~2808 and $\\omega$~Cen are, at the moment, the most extreme examples of a rather complex observational scenario. In fact, evidence of multiple populations has been found in other GCs, like NGC~1851 (Milone et al.\\ 2008), NGC~6388 (Siegel et al.\\ 2007, Piotto et al.\\ 2008), and M54 (Piotto et al., in preparation) in the form of a split in the SGB. In NGC~1851 the presence of a group of RGB stars with enhanced Sr and Ba and strong CN bands, among the majority of CN-normal RGB stars (Yong \\& Grundahl 2008), and the presence of a bimodal HB, agrees with the hypothesis of two stellar generations inferred by the observed SGB split.  The extremely peculiar HB of NGC~6388 (Rich et al.\\ 1997), with the presence of extremely hot HB stars (Busso et al.\\ 2007) was already interpreted in terms of multiple, helium enhanced, population by Sweigart \\& Catelan (1998), and in the more detailed analysis by Caloi \\& D'Antona (2007).  As for M54, it has been recognized as the nucleus of the Sagittarius dwarf galaxy (Da Costa \\& Armandroff 1995; Bassino \\& Muzzio 1995; Layden \\& Sarajedini 2000), and it could simply represent what $\\omega$~Cen was long time ago.  It is worth noting that the clusters showing multiplicities in their CMDs are among the most massive GCs of our Galaxy ($\\rm {M > 10^6 M_\\odot}$) and all of them have peculiar HBs, as well as peculiar abundances, including Na-O anticorrelations. However, we must also note that much less massive GCs, with no evidence of multiple populations identified so far, show very large star-to-star abundance variations. It is noteworthy to recognize here that the Na-O anticorrelation has been found in about 20 GCs (see Carretta et al.\\ 2006 for the most updated list).  In summary, it is clear that GCs are not as simple systems as thought in the past. Up to now, we lack a complete explanation for the mechanisms necessary to understand the observational scenario. A systematic study of the chemical abundances of many stars in GCs is needed in order to better understand the star formation history of these objects.  In this work we present a study on the chemical abundances of the GC NGC~6121 (M4) from high resolution spectra of its RGB stars.  As far as we know, M4 shows no evidence of multiple stellar populations in its CMD, and its mass ($\\rm {log {\\frac {M}{M\\odot}}=4.8}$, Mandushev et al.\\ 1991) is much smaller than the mass of the clusters with the photometric peculiarities discussed above.  Chemical abundances from high resolution spectra of M4 RGB stars have been already measured by different groups of investigators: Gratton, Quarta \\& Ortolani 1986 (hereafter GQO86), Brown \\& Wallerstein 1992 (BW92), Drake et al.\\ 1992 (D92), Ivans et al.\\ 1999 (I99), and Smith et al.\\ (2005). These authors have found a range of [Fe/H] between $-$1.3 dex (BW92) and $-$1.05 dex (D92). For further details, see I99 who have summarized the chemical abundances found in previous studies. A study by Norris 1981 (hereafter N81) of 45 RGB stars showed a CN bimodality in M4, e. g., stars of very similar magnitudes and colors have a bimodal distribution of CN band strengths. By analyzing 4 RGB stars, two selected from the CN-weak group and two from the CN-strong one, D92 (and then Drake et al.\\ 1994) found differences in the Na and Al content. I99 found that O is anticorrelated with N, whereas Na and Al abundances are larger in CN-strong stars. Looking at the evolutionary states of these stars, both I99 and Smith \\& Briley 2005 (SB05) did not find any strong correlation between CN band strength and the position on the CMD. More recently, Smith et al.\\ (2005), studying the fluorine abundance in seven RGB stars in M4, found a large variation in $\\phantom{}^{19}$F, which is anticorrelated with the Na and Al abundances. In these previous studies, both the evolutionary scenario and a primordial one have been taken into account in order to explain the light element variations and the CN bimodality in M4.  In this work, we analyze high resolution spectra in order to study chemical abundances for a large sample of M4 RGB stars and compare our results with those of previous studies. In Section 2 we provide an overview of the observations and target sample. The membership criterion used to separate the probable cluster stars is described in Section 3, and the procedure to derive the chemical abundances is in Section 4. We present our results in Section 5, and discuss them in Section 6 and Section 7. Section 8 summarizes the results of this work.  \\begin{figure}[hpbt] \\centering \\includegraphics*[width=9.5cm]{Fig1.ps} \\caption{Distribution of the UVES target stars on the $\\rm {B}$ vs.\\ $\\rm {(B-I)}$ CMD corrected for differential reddening.} \\label{Fig1} \\end{figure} ",
        "conclusions": "We have presented high resolution spectroscopic analysis of 105 RGB stars in the GC M4 from UVES data.  We have found that M4 has an iron content $\\rm {[Fe/H]=-1.07\\pm0.01}$ (the associated error here is the internal error only), and an $\\alpha$ element overabundance $\\rm {[\\alpha/Fe]=+0.39\\pm0.05}$. Si and Mg are more overabundant than the other $\\alpha$ elements, suggesting a primordial overabundance for these elements. Moreover, we also find a slight overabundance of Al. The ${\\rm [Na/Fe]}$ versus [O/Fe] ratios follow the well-known Na-O anticorrelation, signature of proton capture reactions at high temperature. No Mg-Al anticorrelation was found.  We find a strong dichotomy in Na abundance, and show that it must be associated to a CN bimodality. Tab.~\\ref{Tab7} synthesizes the spectroscopic and photometric differences we found for the two groups of stars. Comparing our results to those in the literature on the CN band strength, the CN-strong stars appear to have higher content in Na and Al, and lower O than the CN-weak ones. Apparently, M4 hosts two different stellar populations. This fact is evident from a chemical abundance distribution, and is also confirmed by photometry. In fact, an inspection of the $\\rm {U}$ vs.\\ $\\rm {(U-B)}$ CMD reveals a broadened RGB, and, as shown in our investigation, the two Na groups of stars occupy different positions (have different colors) along the RGB. Our sample of stars is composed of objects with $\\rm {U}$ brighter than $\\sim$16, but, as shown in Fig.~\\ref{Fig11}, the RGB appears to be similarly spread from the level of the SGB to the RGB tip. Since our photometry has been corrected for differential reddening, we conclude that the observed RGB spread at fainter magnitudes must also be due to the metal content dichotomy.  We did not find any evident dependence of the chemical abundance distribution on the evolutionary status (along the RGB) of the target stars, from below the RGB bump to the RGB tip. This is an additional indication that the abundance spread must be primordial.  The two groups of stars we have identified both spectroscopically and photometrically seem to be due to the presence of two distinct populations of stars. The abundance anomalies are very likely due to primordial variations in the chemical content of the material from which M4 stars formed, and not to different evolutionary paths of the present stellar population of M4. This is somehow surprising, because of the relatively small mass of M4 -- it is an order of magnitude smaller than the mass of $\\omega$~Cen, NGC~2808, NGC~1851, NGC~6388, the other clusters in which a multiple stellar generation has been confirmed. Where did the gas which polluted the material for the CN-Na rich stars come from? Has it been ejected from a first generation of stars? How could it stay within the shallow gravitational potential of M4?  The idea that M4 hosts two generations of stars makes the multipopulation phenomenon in GCs even more puzzling than originally thought. It becomes harder and harder to accept the idea that the phenomenon can be totally internal to the cluster, unless this object is what remains of a much larger system (a larger GC or the nucleus of a dwarf galaxy?). Surely, Because its orbit involves frequent passages at high inclination through the Galactic disk, always at distances from the Galactic center smaller than 5 kpc (see Dinescu et al.\\ 1999), M4 must have been strongly affected by tidal shocks, and therefore it might have been much more massive in the past."
    },
    "0808/0808.1078_arXiv.txt": {
        "abstract": "We demonstrate that the observationally inferred rapid onset of star formation after parental molecular clouds have assembled can be achieved by flow-driven cloud formation of atomic gas, using our previous three-dimensional numerical simulations. We post-process these simulations to approximate CO formation, which allows us to investigate the times at which CO becomes abundant relative to the onset of cloud collapse. We find that global gravity in a finite cloud has two crucial effects on cloud evolution.  (a) Lateral collapse (perpendicular to the flows sweeping up the cloud) leads to rapidly increasing column densities above the accumulation from the one-dimensional flow.  This in turn allows fast formation of CO, allowing the molecular cloud to ``appear'' rapidly. (b) Global gravity is required to drive the dense gas to the high pressures necessary to form solar-mass cores, in support of recent analytical models of cloud fragmentation. While the clouds still appear ``supersonically turbulent'', this turbulence is relegated to playing a secondary role, in that it is to some extent a consequence of gravitational forces. ",
        "introduction": "\\label{s:intro}  One motivation for pursuing a picture of molecular cloud formation by large-scale flows was to resolve the ``crossing time problem'', i.e. the fact that the age spreads of young stars are often much shorter than lateral dynamical timescales (\\citealp{1999ApJ...527..285B}; \\citealp{2001ApJ...562..852H}, HBB01). In this picture, dense star-forming clouds are produced as flows -- expanding H{\\small II} regions, stellar wind bubbles, supernova explosions, spiral density waves (e.g. \\citealp{2006MNRAS.365...37B}) or global gravitational instabilities (e.g. \\citealp{2007ApJ...671..374Y}) -- sweep up and condense diffuse interstellar gas.  The crossing time problem is eliminated because information is not being transmitted along the long dimension of the cloud.  A major attraction of the notion of flow-driven molecular cloud formation is that we see several examples of it in the solar neighborhood, such as in Cep OB2 \\citep{1998ApJ...507..241P}, in addition to the well-known cloud and star formation in galactic spiral arms due to gas inflow (e.g. \\citealp{1979ApJ...231..372E,2007ApJ...668.1064E}; see e.g. \\citealp{2003ApJ...599.1157K} and \\citealp{2007MNRAS.376.1747D} for numerical evidence). While it is not always easy to identify specific driving sources in all cases, this is not surprising given the complexity expected as a result of the interaction of neighboring flows. In addition, the presence of molecular clouds well out of the galactic plane (e.g., Orion) clearly suggests the need for some kind of driving. Given the extensive impact that massive stars have on the interstellar medium -- H{\\small II} regions, stellar wind impacts, and ultimately supernova explosions -- it is difficult to see how further creation of new star-forming locales by flows with scales of several to tens of pc (or even kpc, in the case of spiral arms) could be avoided.  Star-forming clouds in this picture become somewhat accidental associations of gas which are not in virial equilibrium, although rough energy equipartition is expected (HBB01), in adequate agreement with observations \\citep{2006MNRAS.372..443B}. Indeed, it is very difficult to prevent global gravitational collapse motions in finite clouds of many Jeans masses \\citep{2004ApJ...616..288B}, consistent with the empirical evidence for rapid star formation following cloud formation (HBB01).  While global gravitational contraction provides a plausible mechanism for assembling protocluster gas and stars and even overall cloud morphology \\citep{2007ApJ...654..988H}, it poses difficulties for local collapse; non-linear perturbations are probably required to avoid sweep-up in overall collapse modes \\citep{2004ApJ...616..288B}. The necessity of producing star-forming clumps through turbulent motions has long been recognized \\citep{1981MNRAS.194..809L}, based on a large number of numerical simulations over the years (see review by \\citealp{2004RvMP...76..125M}). However, the source of this turbulence has been a matter of debate.  The flow-driven cloud formation picture may provide an answer, in that the dynamical instabilities coupled with rapid cooling and thermal instability which naturally result at the shock interface between driving material and swept-up gas generate turbulence and non-linear density fluctuations (\\citealp{2002ApJ...564L..97K}; \\citealp{2002Ap&SS.281...67I}; \\citealp{2005A&A...433....1A}; \\citealp{2005ApJ...633L.113H}; \\citealp{2006ApJ...643..245V}; \\citealp{2006ApJ...648.1052H}; \\citealp{2007ApJ...657..870V}; \\citealp{2007A&A...465..445H}; \\citealp{2007A&A...465..431H}; \\citealp{2008ApJ...674..316H}; \\citealp{2008A&A...486L..43H}). Based on a consideration of timescales, \\citet{2008arXiv0805.0801H} demonstrate that these instabilities do indeed proceed much faster than global gravity, as required.  As simulations of the relevant processes become more detailed, we are increasingly in a position to test the implications of flow-driven molecular cloud and star formation. In this paper we analyze numerical models published elsewhere (\\citealp{2008ApJ...674..316H}, H\\&08) to illustrate some general predictions of the flow-driven picture, and help clarify some confusing issues about molecular cloud and star-formation region lifetimes. ",
        "conclusions": "\\subsection{What would we observe?}\\label{ss:whatobs}  The cloud evolution that we have presented obviously takes far too long to be observable.  For purposes of comparison with observation one can think in terms of observing four different clouds, each having the same evolution, but studied at different epochs.  We also compare at each stage with some observed clouds of similar properties, with the primary goal of checking whether our prescription for CO formation appears reasonable, either with respect to observations or to theoretical treatments for these clouds.  We restrict attention to model Gf1 initially.  ``Cloud 1'' (the top row of Figure \\ref{f:moviex}) would be difficult to detect in molecular gas emission, but could be inferred through absorption line studies and extinction, especially in the small blobs.  This object would be classified as a diffuse H I cloud with a visual extinction $A_V \\sim 0.4$.  ``Cloud 2'' (second row), with an age of 7.6~Myr, would still be classified as a diffuse H I cloud.  It has a small amount of CO, mostly in localized positions; Figure \\ref{f:comasses} demonstrates that very little of the cloud is CO-bearing at this phase. The average extinction through the cloud would be $A_V \\sim 0.6$, with of course some localized higher column density regions; such properties are roughly consistent with local lines of sight to, for example, $\\zeta$~Oph and $\\zeta$~Per, for which van Dishoeck \\& Black's (\\citeyear{1988ApJ...334..771V}) models are consistent with the presence of small amounts of CO.  Here we should emphasize that our prescription for CO formation relies heavily on results such as those of Figure 5a of \\citet{1988ApJ...334..771V}, which indicate photodissociation timescales $>1$~Myr for $A_V \\gtrsim 1$.  Their analysis implies the presence of substantial amounts of H$_2$ - roughly comparable to the amount of H in the $\\zeta$~Per and $\\zeta$~Oph lines of sight.  However, our models would not form any significant amounts of H$_2$ at the stage of ``Cloud 2'', because essentially none of the gas is at densities high enough to form molecular hydrogen rapidly enough (\\citealp{2004ApJ...612..921B}; formation timescales $>1$ Myr at densities $< 10^3$~cm$^{-3}$, corresponding to local free-fall timescales $>1$~Myr; see Figure \\ref{f:taufftime}).  The only way that there would be significant amounts of H$_2$ present at this time (or in ``Cloud 2'') is if it were present in the inflowing gas (e.g., \\citealp{2001MNRAS.327..663P}).  Thus, if anything we may have overestimated the amount of CO present at this stage.  ``Cloud 3'' (third row) now begins to show significant amounts of CO.  According to Figure \\ref{f:comasses}, the mass in CO is now about $500 M_\\odot$ in the small-perturbation model Gf1 and about $10^3 \\msun$ in the large-perturbation model Gf2. However, no self-gravitating cores have formed at this point.  One might regard ``Cloud 3'' as an example of the most massive of the so-called translucent clouds, which are generally observed at high galactic latitude (e.g., \\citealp{1985ApJ...295..402M}; \\citealp{1996ApJS..106..447M}; \\citealp{2003ApJ...592..217Y}). The constraint on galactic latitude probably represents an observational selection effect rather than a complete absence of such objects in the galactic plane.  Observationally, star formation is rare or absent in this type of cloud, consistent with our simulations. For example, in their study of the H{\\small I} filament associated with the molecular complex including MBM 53, 54, and 55, which has a mass of about $1200 \\msun$ and a length of about 30~pc (assuming a distance of 150~pc), \\citet{2003ApJ...592..217Y} found no evidence for any star-forming cores, though there is evidence for one or two young weak-emission T Tauri stars in the region. Note also that translucent clouds are known to have variations in CO abundances of an order of magnitude, even for various positions within the same cloud (\\citealp{1998ApJ...504..290M}, and references therein) which of course is what would be expected for a turbulent, structured translucent cloud like ``Cloud 3''.  ``Cloud 4'', which appears at $t \\sim 12$~Myr, is now a molecular cloud, with a mass associated with CO of order 1200 $\\msun$.  It now contains a small but significant amount of mass which is both gravitationally bound and is at densities corresponding to free-fall timescales less than 1 Myr; thus we consider that it has now begun to form (low-mass) stars.  Finally, ``Cloud 5'' is now a full-fledged molecular cloud, with a mass associated with CO of order 2500 $\\msun$ and dense cores comprising of order at most 25\\% of the mass associated with CO and of order 5\\% of the total mass of the cloud including H I.  The cores are contained within a strongly confined region spatially, due to the global gravitational collapse of the material.  The behavior of model Gf2, with a larger initial perturbation, is very similar, except that high densities are achieved a little earlier, the cloud collapses into a better-defined filament, and the efficiency of massive core/star formation is larger (Figure \\ref{f:comasses}).  Note that the fraction of mass in gravitationally-bound regions does {\\em not} directly correspond to the rate of star formation.  For example, in model Gf2 at the end of the simulation there is about $2000 \\msun$ in gas containing CO, but only $\\sim 10^2 \\msun$ of this material is actually contained in regions with free-fall times of 1 Myr or less (Figure \\ref{f:taufftime}).  Thus in all cases we expect star formation efficiencies would be on the order of a few percent at most, by the end of the simulations.  The sizes and masses of models Gf1 and Gf2 at the final, ``Cloud 4'' stage, are reasonably consistent with those of one of the major filaments in the Taurus molecular cloud \\citep{2002ApJ...578..914H}.  One may even take this a step further and argue that star (core) formation could have only started in these models at times $\\sim 11-12$~Myr, so that at 14.5 Myr the age spread is only of order 2-3 Myr, again reasonably consistent with that of Taurus \\citep{2003ApJ...585..398H}; but one would then need to shut off star formation to maintain the agreement with observation (\\S 4.6).  The above sequence of ``clouds'' is not meant to suggest that all low-mass atomic clouds will eventually become molecular, star-forming regions. In particular, high-latitude clouds may well remain atomic, simply because they cannot sweep up enough mass.   \\begin{figure} \\begin{center} \\includegraphics[width=0.9\\columnwidth]{f6.eps} \\end{center} \\caption{\\label{f:pressprof}Pressure profiles along the inflow axis for model Gf1. Shown are total, dynamical (``kin''), thermal (``int'') and gravitational (``grv'') pressure.} \\end{figure}   \\subsection{Rapid molecular cloud formation}\\label{ss:rapidcloud}  The simulation sequences shown in Figures \\ref{f:moviex} and \\ref{f:moviez} support previous suggestions of rapid cloud and star formation. Specifically, HBB01 argued that the paucity of substantial molecular clouds without star formation meant that collapse of molecular cloud cores must follow soon after cloud formation.  They further suggested that this was the result of the required shielding column density for molecular gas against the dissociating interstellar radiation field, which was similar to that needed for gravitational collapse.  Our simulations demonstrate this behavior, with molecular (CO) gas setting in near $t \\sim 10$~My and core formation soon after, with filamentary gas structure.  A few principles explain why our simulations show the observationally required behavior. One is that the cloud is formed with relatively high densities even in the atomic phase, and thus is ``pre-formed'' with a short free-fall time making rapid collapse possible.  Figure~\\ref{f:pressprof} shows that the pressure profiles along the inflow axis, averaged over the surface area of the cloud for model Gf1, roughly reflect the ram pressures of the inflows until gravity becomes important.  This means that the density within the cloud will approach \\be \\rho \\sim \\rho_i v_i^2 (k_B T /\\mu m_H)^{-1}\\,, \\en where $\\rho_i$ and $v_i$ are the inflow density and velocity, respectively, and $T$ is the temperature.  In terms of densities and velocities $\\tilde{\\rho}_i$ and $\\tilde{v}_i$ relative to the values used in these simulations ($v_i=7.9$~km~s$^{-1}$ and $n_i=3$~cm$^{-1}$, see \\S\\ref{ss:models}), \\be \\rho \\sim 1.3 \\times 10^{-21} T_{30}^{-1} \\, \\tilde{\\rho}_i \\tilde{v}_i^2 {\\rm g}\\, \\mbox{cm}^{-3} \\,, \\en or \\be n \\sim 750 T_{30}^{-1} \\, \\tilde{\\rho}_i \\tilde{v}_i^2 \\, \\mbox{cm}^{-3} \\,, \\en where $T_{30}$ is the temperature measured in units of 30 K, typical of the cold neutral material.  (Note that in the simulations the mean molecular weight is assumed to be unity, as used in these equations.) The free-fall time is then \\be \\tauff = \\left ( {3 \\pi \\over 32 G \\rho } \\right )^{1/2} \\sim 1.9 T_{30}^{1/2} \\, \\tilde{\\rho}_i^{-1/2} \\tilde{v}_i^{-1} {\\rm Myr}\\,. \\en This estimate is in reasonable agreement with Figure \\ref{f:taufftime}, which shows a substantial amount of gas at $\\tauff \\sim 2$~Myr, with higher temperature material at $\\tauff \\sim 6$~Myr. Thus, simply due to ram pressure, moderately cold gas in the cloud is formed at densities such that free-fall timescales are short.  This allows rapid gravitational collapse once sufficient mass has been accumulated.  Another feature of our model is that, as \\citet{2004ApJ...612..921B} showed using a one-dimensional shock model with chemistry, much of the time spent accumulating cloud mass is spent in the atomic phase; such clouds would not be recognized as ``proto'' molecular clouds.  Beyond this, however, local gravitational collapse tends to occur with CO formation because global gravitational collapse is responsible for accelerating the formation of high-column density, highly-shielded gas which can become molecular. This was already foreseen by Bergin \\etal, who suggested that {\\em lateral} collapse under gravity would be responsible for producing ``runaway'' increases in column density, enhancing the rate of molecule formation beyond what can be achieved in a one-dimensional flow. We observe precisely that behavior in our cloud sequence in Figures \\ref{f:moviex} and \\ref{f:moviez} (compare the left-hand column without gravity to the center and right-hand columns with gravity). \\citet{2008arXiv0806.4312D} discuss the formation of H$_2$ clouds in spiral arms and argue that substantial fractions of H$_2$ can be formed rapidly even without self-gravity. This does not contradict our findings, since -- as discussed in \\S\\ref{ss:comasses} -- H$_2$ is expected to form well before CO, at lower column densities. In addition, the densities of the clouds found in these large-scale simulations are generally smaller than typical molecular clouds in the solar neighborhood.  Finally, even if H$_2$ can form without self-gravity, this does not mean that both H$_2$ {\\em and} CO will not form faster with self-gravity.  \\citet{2004ApJ...616..288B} showed that cold sheets were also susceptible to global gravitational collapse.  Ignoring concentrations due to edge effects, which can occur much faster, they found that the timescale for global sheet collapse with outer radius $r$ is \\be t \\sim \\left ( {r \\over G \\Sigma} \\right )^{1/2}\\,. \\en If we ignore the collapse motions and simply allow mass addition at a rate such that $\\Sigma = \\rho_i v_i t$, then we can estimate the distance scales over which global collapse should be present as \\be r \\sim 5.4 t_7^3 \\, \\tilde{\\rho}_i \\tilde{v}_i \\, {\\rm pc}\\,, \\en where $t_7$ is the time measured in units of 10 Myr. Thus it is not surprising that global gravitational collapse is strongly underway at the time of massive core condensation in the simulations.  \\citet{2004ApJ...616..288B} also demonstrated that finite sheets experience highly non-linear gravitational acceleration as a function of position which tends to cause material to pile up near cloud edges.  Our simulations confirm this; even though the resulting clouds are three dimensional instead of being sheets, they are flat enough that the initial elliptical boundary is echoed in the resulting filamentary collapse (an even stronger version of the edge effect is present in \\citet{2007ApJ...657..870V}, whereas we took steps to suppress the strength of the edge effect by reducing the inflow mass addition near the boundaries).  To summarize, reasonable ram pressures of inflowing material result in cold atomic gas with densities sufficiently large that gravitational collapse can be rapid once sufficient material is accumulated.  Substantial turbulent fragmentation is needed to form stellar mass objects.  Global collapse is generally important in driving up column densities and enhancing the formation of filamentary structures through edge effects. Free-fall times, Jeans masses and lengths will decrease, and global collapse will occur faster for higher inflow rates and ram pressures, which probably are needed to make more massive clouds than those of our simulations.  We emphasize that simulations with periodic boundary conditions cannot capture the evolution seen here, in particular the rapid increase in column density within filamentary structure.  Simulations with periodic gravity can only be relevant if started with substantial perturbations, and limited to scales much smaller than the overall sizes of clouds.  \\subsection{``Accelerated'' Star Formation}\\label{ss:accelsf}  Palla \\& Stahler (\\citeyear{2000ApJ...540..255P}, \\citeyear{2002ApJ...581.1194P}) estimated ages for stars in several local star-forming regions, finding that the rate of star formation generally showed an increase up to the last 1 Myr or so.  As Hartmann~(\\citeyear{2002ApJ...578..914H}, \\citeyear{2003ApJ...585..398H}) pointed out, it is implausible that all of these star-forming regions should be coordinated in their evolution; it must be that the typical lifetimes of star-forming regions are at most a few Myr. Hartmann~(\\citeyear{2002ApJ...578..914H}, \\citeyear{2003ApJ...585..398H}) also pointed out some problems related to neglect of observational errors, logarithmic vs. linear binning in age, and likely isochrone or birthline problems which exaggerate the age spreads of these regions.  Nevertheless, there is evidence that at a small fraction of the stars associated with these areas are a few Myr older than the bulk of the population.  Accelerated star formation is expected in a trivial sense, since starting from a rate of zero star formation to a finite star formation rate mathematically requires an acceleration.  However, it also occurs naturally in our evolutionary model of molecular cloud formation.  The Gf1 and Gf2 simulations shown in Figures \\ref{f:moviex} and \\ref{f:moviez} show that a few dense concentrations show up by 10 Myr or so (see also the core mass history, Fig.~5 in H\\&08). One can easily infer that, if we had sufficient resolution, we would observe a few stars being formed from the very densest of the dense concentrations first; later on, as global and local collapse proceeds, more and more stars would form.  Our simulations emphasize that the global age spread in a star-forming region is an upper bound to the timescales of local collapse, as compressions will never be perfectly coordinated in time across any given cloud.  \\subsection{The Onset of Star Formation}\\label{ss:onsetsf} In our model, turbulence generated through the cloud formation process provides a mechanism to produce thermal and dynamical fragmentation. While this fragmentation is highly efficient in generating cold high-density cloudlets, the ``pre-formation'' of such cloudlets is limited by the minimum Jeans mass achievable. With $c_s$ denoting the (isothermal) sound speed, we have \\be M_J = \\left ( {\\pi c_s^2 \\over G } \\right )^{3/2} \\rho^{-1/2} \\sim 570 T_{30}^2 \\tilde{P}_{ram}^{-1/2} \\, \\msun\\,, \\label{e:mjeans} \\en where $\\tilde{P}_{ram} = \\tilde{\\rho}_i \\tilde{v}_i^2$, again in units of our simulations. To form e.g. solar-mass objects, higher pressures are needed. Since invoking ``supersonic turbulence'' as a fragmentation mechanism in its own right without specifying its physical source is somewhat unsatisfactory, the only remaining free energy source leading to a self-consistent picture of the formation of pre-stellar cores is gravity (e.g. \\citealp{1992ApJ...395..140B}; \\citealp{2008MNRAS.385..181F}).  We have already seen that gravity helps along the ``CO formation'' in our clouds (Figs.~\\ref{f:moviex} and \\ref{f:moviez}), and also that unless the simulations include gravity, they are not going to form the high-density cores with freefall times $\\tauff<1$~Myr (Fig.~\\ref{f:taufftime}). Figure~\\ref{f:moviep} demonstrates that it is gravity which eventually leads to the high pressures required for fragmentation into small masses. It shows the mass-weighted pressure histograms of models Hf1 and Gf1 in the same time sequence as in e.g. Figure~\\ref{f:moviex}. At $t=5.3$~Myr, the distributions are indistinguishable. With increasing time, gravity drives more and more mass to higher pressures, until the pressure distribution follows the Jeans mass (eq.~[\\ref{e:mjeans}]) at the lowest temperature threshold (indicated by the blue dashed line). In contrast to model Gf1, the pressure distribution of model Hf1 stays more or less constant with time.   \\begin{figure} \\begin{center} \\includegraphics[width=0.9\\columnwidth]{f7c.eps} \\end{center} \\caption{\\label{f:moviep}Mass (in $M_\\odot$) within logarithmic pressure bins for models Hf1 and Gf1, at the same times as in e.g. Figure~\\ref{f:moviex}. Colors indicate temperature thresholds. The left vertical black dashed line stands for the thermal pressure of the inflows, while its right-hand counterpart shows the ram pressure of the inflows. The slanted lines in the temperature color code indicate the Jeans mass (eq.~\\ref{e:mjeans}). With increasing time, gravity drives more and more mass to higher pressures, allowing the fragmentation into successively smaller structures \\citep{2008MNRAS.385..181F,2008ASPC..387..240V}.} \\end{figure}   \\subsection{Turbulent Support}\\label{sss:turbsupp} It has been claimed that the loss of turbulent support (e.g. \\citealp{2007ApJ...666..281H}) would lead to an (accelerated) star formation in molecular clouds \\citep{2000ApJ...540..255P,2002ApJ...581.1194P}. A closer look at the energy budget (Fig.~\\ref{f:energies}) in our model clouds however reveals that in the context of flow-driven cloud formation, it is actually the build-up of mass and thus the deepening of the gravitational potential well which leads to collapse, but not the loss of turbulence (see also Fig.~8 of \\citealp{2007ApJ...657..870V}). The energy densities were calculated using only gas at $T<50$~K, approximately corresponding to the molecular (Gf1, Gf2), but at least cold, dense gas (Hf1). The gravitational energies grow faster with time than the kinetic energies drop -- and eventually, once the molecular cloud has ``appeared'' (see Figs.~\\ref{f:moviex} and \\ref{f:moviez}), the kinetic energies start to rise, indicating that the turbulent ``motions'' at that stage are driven to some extent by gravitational collapse (see also \\citealp{2006MNRAS.372..443B}; \\citealp{2007ApJ...657..870V}; \\citealp{2008MNRAS.385..181F}; \\citealp{2008ASPC..387..240V} for a recent discussion). In other words, while the turbulence might decay slightly during the formation of the model cloud (in the atomic phase), it is {\\em increasing} -- driven by gravity -- during the molecular phase. One would expect stellar feedback to even strengthen this trend, if anything. Figure~\\ref{f:losvd} demonstrates that despite the gravitational dominance, the linewidths still can look ``turbulent''. The linewidths shown were calculated for all gas at $T<50$~K, along a single line-of-sight of 16 resolution elements width through the center of the cloud.  We note that the kinetic energies shown in Figure~\\ref{f:energies} are subvirial by a factor of up to $3$ at late times, and that the linewidths (Fig.~\\ref{f:losvd}) are on the smallish side. We believe this is due to the combination of three issues. First, the head-on colliding flows do not allow for large-scale shearing motions. Those would be expected e.g. in molecular clouds forming in spiral arms. Numerical models of this process \\citep{2006MNRAS.365...37B} including a shear component reproduce turbulent linewidth-size relations \\citep{1981MNRAS.194..809L}. Second, non-uniform inflows will increase the level of turbulence in the resulting cloud, as demonstrated by \\citet{2007MNRAS.374.1115D}. They found the slope of the linewidth-size relation to flatten with increasing filling factors of the clumpy inflows. Third, the dynamical range for gravitational collapse in our models is somewhat limited, possibly suppressing turbulent motions once the cloud has collapsed.   \\begin{figure} \\includegraphics[width=\\columnwidth]{f8.eps} \\caption{\\label{f:energies}Energy densities (or pressures) of the three isolated cloud models Hf1, Gf1 and Gf2, for gas at $T<50$~K, approximately corresponding to the molecular gas. The internal energy (black lines) stays constant, while the gravitational energy (green lines) increases with time, offsetting the slight decrease in kinetic energy (red lines).} \\end{figure}   \\begin{figure*} \\includegraphics[width=\\textwidth]{f9c.eps} \\caption{\\label{f:losvd}Line-of-sight velocity distribution for models HF1 (top) and Gf1 (bottom), for a single line of sight parallel to the shorter ({\\em left}) and longer ({\\em right}) axis of the cloud perpendicular to the inflow. Three times as in Figure~\\ref{f:moviex}.} \\end{figure*}   We take the discussion of the perceived loss of turbulent support as an opportunity to demonstrate how the star formation in our models (strictly speaking, we are only forming cores) would be represented by the normalized ``star formation rate'' of \\citet{2005ApJ...630..250K} and \\citet{2007ApJ...654..304K}, \\begin{equation} SFR_{ff}\\equiv \\frac{\\dot{M}}{M(>n)}\\,\\tauff.\\label{e:bogussfr} \\end{equation} Here, $M(>n)$ is the mass of the cloud above a density threshold $n$, $\\dot{M}$ is the star formation rate, and $\\tauff$ is the freefall time at a given density. The choice of this density obviously will affect the results. Figure~\\ref{f:krumholztan} shows the $SFR_{ff}$ as defined above, for various times against the threshold density $n$ (model Gf1). As a proxy for the star formation rate $\\dot{M}$ we use the time derivative of the mass accretion history (see e.g. symbols in Fig.~\\ref{f:comasses}). This includes the formation of new cores as well as the mass accretion onto already formed cores. It is also consistent with the estimates of $\\dot{M}$ from the simulations discussed by \\citet{2007ApJ...654..304K}.   \\begin{figure} \\includegraphics[width=\\columnwidth]{f10.eps} \\caption{\\label{f:krumholztan}The star formation rate normalized by the free fall time $SFR_{ff}$ (eq.~[\\ref{e:bogussfr}]). The plot has been generated following the recipe given by \\citet{2007ApJ...654..304K}, i.e. the free-fall time is computed for the mean ({\\em left}) and median ({\\em right}) density above the density threshold given on the $x$-axis. Symbol sizes indicate simulation times.} \\end{figure}   At early times ($t=11.4, 12.2$~Myr), the models are consistent with the analytical predictions (see Fig.~5 of \\citealp{2007ApJ...654..304K}): the $SFR_{ff}$ ranges around a few percent. Obviously, it is a matter of interpretation whether this low percentage is to be seen as ``slow'' star formation: especially at the lower density thresholds, the freefall time $\\tau_{ff}$ is {\\em not} representative of the local conditions under which the ``stars'' form. At later times, several effects are discernible. For low density thresholds (i.e. considering the ``whole'' cloud, see also Fig.~\\ref{f:taufftime}), still $SFR_{ff}\\lesssim 10$\\%. For higher density thresholds (i.e. selecting for the regions that are locally collapsing in the simulation), the $SFR_{ff}$ increases but stays short of $\\approx 50$\\%. This is not surprising, since (a) these regions are not resolved, i.e. they would not convert all their mass into stars so that the ``real'' $SFR_{ff}$ would be substantially lower and (b) our models do not include feedback. Moreover, the models discussed here are an extreme realization of a range of scenarios whose other extreme would be a perfect shear flow (Williams et al., in preparation). With a larger component of shear or angular momentum, the $SFR_{ff}$ could be reduced even further.  \\subsection{Other limitations} In addition to limited resolution, our models do not include stellar energy input, which is the the most plausible mechanism to reverse the collapse of our self-gravitating molecular clouds.  It is difficult to see how stellar energy input can exactly balance the non-linear variation of gravity with position within the cloud; there is more parameter space for either expansion or contraction. We note that even at the end of our simulations, most of the mass of the cloud, including that of the molecular (CO) regions, has long free-fall times and low densities.  The low filling factors indicated by the $SFR_{ff}$ at $n\\approx 100$~cm$^{-3}$ and by Figure~\\ref{f:taufftime} suggest that stellar feedback can in principle blow away most of the cloud mass and thus keep overall star formation efficiencies low (see also \\citealp{2000ApJ...530..277E}). While this remains an important issue of our scenario which needs further investigation, it should also be pointed out that to our knowledge there are no simulations of long-lived clouds that are {\\em not} computed with periodic boundary conditions, the adoption of which begs the question.  Magnetic fields could in principle suppress or affect the fragmentation (e.g. \\citealp{1965ApJ...142..531F} for the thermal instability; \\citealp{2007ApJ...665..445H} for the dynamical instabilities in colliding flows), although this capability seems to depend strongly on the number of dimensions considered (e.g. \\citealp{2008arXiv0801.0486I}). Also, the sheer accumulation of mass along fieldlines might render the fields less important than commonly thought, as demonstrated in numerical models by \\citet{2008A&A...486L..43H}.    Making a simple but reasonable approximation for the formation of CO from atomic gas, we have shown that the formation of molecular clouds from large-scale flows can be tied closely to the formation of self-gravitating structures.  Our numerical simulations support many parts of the rapid cloud and star formation scenario, in particular the rapid appearance of CO in self-gravitating clouds, quickly followed by the formation of dense cores with free-fall timescales of order 1 Myr, and that global gravitational collapse plays an important role in cloud evolution. The simulations reproduce expected low star formation rates at early times. For the long-term evolution of the star formation history and the molecular cloud itself, the issue of stellar energy input (feedback) needs to be addressed to develop a more comprehensive picture of star formation in the solar neighborhood."
    },
    "0808/0808.1880_arXiv.txt": {
        "abstract": "We examine phantom dark energy models produced by a field with a negative kinetic term and a potential that satisfies the slow roll conditions:  $[(1/V)(dV/d\\phi)]^2 \\ll 1$ and $(1/V)(d^2 V/d\\phi^2) \\ll 1$. Such models provide a natural mechanism to produce an equation of state parameter, $w$, slightly less than $-1$ at present. Using techniques previously applied to quintessence, we show that in this limit, all such phantom models converge to a single expression for $w(a)$, which is a function only of the present-day values of $\\Omega_\\phi$ and $w$. This expression is identical to the corresponding behavior of $w(a)$ for quintessence models in the same limit. At redshifts $z \\alt 1$, this limiting behavior is well fit by the linear parametrization, $w=w_0 + w_a(1-a)$, with $w_a \\approx -1.5(1+w_0)$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.4141_arXiv.txt": {
        "abstract": "We study the stability of standing shock waves in advection-dominated accretion flows into a Schwarzschild black hole by 2D general relativistic hydrodynamic simulations as well as linear analysis in the equatorial plane. We demonstrate that the accretion shock is stable against axisymmetric perturbations but becomes unstable to non-axisymmetric perturbations. The results of dynamical simulations show good agreement with linear analysis on the stability, oscillation and growing time scales. The comparison of different wave-travel times with the growth time scales of the instability suggests that the instability is likely to be of the Papaloizou-Pringle type, induced by the repeated propagations of acoustic waves. However, the wavelengths of perturbations are too long to clearly define the reflection point. By analyzing the non-linear phase in the dynamical simulations, it is shown that quadratic mode couplings precede the non-linear saturation. It is also found that not only short-term random fluctuations by turbulent motions but also quasi periodic oscillations take place on longer time scales in the non-linear phase. We give some possible implications of the instability for quasi periodic oscillations (QPOs) and the central engine for gamma ray bursts (GRBs). ",
        "introduction": "Accretion flows imbedding a shock wave have attracted much attention of researchers. Hydrodynamic instabilities of shocked accretion flows may explain the time variability of emissions from many black hole candidates, since a shock wave is a promising mechanism of transforming the gravitational energy into radiation. The possibility that a shock exists in black hole accretion disks was first suggested by \\citet{Hawley84a,Hawley84b} even though they did not make clear the essential condition for its existence. The possible structure of the shocked accretion flows in equatorial plane was described by \\citet{Fukue87}, who suggested the importance of the multiplicity of critical points for the existence of the standing shock. Multiple critical points exist only for appropriate values of the injection parameters such as the specific angular momentum and Bernoulli constant. Note also that there are generally  two possible shock locations, which are referred to as the inner and outer shocks.   The stability of the standing shock wave in the accretion flows has been also investigated by many authors both analytically and numerically. \\citet{nak94,nak95} showed by linear analysis in the equatorial plane that if the post-shock matter are accelerated, the flow is unstable against radial perturbations, which is true for both Newtonian or general relativistic dynamics. As a result of this theorem, we know for the black hole accretion that the inner shock becomes generally unstable against radial perturbations and that non-rotating steady accretion flows to a black hole cannot have a stable standing shock wave in it. These features were also observed in 1D axisymmetric simulations for a pseudo-Newtonian potential \\citep{cha93,nob94}.   Recently, \\citet{fog01,fog02} pointed out by linear analysis that the outer shock wave, which is stable against radial perturbations, is in fact unstable to non-radial axisymmetric perturbations. He argued that the advective-acoustic cycle could be responsible for the instability. In this mechanism the velocity and entropy fluctuations initially generated at the shock are advected inwards, producing pressure perturbations, which then propagate outwards, reach the shock and generate entropy and velocity fluctuations there, thus repeating the cycle with an increased amplitude. The same instability appears to be working in the accretion flows onto a nascent proto neutron star in the supernova core, in which large scale oscillations of the shock wave of $\\ell=1$ nature are observed \\citep{blo03}, where $\\ell$ stands for the index in the spherical harmonic functions $Y^{\\ell}_{m}$. Although the mechanism is still controversial \\citep{ohn05,fog06sn,blomeza06,Lami07}, the so-called standing accretion shock instability or SASI is currently attracting much attention as a promising explanation of asymmetric explosion of supernova as well as young pulsars' proper motions \\citep{sch04,sch06} and spins \\citep{blomeza07,Iwapre}.     As for the stability of the shock wave against non-axisymmetric perturbations, \\citet{molteni99} did 2D simulations of an adiabatic shocked accretion flow by using the pseudo-Newtonian potential and found a non-axisymmetric instability. They showed that the shock instability is saturated at a low level and a new quasi-steady asymmetric configuration is realized. To investigate the mechanism of this instability, \\citet{fog2003,Gu2005} performed linear analysis for both isothermal and adiabatic flows. They concluded that the instability seems to be of the Papaloizou-Pringle type and the repeated propagations of acoustic waves between the corotation radius and the shock surface are a driving mechanism. These conclusions are based on the WKB approximation and the comparison between the growth rate and the acoustic cycle period.   It is also noted that \\citet{yamasaki2007} investigated the linear stability of the shocked accretion flows onto the proto neutron star against non-axisymmetric perturbations in the equatorial plane. They demonstrated that the counter-rotating spiral modes are significantly damped, whereas the growth rate of the corotating modes is increased by rotation. They also claimed that the instability is not of the Papaloizou-Pringle type, since the stability is not affected by the presence or absence of the corotation point. Instead, they suggested the advective-acoustic cycle based on the WKB analysis. The purely acoustic cycles was found to be stable.   As mentioned above, although many efforts have been made for clarifying the non-radial instability, the complete understanding of the mechanism is still elusive. It is noted that the unperturbed accretion flows should be treated properly, since they strongly affect the instability. In black hole accretions, the gravitation is one of the main factors to determine the flow features such as sonic points and shock locations. So far, however, the shock stability in the accretion flows to black holes has been investigated under the Newtonian or pseudo-Newtonian approximation and there has been no fully GR treatment. As shown later, the range of injection parameters that allows the existence of a standing shock wave is changed substantially when GR is fully taken into account.   In this paper, we investigate fully general relativistically the stability of the shock in the advection-dominated accretion flows into a Schwarzschild black hole by using both linear analysis and non-linear dynamical simulations. In so doing, we consider only the equatorial plane, assuming that the $\\theta$-component of four-velocity ($u^{\\theta}$) and all the $\\theta$-derivatives are vanishing. We use the spherical coordinates ($r,\\theta,\\phi$) in the following. The evolution of the metric is not taken into account, which is justified if the mass of accretion flow is much smaller than the black hole mass. We show that the shock is indeed unstable against non-axisymmetric perturbations and a spiral arm structure is formed as the instability grows. We discuss the instability mechanisms comparing various time scales. Finally, we mention possible implications of our findings for gamma-ray bursts (GRBs) and black hole quasi-periodic oscillations (QPOs).   This paper is organized as follows. In section 2, we describe the steady axisymmetric accretion flows with a shock in them. In section 3, we present the formulation of linear analysis. The numerical method for dynamical simulations is explained in section 4. The main numerical results and their analyses are given in section 5. The implication for GRBs and Black Hole QPOs are mentioned in section 6. Finally we summarize and conclude the paper in section 7. ",
        "conclusions": "We have investigated the non-axisymmetric shock instability in the accretion disk around Schwarzschild black holes, employing the fully general relativistic hydrodynamic simulations as well as the linear analysis. Both the linear and non-linear phases have been analyzed in detail. We have also given some possible implications for astrophysically interesting phenomena such as Black Hole QPOs and GRBs.   The main findings in the present work are as follows:   (a) The standing shock is generally unstable against non-axisymmetric perturbations, and a spiral arm structure is formed as a result of the growth of instability. It is typically one-armed, implying that the dominant mode in the non-linear phase is the $m=1$ mode.   (b) In the linear phase, the dynamical simulations are in good agreement with the linear analysis in such features as stability, oscillation and growth time scales. The progressive modes, in which the deformed shock pattern rotates in the same direction as the unperturbed flow, are unstable and the retrogressive modes are stable. This is consistent with the previous works.   (c) In the non-linear phase, various modes are produced by non-linear couplings, which are mainly of quadratic nature, and the amplitudes are saturated. The axisymmetric mode is also induced by the non-axisymmetric instability, and the shock radius oscillates with large amplitudes. The oscillation periods become slightly longer than in the linear analysis because of larger shock radii.   (d) Even though strong perturbations are added initially, the shock remains to exist. Hence the disk plus shock system is quite robust in this sense.   (e) The comparison of various cycle time scales with the linear growth times seems to support the claim that the instability is induced by the acoustic-acoustic cycle, although the inner reflection point is not identified unambiguously. It is important to note in this respect that the wavelength of perturbations is longer than the scale height, which does not allow the WKB approximation.   (f) The Black Hole SASI found by \\citet{molteni99} may be a promising candidate for the sources of the Black Hole QPOs and fluctuations in GRB jets.    In the present study, we have also found that the non-axisymmetric instability is sensitive to the structure of the unperturbed steady flow. The general relativity is important in this respect. It should be stressed that the injection parameters that allow the existence of a standing shock wave are different between the GR and pseudo-Newtonian treatments. In fact, we have found by the direct comparison that the maximum specific angular momentum for the existence of multiple sonic points is different by more than 60\\% for the Bernoulli constant $E \\le 1.003$.  Note also that the general relativity is indispensable in discussing the accretion into a Kerr black hole, since the frame-dragging will play an important role. This is currently undertaken \\citep{nak08v2}. For more detailed comparison with observations, it is necessary to include the cooling and heating for GRBs case, and the magnetic field and viscosity for Black Hole QPOs. Last but not least, the discussed simulations including the polar dimension are inevitable."
    },
    "0808/0808.1849_arXiv.txt": {
        "abstract": "{Applying photometric catalogs to the study of the population of the Galaxy is obscured by the impossibility to map directly photometric colors into astrophysical parameters. Most of all-sky catalogs like ASCC or 2MASS are based upon broad-band photometric systems, and the use of broad photometric bands complicates the determination of the astrophysical parameters for individual stars.} {This paper presents an algorithm for determining stellar astrophysical parameters (effective temperature, gravity and metallicity) from broad-band photometry even in the presence of interstellar reddening. This method suits the combination of narrow bands as well.} {We applied the method of interval-cluster analysis to finding stellar astrophysical parameters based on the newest Kurucz models calibrated with the use of a compiled catalog of stellar parameters.} {Our new method of determining astrophysical parameters allows all possible solutions to be located in the effective temperature-gravity-metallicity space for the star and selection of the most probable solution.} {} ",
        "introduction": "Estimating the astrophysical parameters of stars is one of the core problems in astrophysics and the key solving a number of astrophysical questions. Indeed, the knowledge of effective temperature, luminosity class (or gravity at the surface of the star), and metallicity of the star allows the distance to the star, the age of the star, and the extinction in the interstellar medium between the star and the observer to be estimated.  The number of stars with detailed spectral information is very limited compared to the wealth of data in photometry. The problem of determining astrophysical parameters with photometry is usually solved with the help of the narrow-band photometric systems (for instance, the Str\\\"omgren photometric system) that provide quite good precision for a temperature, gravity, and metallicity. For example, Nordstr\\\"om (\\cite{Nordstroem}) estimates the metallicity from Str\\\"omgren photometry with the final average 0.12 dex dispersion around the spectroscopic values. The problem becomes much more complicated in the case of broad-band photometry where all fine features in the photometric bands (tracers of the differences in metallicity and especially in gravity) are effectively smoothed. The estimation of $T_{eff}$, $\\log g$, and [M/H] is crucial for the study of the Galaxy on the basis of broad-band all-sky surveys like 2MASS or the broad-band deep surveys like SDSS/SEGUE or UKIDSS.  The history of stellar parameter determination from broad-band photometry was started by Johnson in the middle of the 50s of the last century (see, for example, Johnson~\\cite{Johnson53} where the author defined extinction-free colors for the UBV photometric system).  Later, some broad-band photometric systems showed the usefulness of extinction-free colors for determining of temperatures and gravities with some assumptions about the metallicity (for example, the Vilnus photometric system). The use of extinction-free colors is discussed in detail in Strai\\v{z}ys (\\cite{Straizys}).  A major problem of using extinction-free colors is how to determine extinction-free colors themselves  from the observed ones. Indeed, as we show below, to calculate the extinction-free indices and to use them for the estimation of astrophysical parameters we have to accept some initial guesses about the estimated astrophysical parameters of the star {\\it a priori}.  In this study we minimize  the amount of information to be assumed  {\\it a priori} for the estimation of astrophysical parameters of the star. We will not make guesses about the astrophysical parameters of the star (temperature, gravity or metallicity) based on the position of the star, e.g. relative to the galactic plane.  We do not discuss the dependence of the estimated parameter on the synthetic model grid adopted in this study and do not compare different synthetic model grids. The estimated parameters will always depend on the consistency of the synthetic models used in the study. However, in the case of Johnson B,V,J,H, and $K_S$ photometry, we calibrated the parameters obtained from the model fluxes by our method with published parameters from different authors. ",
        "conclusions": "\\begin{table}[b] \\caption{Reason for the failure of the parameter estimation} \\label{fail_reas} \\begin{tabular}{|c|c|c|c|} \\hline \\multicolumn{4}{|c|}{$\\log T_{eff}$} \\\\ \\hline name  &  $\\log T_{eff}^{sp}$ & $\\log T_{eff}^{ph}$ &  \\\\ \\hline HD 18078 & 4.003 & 4.654 & A0p \\\\ HD 96446 & 4.401 & 4.635 & B2IIIp \\\\ \\hline \\multicolumn{4}{|c|}{$\\log g$} \\\\ \\hline name  &  $\\log g^{sp}$ & $\\log g^{ph}$ &  \\\\ \\hline BD +30 2611 & 1.13 & 3.20 & \\\\ HD 101602   & 1.40 & 4.88 & UZ Cen, \\\\ &      &      & Cep-type variable\\\\ \\hline \\multicolumn{4}{|c|}{[M/H]} \\\\ \\hline name  &  $[M/H]^{sp}$ & $[M/H]^{ph}$ &  \\\\ \\hline BD +30 2611 & -1.40 & -0.33 & \\\\ HD 101602   &  0.14 & -2.15 & UZ Cen, \\\\ &       &       & Cep-type variable \\\\ HD 154338   &  0.0  & -1.33 & V991 Sco, \\\\ &        &      & RS CVn-type variable\\\\ HD 3940     &  -0.39 & -1.85 & V755 Cas, \\\\ &        &       & Algol-type \\\\ HD 6582     & -0.73  & -1.93 & SB\\\\ \\hline \\end{tabular} \\end{table}     We developed and tested a new method determining the astrophysical parameters from broad-band photometry. This method does not rely on a predefined model for the stellar population, but requires assuming the law for the extinction in the direction of the estimated star. In principle, the  extinction law can be determined by the method itself with the use of stars with a unique solution for their astrophysical parameters, which is independent of the size of extinction.  We used an interval-cluster analysis as the mathematical core for the method. For a given input photometric system (here Johnson B,V, and 2MASS J, H, and $K_S$), this method allows us to select the regions in the parameter space where a (unique) solution is possible and to identify regions in the parameter space where no solutions are possible due to the complicated topology of the color and parameter spaces. The astrophysical parameters determined with the use of this method are defined in the system of the selected models (synthetic or empirical) and suffer from all the systematic errors which are inherent in these models.  We have compiled a catalog of stars with  astrophysical parameters known from spectroscopic observations to calibrate the Kurucz models and the estimated parameters.   In the Johnson B,V, and 2MASS J, H, and $K_S$ photometric system and with mean errors of input colors better than 0.01 mag the best results are achieved for main sequence G-K stars (most preferable region, G6 to K9). In this case the parameters are accurate to $\\sigma_{\\log T_{eff}} < 0.05$, $\\sigma_{\\log g} < 0.3$, and $\\sigma_{[M/H]} < 0.3$. If the accuracy of the color measurements drops to 0.05 mag, the most preferable region with the accuracy defined above ($\\sigma_{\\log T_{eff}} < 0.05$, $\\sigma_{\\log g} < 0.3$, and  $\\sigma_{[M/H]} < 0.3$) shrinks to main sequence K6-K9 stars; meanwhile, the accuracy for the rest of the most preferable region drops dramatically, especially for gravity. Nevertheless, even for 0.05 mag the accuracy of the determined parameters for the rest of most preferable regions (main sequence G6 to K5) stays within $\\sigma_{\\log T_{eff}} < 0.1$, $\\sigma_{\\log g} < 2.0$, and $\\sigma_{[M/H]} < 0.4$. The definition of ``the most preferable regions'' depends on the combination of photometric bands, desired precision for the result and the precision of the input photometry.  We find that the biggest problem in determining astrophysical parameters is not to determine a solution but to distinguish between possible solutions that correspond to the input colors. We propose and tested a simple criterium for the selection of a solution based on the estimated extinction. Final tests with the compiled catalog proves the reliability and effectiveness of the method by comparing spectroscopic (literature) and photometric (our method) determinations of parameters.  The proposed method was  initially designed to search for sparse stellar groups in deep sky surveys, but it also suits the study of populations in the Galaxy without a predefined model for the distribution of stars. It is possible to use the method to determine an extinction law."
    },
    "0808/0808.2883_arXiv.txt": {
        "abstract": "Sgr A* is the closest massive black hole and can be observed with the highest angular resolution. Nevertheless, our current understanding of the accretion process in this source is very poor. The inflow is almost certainly of low radiative efficiency and it is accompanied by a strong outflow and the flow is strongly variable but the details of the dynamics are unknown. Even the amount of angular momentum in the flow is an open question. Here we argue that low angular momentum scenario is better suited to explain the flow variability. We present a new hybrid model which describes such a flow and consists of an outer spherically symmetric Bondi flow and an inner axially symmetric flow described through MHD simulations. The assumed angular momentum of the matter is low, i.e. the corresponding circularization radius in the equatorial plane of the flow is just above the innermost stable circular orbit in pseudo-Newtonian potential. We compare the radiation spectrum from such a flow to the broad band observational data for Sgr A*. ",
        "introduction": "Accretion of matter onto a black is a complex non-stationary phenomenon, and an inflow of material is accompanied by an outflow. It is widely believed that magnetic fields play a key role but the details of accretion/outflow process are not understood.  Sgr A* is the closest massive black hole and therefore an obvious target for any future high spatial resolution observations. The Schwarzschild radius for the black hole mass of $3 \\times 10^6 M_{\\odot}$ at the distance of 8 kpc corresponds to $\\sim 8 $ micro arc seconds. The next massive black hole (M87, distance of 16 Mpc, black hole mass $3 \\times 10^9 M_{\\odot}$) is by a factor of 2 smaller with respect to the angular size.  At the same time, Sgr A* is an example of a source with extremely low value of the Eddington ratio. It is by no means a spectacularly bright source, and the point like source emission in radio, IR and X-ray band was only detected and monitored in recent years. However, studying weakly active galaxies is important for two reasons: (i) galaxies spend most of their time in such phase (ii) the outflowing material may carry significant amount of kinetic energy so the feedback between the central activity and the galaxy itself may be also significant during the long inactive phase.  This studying Sgr A* is important but not easy. Models which apply to bright, high Eddington ratio sources like quasars may not necessarily apply to low level activity. We argue that accretion in Sgr A* is most likely an example of a low angular momentum flow, with circularization radius at a few or a few tenths of Schwarzschild radii, since such a flow is naturally variable and accompanied by a strong outflow so only a small fraction of the material actually reaches the horizon of a black hole. ",
        "conclusions": "The Sgr A* is certainly an extremely interesting object and a promising laboratory for the study of the accretion flow onto a black hole due to its unique large angular size. On the other hand, the accretion process in this source is particularly complex which is a characteristic property of the Low Luminosity AGN. The accretion/outflow pattern is highly non-stationary and it is likely to depend very strongly also on the outer/initial boundary conditions of the flow. This means that even with an advancement in the theory we cannot uniquely reconstruct the flow without direct observational information based on time-dependent imaging of the source with high spatial resolution. Thus Sgr A* is a much more difficult case to understand than bright quasars where a fairly standard Shakura-Sunyaev disk well accounts for most of the source properties.  This pessimistic aspect does not mean that there is nothing to be done in the meantime. As for the observations, it is important to continue the multi-wavelength monitoring as this gives direct constraints to the clumpiness of the flow as discussed by Baganoff and Eckart (this proceedings). It is also important to attempt to determine the density, temperature and emissivity map as close to the black hole as possible since the flow is not expected to be spherically symmetric at the Bondi radius. Instead, it is rather likely to show some axial symmetry, and the general level of the departure from a spherical symmetry is likely to measure the net angular momentum of the material in the circumnuclear region. On the theoretical side, simulations covering both the inner region and the outer resolved region are certainly needed, and the description of the outer region should contain the information about the motion of the donor stars as in Cuadra et al. (2008) as well as possible variations (time-dependence and clumpiness) of the stellar winds from individual stars. There are also many specific aspects of modeling which should be addressed. For example, the derivation of the electron temperature from the dynamical simulations in the inner MHD region require significant smoothing of the local energy density if the electron advection is also to be included (see Mo\\' scibrodzka et al. 2007); otherwise, in some regions there is no self-consistent solution for the electron temperature. This means that some important physics is still missing in the models.  So far the amount of angular momentum and its distribution in the accreting matter is an open question to be addressed in the future but in our opinion the models with low angular momentum content are most promising. In this case a kind of strongly variable inner torus forms in a natural way, and the strong short lasting outbursts are easily produced in such scenario. The torus itself, pushed by the continuously inflowing material, shows a kind of oscillations (Mo\\' scibrodzka \\& Proga 2008, in preparation) which may be of interest in the context of the quasi-periodic oscillations possibly seen in Sgr A* (in X-ray band: Genzel et al. 2003b, Aschenbach et al. 2004; in IR band: Meyer et al. 2006, Eckart et al. 2006, Trippe et al. 2007) although torus oscillations are of somewhat lower frequencies than $\\sim 17$ minutes timescale seen in the data.  \\subsection*"
    },
    "0808/0808.2551_arXiv.txt": {
        "abstract": "We in the paper study the metric perturbations generated in a bouncing universe driven by the Quintom matter. Firstly, we review the background evolution of Quintom Bounce and the power spectrum of scalar perturbations. Secondly, we study the non-Gaussianity of curvature perturbations and then calculate the tensor perturbations of the model. ",
        "introduction": "Inflation was invented to resolve problems existing in hot Big Bang cosmology, such as flatness, horizon, primordial monopole problem\\cite{Guth:1980zm} (for some early attempts see Refs. \\cite{Starob}). After more than twenty years' development, one has obtained a deep sight at this theory. However, it is still puzzled by an initial singularity which exists in usual inflationary models\\cite{Borde}.  It has been suggested that, bouncing cosmology, which requires our universe initially experience a contracting phase before the hot Big Bang expansion, could provide a solution to the problem of the initial singularity. For years, models of a bouncing universe have received a lot of attentions and there have been a number of works on constructing this scenario. For example, there are models with singular bounce such as the Pre-Big-Bang \\cite{PBB} and cyclic/Ekpyrotic \\cite{Ekp}; also in Refs. \\cite{Bojowald:2001xe,Brustein:1997cv,Biswas:2005qr} some non-singular bounce models were considered where the gravitation action was modified by higher order corrections; and, see Ref. \\cite{Novello:2008ra} for a recent review on various models of bouncing cosmology.  Recently, a new bounce model has been proposed \\cite{Cai:2007qw}, dubbed Quintom Bounce. In this model, a bouncing universe was obtained within the standard 4-dimensional Friedmann-Robertson-Walker (FRW) framework by making use of Quintom matter\\cite{Feng:2004ad}. The key feature of this model is that, the null energy condition (NEC) has to be violated for a short while around the bounce point and after that the equation-of-state (EoS) of our model $w$ is able to transit from below $-1$ to above $-1$ which makes the universe enter into normal expanding history.  The model of Quintom Bounce has presented a very interesting picture of the early universe. Firstly, the investigation of the cosmological perturbations in this model \\cite{Cai:2007zv} has shown a combination of some aspects found in some recently studied non-singular bounce models \\cite{Abramo,Brandenberger:2007by} and some others in singular bounce models \\cite{BGGMV,Lyth,Hwang2,Fabio}. Secondly, in a recent study on Quintom Bounce, the authors of Ref. \\cite{Cai:2008qb} have found that a concrete model of Quintom Bounce are able to provide a scale-invariant spectrum in ultra-violet regime, and also give rise to an oscillation signature which could be verified in the forthcoming astronomical observations.  One important lesson of studying the primordial curvature perturbations is to investigate its bispectrum which describes the non-Gaussianity of the power spectrum\\cite{Bartolo:2004if}. There is a number of information coded in this bispectrum, such as magnitude, shape, sign, and even running. To probe these signatures, we expect to distinguish various models of the very early universe. For example, in the usual slow-roll inflationary model, it is pointed out that non-Gaussianity is negligible due to the suppression of slow-roll parameters\\cite{Maldacena:2002vr,Acquaviva:2002ud}; however, in the models of Ekyrotic/cyclic \\cite{Koyama:2007if,Buchbinder:2007at,Lehners:2007wc} or island universe\\cite{Piao:2007cj}, a large non-Gaussianity is predicted. Observationally, current cosmological data\\cite{Yadav:2007yy,Komatsu:2008hk} is consistent with Gaussian distribution, however, the forthcoming observations will be sensitive to the non-Gaussianity with much higher precision, such as Planck satellite \\cite{Planck}. Given the considerations above, we in this paper calculate the bispectrum of a Quintom Bounce model. Our results show that non-Gaussianity in this model is still suppressed by slow-roll parameters, but there is an interesting oscillation signal on the non-linear parameter and its maximal value is mildly bigger than that in the usual scenario of slow-roll inflationary models.  Another important clue to discover the information of the very early universe is to study the relic gravitational wave background (GWB) formed by primordial tensor fluctuations. There have been a number of detectors operating for the signals of primordial GWB, e.g. Planck \\cite{Planck}, Big Bang Observer (BBO) \\cite{BBO}, LIGO \\cite{Abbott:2004ig}. Moreover, the indirect detection of GWB attracts a lot of interests of the next generation of CMB observations, see related analyses \\cite{Verde:2005ff, Smith:2005mm, Boyle:2005ug,TFCR}. The basic mechanism for the generation of primordial GWB in cosmology has been discussed in Refs. \\cite{Grishchuk:1974ny,Allen:1987bk}. Usually, inflation theory predicts that the power spectrum of primordial tensor fluctuations is scale-invariant and its value is roughly equal to scalar spectrum times a slow-roll parameter defined as $\\epsilon\\equiv-\\dot H/H^2$\\cite{Starobinsky:1979ty,Stewart:1993bc}. However, this so-called consistency relation is not necessary to be valid in a bounce scenario. For example, Ref. \\cite{Boyle:2003km} has investigated the primordial gravitational waves in singular bounce models and predicted an undetectable GWB. We in this paper study the relic gravitational waves in Quintom Bounce with Coleman-Weinberg potential. Interestingly, we find that the power spectrum of primordial gravitational waves is scale-invariant both in ultra-violet and infrared regimes but strongly oscillate in the middle band which is related to the bounce, and also find a sunken signature on the energy spectrum. These results on one side support the conclusions of Ref. \\cite{Cai:2008qb} in which the scalar perturbations were discussed, and on the other side provide another approaches to searching the signals of a bounce.  The outline of this paper is as follows. In Section II, we briefly review the basic scenario of a Quintom Bounce model with a the Coleman-Weinberg potential, and the curvature perturbation. In Section III, we make the calculation of non-Gaussianity in the model of Quintom Bounce, and find that the non-linear parameter $f_{NL}$ is oscillating around a small value which coincide with that in usual inflation model. In Section IV, we investigate the tensor part of gravitational perturbations in this model, and then give the power spectrum and spectral index correspondingly. Considering the influences of possible factors during the late time evolution, we discuss the transfer function for the gravitational waves and finally obtain the energy spectrum of GWB we may probe in future observations. Section V contains discussion and conclusions. ",
        "conclusions": "Bouncing cosmology, due to the avoidance of the initial singularity, has attracted a lot of interests in the literature \\cite{Peter,Veneziano,Wands,Piao:2004me,Senatore,Buchbinder:2007ad,ArmendarizPicon:2003qk} (and see Ref. \\cite{Novello:2008ra} for a recent review). However, since it happens in extremely high energy regime, we hardly observe a bounce by experiments directly. So it is a debate whether a bounce has taken place or not. To find the evidences of a bounce, we need to know what can a bounce leave for observations. This question is still discussed drastically in the literature, and one potential clue is to study the primordial gravitational fluctuations. In the context of the Pre-Big-Bang scenario and in the cyclic/Ekpyrotic cosmology, the primordial curvature perturbation strongly depends on the physics at the epoch of thermalization, and thus an uncertainty of a thermalized surface is involved \\cite{BGGMV,Lyth,Hwang2,Fabio,Tolley:2003nx}. In the frame of loop quantum cosmology, it is argued that fluctuations before and after the bounce are largely independent \\cite{Bojowald:2007zza} (yet see Ref. \\cite{Corichi:2007am} for some criticisms). We in this paper have studied the perturbation theory of a Quintom Bounce model detailedly and show that there are some imprints of the bounce on CMB observations at large scales. In the main content we have analyzed both the linear and non-linear evolutions of scalar modes, and the tensor perturbations are also considered.  The model we considered is constructed by a double field Quintom model with a Coleman-Weinberg potential. We firstly have reviewed the background dynamics of this model, and obtained a scale-invariant scalar spectrum in virtue of an asymmetry of the background evolution around the bounce point. A similar but more phenomenological scenario has been studied in Refs. \\cite{Piao:2003zm} as a possible solution to the suppressed low multi-poles of the CMB anisotropies. Moreover, since the gravitational perturbations in sub-hubble region would change their propagations when pass through the surface between the contracting phase and the bounce, there would be an oscillation signature generated both on the linear scalar modes and non-linear ones. We have also calculated the non-Gaussianity and shown that the maximal value of non-linear parameter $f_{NL}$ predicted by our model is mildly bigger than the usual one in single scalar slow-roll inflation, but the central value is still suppressed by the slow-roll parameters. So we expect that a large non-Gaussianity might be generated by other mechanisms\\cite{Li:2008fm} in the frame of Quintom Bounce.  We in the last part of this paper focus on the behavior of the gravitational waves in Quintom Bounce. Due to the effects of a bounce, the solution of the tensor perturbation is quite different from the usual one. In our analysis, we find that the physics of a bounce would affect the evolution of primordial tensor perturbations at large scales of the universe, which corresponds to the physics in very early time. The behavior of the energy spectrum of the GWB in our model is similar to that in the single scalar inflationary model in high-frequency regime. However, for low-frequency regime the difference becomes larger. Moreover, there is a sunken area in the middle band which links high-frequency regime and low-frequency regime. If these signals would be detected, these might act as a smoking gun to the bouncing cosmology."
    },
    "0808/0808.0886_arXiv.txt": {
        "abstract": "We present an analysis of \\ion{Kr}{1} $\\lambda$1236 line measurements from 50 sight lines in the {\\it Hubble Space Telescope} Space Telescope Imaging Spectrograph and Goddard High Resolution Spectrograph data archives that have sufficiently high resolution and signal-to-noise ratio to permit reliable krypton-to-hydrogen abundance ratio determinations. The distribution of Kr/H ratios in this sample is consistent with a single value for the ISM within 5900 pc of the Sun, log$_{10}$(Kr/H) = $-9.02\\pm0.02$, apart from a rough annulus from between $\\sim$600 and 2500 pc distant. The Kr/H ratio toward stars within this annulus is elevated by approximately 0.11 dex, similar to previously noted elevations of O/H and Cu/H gas-phase abundances beyond $\\sim$800 pc. A significant drop in the gas-phase N/O ratio in the same region suggests that this is an artifact of nucleosynthetic history. Since the physical scale of the annulus' inner edge is comparable to the radius of the Gould Belt and the outer limit of heliocentric distances where the D/H abundance ratio is highly variable, these phenomena may be related to the Gould Belt's origins. ",
        "introduction": "Where krypton is detectable, its interstellar abundance has the capacity to be a singularly reliable gauge of the degree of nucleosynthesis in a given region of the interstellar medium (ISM). Krypton is a noble gas whose outer electron shell is spherically symmetric. Consequently, it is not prone to forming either mechanical or chemical bonds in diffuse interstellar clouds where, given its ionisation potential relative to that of hydrogen (13.999 eV versus 13.598 eV), it should be predominantly in neutral form. These properties imply that a constant krypton to hydrogen interstellar abundance ratio is synonymous with a well-mixed ISM on the length scale over which the gas is sampled. Conversely, any significant departures from an established large-scale mean represent nucleosynthetic artifacts that mixing has not yet erased. Identifying such artifacts is crucial to understanding topics as varied as the details of dust composition and the large-scale processes of galactic chemical evolution. In particular, analyses of abundance departures from interstellar means constrain the observed efficacy of mixing at different Galactic length scales, the influence of isolated and/or prolonged but localized star formation events on elemental abundance ratios, the yields of supernovae, and the amount and composition of material available for dust formation. Because it is undepleted in the diffuse ISM, elemental abundance ratios involving krypton can be very sensitive probes for these fields of study.  Krypton to hydrogen abundance ratios have been well-studied within a kiloparsec of the Sun. The krypton line at 1235.838 {\\AA} was first detected using the Goddard High Resolution Spectrograph (GHRS) aboard {\\it Hubble Space Telescope (HST)} during observations of $\\zeta$ Oph \\citep{car91}. A comprehensive study of GHRS detections of \\ion{Kr}{1} $\\lambda$1236 later demonstrated that the Kr/H abundance ratio within 500 pc of the Sun is constant at the level of log$_{10}$(Kr/H) = $-9.02\\pm0.02$ \\citep{car97}. This conclusion was consistent with GHRS results from oxygen \\citep{mey98} and nitrogen \\citep{mey97}, two of the most abundant elements in the ISM, which also exhibited a constant local abundance relative to hydrogen. Moreover, since the interstellar krypton abundance ratio is roughly 50\\% of the Solar value ($-8.72\\pm0.08$; \\citealt{lod03}), this study established that the deficit in the interstellar krypton abundance with respect to its Solar value was similar to the deficits for other elements (e.g., carbon, oxygen, and nitrogen), using the then-accepted Solar abundances. Myriad spectral observations using the Space Telescope Imaging Spectrograph (STIS) and revisions both to oscillator strengths and Solar abundances have lessened the similarity between the interstellar abundance deficits for these elements; nevertheless the uniformity of their elemental abundances in the local ISM has only {\\it gained} further support from the new data. Indeed, since krypton does not readily deplete into grains or form molecules, the constancy of its abundance relative to hydrogen, independent of molecular hydrogen fraction [$f$(H$_2$)] or mean hydrogen sight line density [${\\langle}n_{\\rm H}\\rangle$], has been used to establish that the diffuse ISM is well-mixed on lengthscales of a few hundred parsecs \\citep{car03,car06}. This result is independently supported by other analyses, specifically some involving the composition of stellar atmospheres \\citep{red03} and pre-solar meteoritic dust grains \\citep{nit05}. As a result, it can be concluded that significant deviations from the mean interstellar Kr/H abundance ratio identify sites or regions within the Galaxy where the nucleosynthetic history differs from that of the local ($d \\lapprox$ 500-700 pc) ISM.  In order to apply this tool to exploring the limits of this well-mixed interstellar region and to probe for variations in the Kr/H abundance ratio based on physical conditions not found locally, it was necessary to search for krypton along a wider variety of paths than were previously studied. Thus, we have plumbed the depths of the MAST Archive at STScI for detections of \\ion{Kr}{1} $\\lambda$1236, measuring the krypton abundance for each sight line with a significant absorption feature. In this paper, we discuss the relevant GHRS, STIS, and {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} observations (\\S~\\ref{observations}), then examine the uniformity of the Kr/H abundance ratio with respect to key sight line properties (\\S~\\ref{krypton_distribution}). Finally, in \\S~\\ref{distance_effect}, variations in Kr/H, O/H, O/Kr, N/O, and Cu/H are compared, in search of reinforcement for the identified Kr/H trends, and possible explanations. ",
        "conclusions": "This paper has combined new observations from an {\\it HST} observing program abbreviated by the STIS power failure with Archival detections and previously published measurements of \\ion{Kr}{1} $\\lambda$1236 to compile a database of 50 Kr/H abundance ratios for sight lines up to 5.9 kpc in length and with average hydrogen densities from less than 0.10 cm$^{-3}$ to more than 10.0 cm$^{-3}$. Collectively, the data strengthen earlier evidence that krypton is undepleted in the ISM and that the Kr/H ratio is remarkably uniform within 750 pc of the Sun. Furthermore, these two properties underline the significance of an apparent enhancement of 0.10 dex in krypton abundance for sight lines of length 750---2500 pc. This abundance elevation is matched by enhanced O/H abundance ratios for the same sight lines, and similar distance dependencies in both gas-phase oxygen and copper abundances among other sight lines. However, N/O abundance ratios are reduced by a factor of two thirds in the same annulus. These disparate pieces of evidence suggest a nucleosynthetic origin for the abundance variations, while their spatial extent is comparable to the size of the Gould Belt. It is possible that the noted elemental abundance variations for Kr, O, Cu, and N arise out of star formation engendered as the Gould Belt was formed. Given the spatial extent of D/H variations, the Gould Belt may be at the bottom of this trend as well."
    },
    "0808/0808.0138_arXiv.txt": {
        "abstract": "Recently, a few peculiar Type Ia supernovae (SNe) that show exceptionally large peak luminosity have been discovered.  Their luminosity requires more than 1~\\Msun\\ of $^{56}$Ni ejected during the explosion, suggesting that they might have originated from super-Chandrasekhar mass white dwarfs.  However, the nature of these objects is not yet well understood.  In particular, no data have been taken at late phases, about one year after the explosion. We report on Subaru and Keck optical spectroscopic and photometric observations of the SN Ia 2006gz, which had been classified as being one of these ``overluminous\" SNe~Ia.  The late-time behavior is distinctly different from that of normal SNe Ia, reinforcing the argument that SN 2006gz belongs to a different subclass than normal SNe~Ia.  However, the peculiar features found at late times are not readily connected to a large amount of $^{56}$Ni; the SN is faint, and it lacks [Fe~II] and [Fe~III] emission. If the bulk of the radioactive energy escapes the SN ejecta as visual light, as is the case in normal SNe~Ia, the mass of $^{56}$Ni does not exceed $\\sim 0.3$~\\Msun. We discuss several possibilities to remedy the problem. With the limited observations, however, we are unable to conclusively identify which process is responsible.  An interesting possibility is that the bulk of the emission might be shifted to longer wavelengths, unlike the case in other SNe~Ia, which might be related to dense C-rich regions as indicated by the early-phase data. Alternatively, it might be the case that SN 2006gz, though peculiar, was actually not substantially overluminous at early times. ",
        "introduction": "There is general agreement that Type Ia supernovae (SNe Ia) are thermonuclear explosions of white dwarfs (WDs; e.g., Nomoto, Iwamoto, \\& Kishimoto 1997; Hillebrandt \\& Niemeyer 2000).  Thanks to the uniformity of their light-curve shapes after applying a correction factor (Phillips 1993; Phillips et al. 1999), they can be used as ``standard candles'' to measure cosmological parameters, which led to the discovery of the accelerating expansion of the universe and dark energy (Riess et al. 1998; Perlmutter et al. 1999; see Filippenko 2005 for a review). However, the nature of their progenitor systems has not been resolved (e.g., Livio 2000; Hillebrandt \\& Niemeyer 2000; Nomoto et al. 2003), making it difficult to reliably predict potential evolutionary effects that could add systematic errors to the determination of cosmological parameters.  The progenitors of normal SNe Ia (Branch, Fisher, \\& Nugent 1993; Li et al. 2001) are believed to be WDs having nearly the Chandrasekhar mass (hereafter Ch-SN Ia and Ch-WD).  The recent discovery of extremely luminous SNe Ia raises the possibility that not all SNe Ia originate from a single type of progenitor system.  Howell et al. (2006) reported that SN Ia 2003fg (SNLS-03D3bb) reached an absolute $V$-band magnitude of $M_V = -19.94$ ($H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M} = 0.3$, flat universe).  Assuming that this luminosity is powered by the decay chain $^{56}$Ni $\\to$ $^{56}$Co $\\to$ $^{56}$Fe as in other SNe Ia, they estimated $M_{\\rm 56Ni} \\approx 1.29$~\\Msun\\ (hereafter $M_{\\rm 56Ni}$ is the mass of $^{56}$Ni produced and ejected during the explosion).  Combining this with other elements whose existence in the ejecta is evident from the spectra, the ejecta and progenitor masses should exceed the Chandrasekhar limit of a nonrotating WD. This was the first observationally based suggestion of an SN Ia from a super-Chandrasekhar WD (hereafter SupCh-SN Ia and SupCh-WD).  Since then, two other possibly overluminous SNe~Ia have been reported: SN 2006gz (Hicken et al. 2007) and SN 2007if (Yuan et al. 2007).  \\begin{deluxetable}{llll} \\tabletypesize{\\scriptsize} \\tablecaption{Spectroscopy ($+ 341$ days)\\tablenotemark{a} \\label{tab:log_01}} \\tablewidth{0pt} \\tablehead{ \\colhead{Filter} & \\colhead{Wavelengths} & \\colhead{Airmass} & \\colhead{Exposure (s)} } \\startdata O58 + R300  & 5,800--10,200~\\AA & 1.056 & 1200 \\\\ None + B300 & 4,700--9,000~\\AA  & 1.085 & 1200 \\\\ Y47 + B300  & 3,800--7,200~\\AA  & 1.165 & 1200 \\\\ \\enddata \\tablenotetext{a}{The slit PA was 0 deg. The flux was calibrated with the standard star G191B2B (Massey \\& Gronwall 1990). } \\end{deluxetable}  All of the available data for these SNe~Ia are only for the early phases, $t \\lsim 100$ d (hereafter $t$ is the time since the explosion).  At later times, SNe Ia enter the nebular phase, when the Fe-rich innermost ejecta, which are hidden at early phases, can be directly observed (e.g., Axelrod 1988).  In this paper, we report late-time spectroscopy and photometry of the potentially overluminous\\footnote{In this paper, we define overluminous SNe Ia by the peak luminosity, corresponding to $M_{\\rm 56Ni}$ (mass of $^{56}$Ni) $\\gsim 1$ \\Msun. } SN 2006gz taken with the 8.2-m Subaru telescope and the 10-m Keck I telescope. In \\S 2 we present the observations and data reduction.  Results are shown in \\S 4, along with some comparisons to model calculations (\\S 3).  SN 2006gz is intrinsically different from normal SNe Ia.  However, our results are not readily interpretable in the SupCh-SN Ia scenario, raising questions about the nature of this new subclass of objects, as discussed in \\S 5. ",
        "conclusions": "We have presented late-time photometry and spectroscopy of SN 2006gz obtained with the Subaru and Keck I telescopes.  These are the first late-time observational data for an SN Ia that was claimed to be overluminous at early phases.  Such SNe have been suggested to be the explosions of SupCh-WD progenitors.  The spectrum is characterized by the weakness of the [Fe~II] and [Fe~III] lines in the blue ($\\lsim 5200$~\\AA) relative to emission lines in the red (either [Fe~II] or [Ca~II] at $\\sim 7200$~\\AA).  This does not fit into the sequence of the late-time spectral behavior of normal SNe Ia, confirming that SN 2006gz belongs to a different subclass of objects than normal SNe Ia.  Coupled with this is the problem of the observed faintness of SN 2006gz, especially in the $V$ band.  The SupCh2 model (with a 2~\\Msun\\ WD progenitor synthesizing 1~\\Msun\\ of $^{56}$Ni) is in good agreement with the early-phase bolometric light curve (\\S 4.1; but see \\S 5.1), although it predicts a higher late-time luminosity than the Ch-SN Ia models because of the larger $M_{\\rm 56Ni}$ and $\\gamma$-ray deposition rate.  The luminosity of Sup-Ch models exceeds that observed by more than 2 mag in the $V$ band. Furthermore, one derives $M_{\\rm 56Ni} \\lsim 0.3$~\\Msun, under the usual assumptions that the bulk of the deposited radioactive luminosity emerges at visual wavelengths, and that positrons are almost fully trapped within the ejecta.  In what follows, we list possible solutions for this inconsistency.   \\subsection{Reconsidering the Super Chandrasekhar Model?}  We should first mention that the peak luminosity of SN 2006gz derived by Hicken et al. (2007) involves large uncertainty, as the host-galaxy extinction is not well constrained.  If the host extinction is negligible, the peak luminosity of SN 2006gz is close to that of normal SNe~Ia ($M_{V} = -19.19$ mag).  In this case, SN 2006gz might be a less energetic explosion of a Ch-WD, possibly because of insufficient burning of the C-O layer to intermediate-mass elements, to be consistent with the observed slow light-curve evolution at early times.  It is not clear if the early-phase spectroscopic features can be explained by this model: the early-phase spectra (Hicken et al. 2007) indicate that the photospheric temperature is high (e.g., the weak Si II 5972) and the velocity is normal ($v_{\\rm Si~II} \\approx 11,000$--12,000 km s$^{-1}$ at $B$-band maximum). Thus, the most straightforward interpretation is that the peak luminosity is also high.  Although the Ch-WD model is an interesting possibility and deserves further investigation, it should not affect most of our conclusions about the peculiar late-time behavior, as our arguments are based mainly on the magnitude difference between the peak and the tail. In what follows, we assume the host extinction adopted by Hicken et al. (2007), but the above caveat should be kept in mind.\\footnote{For example, if the host extinction is smaller, then the estimated values for $M_{\\rm 56Ni}$ in the early and late-phases both become smaller.}  From the late-time data, we find that $m_V$ is $\\sim 1.3$ and $\\sim 2.2$ mag larger (that is, fainter) than expected from the SW7 ($M_{\\rm 56Ni} = 0.6$~\\Msun) and SupCh2 ($M_{\\rm 56Ni} = 1.0$~\\Msun) models, respectively.  This would imply that 0.1--0.2~\\Msun\\ of $^{56}$Ni powers the late-time emission, smaller than the value expected from the light-curve peak ($M_{\\rm 56Ni} \\approx 1$~\\Msun) by a factor of $\\gsim 5$.  Such a large discrepancy cannot be attributed solely to the effect of viewing angle, even if SN 2006gz was a result of an extremely off-axis explosion (see, e.g., Sim et al. 2007 for the effect of viewing angle; see also Maeda, Mazzali, \\& Nomoto 2006a).\\footnote{This does not rule out the possibility that SN 2006gz is a Ch-SN Ia with a large offset and viewed from the side of the $^{56}$Ni blob, as suggested by Hillebrandt et al. (2007) for another overluminous SN~Ia, SN 2003fg. Their model would have the same problem in reproducing the late-phase data presented in this paper, and would require additional processes to explain it, as do the SupCh models.}  Alternatively, one may hypothesize that the energy source at early times was not $^{56}$Ni/Co/Fe decay.  Strong circumstellar interaction as seen in a peculiar class of SNe Ia (i.e., SNe Ia/IIa 2002ic-like events: Hamuy et al. 2003; Deng et al. 2004; Wang et al. 2004) is not favored, because of the lack of evidence for interaction in the early-phase spectra.  Moreover, the SupCh models yield a reasonable fit to the early-time {\\it UBVRI} bolometric light curve without the fine tuning of parameters.  The early-phase spectroscopic features also seem to be consistent with SupCh models (\\S 4.2; K. Maeda, in preparation). Thus, $^{56}$Ni/Co energy input is the most likely mechanism to power the early-phase emission.  \\subsection{Positron Escape, Infrared Catastrophe, or Dust Formation?}  If the early-phase emission is powered by $\\sim 1$~\\Msun\\ of $^{56}$Ni, then one is left with the possibility that some mechanism, which is not at work in normal SNe~Ia, must be affecting the late-time emission and the thermal conditions within the ejecta of SN 2006gz.  One possibility is the escape of positrons produced by the $^{56}$Co decay out of the SN ejecta.  The issue has been comprehensively explored by Milne, The, \\& Leising (2001). They showed that a fraction of the positrons can escape for a radially combed and/or weak magnetic field, leading to the changing light curve at late times.  However, at $t \\approx 400$ d, this effect can account for at most $\\sim 0.5$ mag, which is not large enough to completely remedy the present problem.  Another possibility is the thermal catastrophe within the $^{56}$Ni-rich region, which shifts the bulk of the emission from optical to infrared wavelengths (Axelrod 1988). This so-called ``infrared catastrophe'' (IRC) is expected to take place after the temperature drops below a few 1000~K, depending on the electron density.  Observationally, the IRC has not been clearly detected in any SNe Ia.  The IRC does not take place at $t \\approx 1$ yr in any of our models: the heating rate per unit volume in the Fe-rich emission region is only a factor of $\\sim 1.5$ smaller in SupCh2 than in SW7, and thus the thermal conditions are similar in the two models.  Finally, it may be possible that the visual light from the $^{56}$Ni-rich region is converted to near-IR/mid-IR wavelengths by dust formed in the C+O-rich region. Before maximum light, SN 2006gz showed the strongest C lines of any SN~Ia ever observed, with little evolution in the absorption velocity. This indicates that a dense C+O-rich shell or clumps are present in SN 2006gz, at least in the outermost region.\\footnote{Khokhlov, M\\\"{u}ller, \\& H\\\"{o}flich (1993) assumed that explosive carbon burning is ignited at the center of the sub-Chandrasekhar mass ($\\sim 1.2~ M_\\odot$) degenerate C+O core (with a central density as low as $\\rho_c \\approx 2 \\times 10^8$ g cm$^{-3}$) surrounded by a spherical extended envelope.  For such a configuration, they demonstrated that interaction between the ejecta and the hypothesized spherical envelope can give rise to a dense shell.  However, the WD should evolve until the central density becomes sufficiently high ($\\rho_c \\approx 2 \\times 10^9$ g cm$^{-3}$) for explosive carbon burning to occur.  At this point, the whole mass would be in an almost hydrostatic rotating WD (Uenishi, Nomoto, \\& Hachisu 2003; Yoon \\& Langer 2004) surrounded by a geometrically thin accretion disk.}  The presence of carbon may be a key to the understanding of the peculiar behavior at late phases, since carbon has a relatively high condensation temperature.  The W7 model has $n \\approx 10^7$ cm$^{-3}$ at the inner edge of the C+O-rich region\\footnote{ This relatively dense region is formed at the interface between burning and non-burning layers. The density structure is almost frozen because the time scale of the Rayleigh-Taylor instability becomes longer than the time scale of the bulk expansion of the WD. Such a density enhancement at the (non-spherical) interface is thus also expected in multi-dimensional explosions.} at 100~d after the explosion.  The corresponding density in the SupCh models is a few times higher. Recently, Nozawa et al. (2008) theoretically investigated dust formation in SNe~Ib.  They found a large amount of carbon grain formation in the C-enhanced outer He region, with a density comparable to that in the SN~Ia models, in their model at $t \\gsim 50$ d. Thus, carbon grain formation in SNe~Ia may also be possible if carbon is abundant in the outermost region.  For normal SNe~Ia, accelerated fading like that of SN 2006gz has never been reported.  This may be consistent with the estimate that the upper limit of the carbon mass fraction is as low as 0.01 (e.g., Tanaka et al. 2008).  Details of the dust formation process, addressing the difference between SN Ib and SN Ia models (heating rate by $^{56}$Ni, different oxygen mass fraction, etc.), should be investigated, as well as details of expected observational signatures during the dust formation. In this interpretation, a part of the outer Ca-rich region might be mixed with the C+O-rich region. It may partly explain the relative strength of the feature at $\\sim 7300$~\\AA\\ relative to the [Fe~II] and [Fe~III] in the blue, since [Ca~II] arising from the Ca-rich region could be less diluted than the Fe emission, contributing to the 7300~\\AA\\ feature.  \\subsection{Future Observations}  The possible interpretations we listed above are only speculative, but they predict distinct observational signatures, to be tested in future overluminous SNe~Ia.  First, late-time spectroscopy of other potentially overluminous SNe~Ia would be extremely useful to see whether SN 2006gz is special even among these peculiar SNe~Ia.  If some of them show the SN 2006gz-like peculiarities at late phases, the optical light curves spanning from early to late times should provide a probe: (a) the IRC and dust scenarios predict a sudden decrease in the visual luminosity, (b) the positron-escape scenario would manifest itself with a relatively gentle decrease of the luminosity, and (c) other heating scenarios would not necessarily follow the quasi-exponential decline.  Near-IR light curves could directly show the effects of the IRC and dust distinctly from other scenarios: we expect that a rapid increase in the near-IR brightness should occur simultaneously with a rapid decrease in the visual brightness, while other scenarios do not necessarily predict this behavior.  A temporal sequence of optical spectra would also be highly useful. In the dust formation scenario, we may see a transient red continuum and emission-line shifts at intermediate phases, as was seen in the peculiar SN Ib 2006jc (Smith, Foley, \\& Filippenko 2008).  Finally, the IRC and dust scenarios could be unambiguously distinguished with near-IR through mid-IR spectra. Blackbody radiation from the dust particles should be seen in the dust scenario, while line emission is the dominant cooling process in the IRC scenario."
    },
    "0808/0808.1812_arXiv.txt": {
        "abstract": "IRAS 21413+5442 and IRAS 21407+5441 are two massive star forming regions of high luminosity, likely associated with each other. Near-infrared photometry on these two IRAS sources was performed at UKIRT using the UFTI under excellent seeing conditions yielding an angular resolution of $\\sim$ 0.5 arcsec. Our results reveal details of stellar content to a completeness limit (90\\%) of J = 18.5, H = 18.0, and K = 17.5 mag in the two regions. In IRAS 21413+5442, we identify a late O type star, having large (H-K) color, to be near the centre of the CO jets observed by earlier authors. The UKIRT images reveal in IRAS 21407+5441, a faint but clear compact HII region around a central high - intermediate mass star cluster. We have detected a number of sources with large (H-K) color which are not detected in J band. We also present the GMRT radio continuum map at 1.28 GHz covering the entire region surrounding the two star forming clouds. The radio continuum fluxes are used to estimate the properties of HII regions which seem to support our near-IR photometric results. Based on our radio continuum map and the archival MSX 8.2 $\\mu$m image, we show that the two IRAS sources likely belong to the same parent molecular cloud and conjecture that a high mass star of large IR colors, present in between the two sources, might have triggered star formation in this region. However one can not rule out the alternative possibility that Star A could be a nearby foreground star. ",
        "introduction": "Appearance of a compact (CHII) or an ultra-compact HII (UCHII) region embedded in a molecular cloud signifies the formation of a massive star of spectral type earlier than $\\sim$ B3 (Shepherd \\& Churchwell \\cite{shep96}; Churchwell \\cite{chur02}). During this phase the massive star is believed to be close to its zero-age main-sequence (ZAMS), although pre-main-sequence manifestations such as outflows have been detected in some cases (e.g., Weintraub \\& Kastner \\cite{wein96}; Beuther et al. \\cite{beut02}; Kumar et al. \\cite{kuma06}). Massive star formation occurs always in clusters. Further the environment around the massive stars gets affected by their strong winds and energetic radiation. It is therefore important to study such CHII/UCHII regions in order to find out possible evolutionary stage of stellar content in their vicinity. Near-Infrared photometry was shown to be very useful for this purpose (e.g., Lada \\cite{lada85}). Supplementary radio continuum measurements can provide important physical parameters concerning the object under certain assumptions (e.g., Scheffler \\& Elsasser \\cite{sche88} and references therein). In this paper we describe a study in near-IR and radio continuum in and around two IRAS sources that are believed to be massive star forming regions by virtue of the presence of CHII/UCHII.  IRAS 21413+5442 (Object 1) is one of the highly luminous, massive young stellar objects (YSOs) in our Galaxy and is situated at an estimated distance of 7.4 kpc (Wouterloot \\& Brand \\cite{wout89}; Yang et al \\cite{yang02}). This source is identified with the presence of a UCHII region, called IRAS 21413+5442S, about 20 arcsec south of a compact HII region, called IRAS 21413+5442N (Miralles et al. \\cite{mira94}). The far infrared luminosity from IRAS fluxes at 12, 25 and 60 $\\mu$m was estimated to be 3.2 $\\times$ 10$^5$ L$_\\odot$ (Campbell et al. \\cite{camp89}). The CO surveys of Shepherd \\& Churchwell \\cite{shep96} have classified this source as a massive star forming region with outflows of high velocity gas. This indicates therefore the pre-main-sequence signature of the source that is believed to be at/near its ZAMS stage (due to the appearance of a UCHII region). Bronfman et al. \\cite{bron96} detected this object in their survey of CS(2-1) rotational emission (at 97.981 GHz) that is believed to be a signature of high density molecular gas. Ishii et al. \\cite{ishi98,ishi02} have identified the 3.1 $\\mu$m H$_2$O ice in absorption which indicates the presence of a disk or thick nebular matter around the object. Ishii et al. \\cite{ishi98,ishi02} classify this object as heavily obscured due to the larger excess in near-IR colors. Interestingly, Shepherd \\& Churchwell \\cite{shep96} have clearly mentioned that the source engine of the outflows (or jets) found in this object must be a heavily obscured (pre-main-sequence) star yet to be discovered.  The second object in our study namely IRAS 21407-5441 (Object 2) was classified as a UCHII region from its IRAS colors and high luminosity of 8 $\\times$ 10$^{4}$ L$_{\\odot}$ (Wood \\& Churchwell \\cite{wood89}: S$_{25}$/S$_{12}$ $\\geq$ 3.7 and S$_{60}/S_{12}$ $\\geq$ 19.3). The object however fell short of the sensitivity limit of Bronfman et al. \\cite{bron96} survey of CS(2-1) line, possibly due to its low density gas. This object is situated at about the same (kinematic) distance ($\\sim$ 8 kpc) as Object 1, and is in all likelihood associated with it (as quoted in Carral et al \\cite{carr99} based on the rotation curve due to Wouterloot et al \\cite{wout90}). At this distance, the two Objects are separated in projection by about 20 pc (9.5 arcmin). The MSX A band (8.2 $\\mu$m) image of the region seems to support the association of the two IRAS sources; it also shows an interesting bright point source nearly mid-way between the two IRAS sources (called Star A in Fig 1 and subsequent sections of the paper).  In this paper we describe the first sub-arcsec UKIRT near-infrared photometry and 1.28 GHz radio continuum mapping in an attempt to classify the stellar content in and around the two IRAS sources and to try and establish their association. We also examine the nature of Star A and its possible role in the region. We have also made use of archival 2MASS and MSX data to supplement our study. In section 2.1, the near-infrared photometry is presented and in section 2.2, we present 1.28 GHz radio continuum mapping of the two objects. In section 3, we describe the results from near-infrared photometry and radio continuum observations. Section 4 gives discussion of the results that reveal some new interesting insights on the two objects. Section 5 lists important conclusions of the present work.  \\section {Observations}  Fig 1, generated from the 2MASS archival images, shows the two regions of our interest namely IRAS 21413+5442 (to the left, Object 1) and IRAS 21407+5441 (to the right, Object 2). Marked in the figure are three clusters - one in IRAS 21413+5442 (Cluster 1) and two in IRAS 21407+5441 (Clusters 2C and 2N). A bright star with large visual extinction, marked as Star A, is present in between the two IRAS sources. Also shown in the figure are the MSX low resolution contours in A band (8.2 $\\mu$m) overlaid on the 2MASS image. One may notice the apparent association of the two IRAS sources with the Star A nearly at the middle. The observations made on these sources are described in the following two subsections.  \\subsection{UKIRT-UFTI Photometry at sub-arcsec resolution} Sub-arcsec resolution JHK photometry was performed at UKIRT during 17 July 2005 on the two objects under excellent seeing conditions ($\\sim$ 0.5 arcsec) using the UFTI (HAWAII-1 1024 $\\times$ 1024) camera. The overall integration times were 75 sec for each of the three bands. The plate scale is 0.09 arcsec/pixel that gives a useful central field of view of $\\sim$ 75$^{\\prime\\prime}$ on the two objects. The standard star used was FS 151 (Hawarden et al. \\cite{hawa01}) for both the sources.  \\begin{figure*} \\centering \\includegraphics[scale=0.75]{fig1.eps} \\caption{2MASS Ks-band image of the combined field of the Clusters 1 and 2. For Cluster 2 the central source is marked as Cluster 2C and the north-eastern source as Cluster 2N. At the middle of the image, the highly reddened object Star A is shown. The abscissa (RA) and ordinate (Dec) are for J2000 epoch. The overlaid contours represent MSX low resolution image of the region; with contour levels 2, 3, 4, 5, 10, 15 and 20 W/m$^{2}$/Sr. The linear size in the region shown at the bottom left assumes a distance of 7.4 kpc.} \\label{2mass12}% \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[scale=0.75]{fig2.eps} \\caption{UKIRT K-band image of IRAS 21413+5442 (Cluster 1). The inset at the bottom-left represents the central region (see dashed lines in the main fig) and shows the identified stars with large IR excess falling in the T Tauri star region and beyond (see text and Table 1). The star (No. 7) that was shown encircled is situated at the centre of the CO jets. There are several \"red objects\" that are detected only in H and K bands with large (H-K) color (see text, Fig 5 and Table 1). The abscissa (RA) and ordinate (Dec) are for J2000 epoch.} \\label{ukirkob1}% \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[scale=0.75]{fig3.eps} \\caption{UKIRT K-band image of IRAS 21407+5441: Central region (Cluster 2C) revealing the HII region with filamentary structures. The inset at the bottom-left shows stars identified, as in Fig 2 (see also text and Table 2). There are several \"red objects\" that are detected only in H and K bands with large (H-K) color (see text, Fig 5 and Table 2). The abscissa (RA) and ordinate (Dec) are for J2000 epoch.} \\label{ukirkob2c}% % \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[scale=0.75]{fig4.eps} \\caption{UKIRT K-band image of IRAS 21407+5441: North-eastern region (Cluster 2N). Stars identified are numbered (see text and Table 3). There are several \"red objects\" that are detected only in H and K bands with large (H-K) color (see text, Fig 5 and Table 3). The abscissa (RA) and ordinate (Dec) are for J2000 epoch.} \\label{ukirkob2n}% \\end{figure*}  The image processing was done by using standard IRAF tasks. Only those sources are considered which have photometric errors less than 0.1 mag for all the three bands. The 90\\% completeness limits are J = 18.5; H = 18.0 and K = 17.5 mag for the observations on the two objects.  We found from the MSX image archives (low resolution mosaic images at an angular resolution of 18$^{\\prime\\prime}$, see Price et al. \\cite{pric01}), that the two IRAS sources in question seem to be physically linked, being part of a region that encompasses a large sky area (about $8^\\prime \\times 8^\\prime$). Since the UKIRT photometry does not cover the entire region of MSX image data for the region, we have extracted 2MASS JHK$_{s}$ photometric data for regions in between the two IRAS sources, including Star A and its environs.  \\subsection{1.28 GHz Radio Continuum Observations at GMRT}  The radio continuum observations in the 1.280 GHz frequency band (with a bandwidth of 16 MHz) were made on 1 August 2003 at the Giant Metre-Wave Radio Telescope (GMRT), operated by National Centre for Radio Astrophysics (NCRA), Tata Institute of Fundamental Research (TIFR) near Pune (details of the telescope can be found in Swarup et al. \\cite{swar91}). Out of the total 30 antennae, 26 had operated on the day of observations. The calibration sources used were 3C286 and 3C48 for flux and 2202+422 for phase. The field of view includes both the objects encompassing a region of 10$^\\prime \\times 10^\\prime$. The image processing was done by using AIPS software. The data sets were checked for dead antennas, bad baseline, bad time ranges, spikes etc using the tasks UVPLT and VPLOT. The tasks UVFLG and TVFLG were used for subsequent editing. Maps of the field were generated by Fourier transform and subsequent cleaning and deconvolution using the IMAGR task. The final angular resolution on the map is $\\sim$ 30 arcsec (beam size), after using adequate UV tapering on the available base-lines. The rms noise in the radio map is 1.4 mJy/beam.   \\begin{figure*} \\centering \\includegraphics[scale=0.7,angle=-90]{fig5.eps} % \\caption{Color-Color diagrams (left) for the sources in the central HII regions in IRAS 21413+5442 (Cluster 1, in top panel) and Cluster 2C (middle panel) and Cluster 2N (bottom panel) in IRAS 21407+5441. The thick dashed curve represents the unreddened main-sequence stars; the thin dotted curve shows the unreddened giant stars. The straight lines indicate the extinction vectors for A$_{V}$ = 30 mag and the crosses on the vectors indicate the visual extinction for 5 mag intervals. The T Tauri stars fall on the dashed straight line, while the dot-dashed straight line represents the extinction vector for the T Tauri stars. Color-Magnitude diagrams (rightside panels) are also shown. The vertical lines indicate the unreddened main-sequence stars of different spectral types (on the left-side) and the same with an extinction of A$_{V}$ = 30 mag (on the right-side); the slanting lines showing the extinction vectors. The stars that are common for all the JHK bands are shown by asterisks and the stars that are detected only in HK bands are shown by open circles. The filled square represents the embedded star shown circled in Fig 2. The diamond symbol marked as 'A' represents the position of Star A, shown for comparison.} \\label{i1ccmd}% \\end{figure*} ",
        "conclusions": "The important conclusions of this work are:  (i) Seeing limited near-infrared photometry was made on two massive star forming regions IRAS 21413+5442 and IRAS 21407+5441 using UKIRT facility. A radio continuum mapping of the two regions was also done using GMRT facility in the 1.28 GHz band.  (ii) Several embedded early B type stars are detected in both IRAS 21413+5442 (Cluster 1) and IRAS 21407+5441 (Clusters 2C and 2N). Several pre-main-sequence stars are detected only in H and K bands having large (H-K) colors (red objects) in all the three clusters. One such red object in Cluster 1, possibly a late O type star is located very close to the centre of the CO outflows in IRAS 21413+5442. A new CHII region was identified in the Cluster 2C which is powered by early B type or late O type stars.  (iii) Our infrared photometry and the radio map along with the MSX A band image indicate that the two IRAS sources are associated with each other. A highly obscured massive (O4 or earlier) star is shown to be present in the region between the two IRAS sources. It is conjectured that this single massive star might have triggered the star formation in the clusters. Further investigation is needed to establish this hypothesis beyond doubt.  (iv) The radio continuum observation at 1.28 GHz confirms the earlier results that the UCHII region in IRAS 21413+5442 is optically thick with the spectral index 1.0. The radio properties seem to support the NIR photometric results. \\\\"
    },
    "0808/0808.3817_arXiv.txt": {
        "abstract": "{Extrasolar planets observation and characterization by high contrast imaging instruments is set to be a very important subject in observational astronomy. Dedicated instruments are being developed to achieve this goal with very high efficiency. In particular, full   spectroscopic characterization of low temperature planetary companions is an   extremely important milestone.} {We present a new data analysis method for long slit spectroscopy (LSS) with   coronagraphy, which allows characterization of planetary companions of low   effective temperature. In a speckle-limited regime, this method allows an   accurate estimation and subtraction of the scattered starlight, to extract a clean spectrum of the planetary companion.  } {We performed intensive LSS simulations with IDL/CAOS to obtain realistic   spectra of low ($R=35$) and medium ($R=400$) resolution in the J, H, and K   bands. The simulated spectra were used to test our method and estimate   its performance in terms of contrast reduction and extracted spectra   quality. Our simulations are based on a software package dedicated to the development of SPHERE, a second generation instrument for the ESO-VLT.} {Our method allows a contrast reduction of 0.5 to 2.0 magnitudes compared to the   coronagraphic observations. For M0 and G0 stars located at 10~pc, we show that   it would lead to the characterization of companions with \\teff of 600~K and   900~K respectively, at angular separations of 1.0\\as. We also show that errors   in the wavelength calibration can produce significant errors in the   characterization, and must therefore be minimized as much as possible.} {} ",
        "introduction": "\\label{section:introduction}  Detection of exoplanets through indirect methods, such as radial velocities or transits has become a subject of intensive research that has led to the discovery of more than 270 exoplanets. A future generation of instruments will be able to deliver very high contrast coronagraphic images of nearby stars to detect possible orbiting companions. These direct detections are particularly effective at detecting relatively hot companions at large separations from their host stars, making them the perfect complement to radial velocity and transit methods that are usually able to detect planets with short orbiting periods.  The European project SPHERE (Spectro-Polarimetic High contrast imager for Exoplanets REsearch) for the VLT second generation of instruments aims to detect exoplanets down to a contrast of $10^{-6}$ at angular separations as small as 0.1\\as with dedicated extreme adaptive optics and coronagraphs \\citep{dohlen2006}. One of the 3 science channels of SPHERE will be equipped with a near infrared camera, IRDIS (InfraRed Dual Imaging Spectrograph), that will offer several observing modes for exoplanet detection and characterization including long slit spectroscopy (LSS). The prime objective of IRDIS is exoplanet direct detection with simultaneous differential imaging (SDI), a technique that has been extensively described in the literature \\citep{racine1999,marois2000} including possible improvements \\citep{marois2006}. It was also used successfully to detect cool companions with the NACO instrument on the ESO-VLT \\citep{lenzen2004}. This technique relies on the fact that the planetary spectrum has deep absorption features due to its atmosphere composition (e.g. CH$_{4}$ in the H band), whereas the starlight has a relatively smooth spectrum over the observation band. Recording two simultaneous images inside and outside of the planet absorption band, and subtracting them allows removal of most of the starlight, leaving mainly the planetary light in the subtracted image residuals.  Although very powerful at detection, this technique appears less suitable for the characterization of exoplanets. Current models predict that giant gas planet atmospheres most probably contain methane and water \\citep{allard2001,allard2003,burrows1997,burrows2006}. The reliance on the related molecular features for detection is a primary step in characterization, since it identifies planetary objects with atmosphere deprived of these compounds. However, a full characterization is not possible with a flux difference of only 2 narrow and adjacent bands, which is the information provided by the SDI technique: some degeneracy in planet parameters is probable if only the methane band in H is observed. It can also be difficult to exclude signal contamination by a background object without more information. This explains why follow-up spectroscopic observations over a significantly wide wavelength range are required.  In IRDIS, this observation mode is achieved with long slit spectroscopy at small ($R=35$) and medium ($R=400$) resolutions and coronagraphy over J, H and K bands. In this context, we adapted and improved the data reduction technique first developed by \\citet{sparks2002} and used by \\citet{thatte2007} with AO-fed Integral Field Spectroscopy (IFS), to extract a full companion spectrum from simulated high-contrast images. In contrast to its use with an IFS, LSS provides spatial information in only one dimension (along the slit), making it necessary to develop a specific data analysis method.  Our method is very general and can be used on coronagraphic images as well as non-coronagraphic ones. Moreover, it is fairly simple to implement, and does not require a large number of operations. It is fully described here in the context of SPHERE, since it is intended to be used for LSS data reduction on this particular instrument. Section \\ref{section:data_analysis} describes the concept and practical implementation details, Sect. \\ref{section:lss_simulations} provides a global overview of the simulations performed to simulate realistic LSS data, and Sect. \\ref{section:results} details the global performance of the method and the influence of a few key parameters. ",
        "conclusions": "\\label{section:conclusions}  We have implemented a promising method of characterizing planetary companions using long slit spectroscopy with coronagraphy, which has been considered for very high contrast imaging instruments, such as VLT-SPHERE. The need to develop a specific method was mandatory to remove the scattered light residuals fully from the random speckles pattern in the spectra. Using the linear wavelength dependence of the speckle pattern, we were able to evaluate to high precision the precise contribution of the star to the spectrum, and remove this contribution before extracting a clean planetary companion spectrum.  The simulations, performed using IDL, allowed us to test our method on realistic data for various cases of contrast in the case of low ($R = 35$) and medium ($R = 400$) resolution spectroscopy with extreme AO and Lyot coronagraphy. The final gain of the method for a 1 hour exposure on M0 and G0 stars at 10~pc is of the order of 0.5 to 2 magnitudes in terms of contrast over the J, H and K bands, compared to the coronagraphic profile. Although not very high in K band, the gain is substantial in J and H, allowing us to study companions with \\teff as low as 600~K. Following the evolutionary models of \\citet{baraffe2003}, a 600~K companion in a 500~Myr system corresponds to a mass of $\\sim$10~$M_{\\mathrm{Jup}}$.  To estimate the method efficiency, a quality factor has been introduced to measure the correlation between the input planetary spectrum in the simulation and the extracted spectra, as well as the discrepancy between the two. This allowed us to be confident that our method will allow precise characterization at LRS and MRS of companions as cool as 600~K around M0 stars at 10~pc, or 900~K around G0 stars at 10~pc, for angular separations larger than 1.0\\as. For smaller angular separations, results should be extremely valuable for slightly warmer companions.  Finally, we have estimated the influence of two key data analysis parameters at LRS. It is necessary to mask the planet signal during the analysis to avoid overestimating the scattered light level, which would lead to a mis-estimation of the planetary continuum. The method also relies on accurate wavelength calibration to be able to use the known wavelength dependence of the speckles to eliminate them. A systematic error as small as 10~nm can have major consequences on the characterization of the observed objects.  \\vspace{0.5cm} SPHERE is an instrument designed and built by a consortium consisting of LAOG, MPIA, LAM, LESIA, LUAN, INAF, Observatoire de Gen\\`eve, ETH, NOVA, ONERA and ASTRON in collaboration with ESO."
    },
    "0808/0808.3998_arXiv.txt": {
        "abstract": "We report the first discovery of the emergence of a high-velocity broad-line outflow in a luminous quasar, J105400.40+034801.2 at redshift $z\\sim 2.1$. The outflow is evident in ultraviolet \\civ\\ and \\siiv\\ absorption lines with velocity shifts $v\\sim 26,300$ \\kms\\ and deblended widths FWHM $\\sim$ 4000 \\kms . These features are marginally strong and broad enough to be considered broad absorption lines (BALs), but their large velocities exclude them from the standard BAL definition. The outflow lines appeared between two observations in the years 2002.18 and 2006.96. A third observation in 2008.48 showed the lines becoming $\\sim$40\\% weaker and 10\\% to 15\\% narrower. There is no evidence for acceleration or for any outflow gas at velocities $\\la$23,000 \\kms . The lines appear to be optically thick, with the absorber covering just 20\\% of the quasar continuum source. This indicates a characteristic absorber size of $\\sim$$4\\times 10^{15}$ cm, but with a BAL-like total column density $\\log N_H ({\\rm cm}^{-2})\\ga 21.2$ and average space density $n_H\\ga 2\\times 10^5$ \\cmn . We attribute the emergence of the outflow lines to a substantial flow structure moving across our line of sight, possibly near the ragged edge of the main BAL flow or possibly related to the onset of a BAL evolutionary phase. ",
        "introduction": "Accretion disk outflows are an important part of the quasar phenomenon. They are known to be present in at least 50\\% of optically-selected quasars \\citep{Hamann04,Nestor08,Ganguly08,Paola08}. They might play a critical role in the accretion physics and in the observational properties of quasars, and they might expel enough metal-rich gas and kinetic energy to have a major impact on star formation and galaxy evolution in their surroundings \\citep{Kauffmann00,Richards02, diMatteo05,Everett05,Proga05}. The outflows appear most conspicuously in quasar spectra as broad absorption lines (BALs), which typically have velocity widths of 5000--20,000 \\kms . However, there is also a variety of weaker and narrower outflow lines that are just as common as BALs. These include the so-called ``mini''-BALs, which appear at velocities from near 0 to $\\sim$0.2$c$ with profile widths from the BAL range down to a few hundred \\kms . Understanding the relationships between the various outflow lines is fundamental to our understanding of the outflows themselves. For example, in one scenario, the different line types represent different manifestations of a single outflow phenomenon viewed at different angles \\citep{Elvis00,Ganguly01}. Alternatively, they might represent an evolutionary sequence, where weak mini-BALs appear near the beginning or end of a more powerful BAL outflow phase \\citep{Hamann04}.  Variability studies can help test these ideas by providing constraints on the outflow dynamics, stability, location and basic physical properties. Recent studies have shown, for example, that mini-BALs tend to be more variable than the broader and stronger BALs \\citep{Gibson08,Paola08,Capellupo08}. Approximately half of mini-BALs show significant variations on time scales from a few years to few months in the quasar rest-frame. In rare cases, weak mini-BALs become much stronger and broader like BALs, or they disappear altogether. Following the emergence of a quasar outflow would be particularly valuable, but we know almost nothing about such occurances because variability programs start with quasars that are already known to have outflow lines. The only known case of an emerging BAL outflow was in a narrow line Seyfert 1 galaxy at redshift $\\sim$0.03 \\citep{Leighly05,Leighly08}. The dramatic appearance of BALs in this object is surprising because it is much less luminous (by a factor of $\\ga$100) than typical BAL quasars.  In this paper, we report the first discovery of an emerging high-velocity BAL-like outflow in a luminous quasar, J105400.40+034801.2 at redshift $\\sim$2.1. We observed this quasar as part of a program to study metal-strong damped Ly$\\alpha$ (DLA) absorption systems \\citep{Herbert-Fort06,Kaplan08}. Those data showed the new appearance of broad outflow lines compared to spectra obtained several years earlier by the Sloan Digital Sky Survey (SDSS). In the sections below we discuss the observations (\\S2), the measured properties of the outflow lines (\\S3), and some theoretical implications of our results (\\S4).  \\section[]{Observations}  The Sloan Digitial Sky Survey (SDSS) obtained the first high-quality spectrum of J105400.40+034801.2 on 5 March 2002 (2002.18). We retrieved the fully reduced spectrum from the SDSS archives, Data Release 6. This spectrum covers wavelengths from 3805 \\AA\\ to 9210 \\AA\\ at resolution $R\\equiv \\lambda/\\Delta\\lambda\\approx 2000$ (150 \\kms ). The absolute fluxes measured through through a $\\sim$3$^{\\prime\\prime}$ aperture fiber are expected to be accurate to within a few percent \\citep{Adelman08}. The emission line redshift and the blue and red magnitudes reported by the SDSS are $z_{em}=2.095$, $g=18.11$ and $r=17.98$, respectively. We also note that this quasar is radio-quiet based on a non-detection in the FIRST radio survey \\citep{Becker95}.  We observed J105400.40+034801.2 on 16 December 2006 (2006.96) using the Multi-Mirrored Telescope (MMT) Spectrograph in the blue channel with the 800 groove/mm grating. Combined with a 1.5$^{\\prime\\prime}$ wide slit, this setup provided wavelength coverage from 3090 \\AA\\ to 5095 \\AA\\ at resolution $R \\approx 1370$ (220 \\kms ). The exposure time was 720 s. After noticing the broad outflow lines in the MMT data, we obtained another spectrum on 30 May 2008 (2008.48) using the Low Resolution Imaging Spectrograph (LRIS) at the Keck Observatory. In this case, a 600 groove/mm grating blazed at 6000 \\AA\\ and a 1$^{\\prime\\prime}$ entrance slit provided wavelength coverage from 3105 \\AA\\ to 5600 \\AA\\ at resolution $R\\approx 1000$ (300 \\kms ). The total exposure time was 400 s. The observations at the MMT and Keck were both performed with the slit at the parallactic angle. The spectra were extracted and processed by standard techniques using the Low Redux Pipeline\\footnote{http://www.ucolick.org/$\\sim$xavier/LowRedux/index.html} software package. The relative fluxes were calibrated using spectrophotometric standards observed on the same night. The absolute fluxes have uncertainties up to 50\\%.  \\section[]{Results}  Figure 1 shows the three spectra obtained in 2002.18 (SDSS), 2006.96 (MMT) and 2008.41 (Keck). The wavelengths in the quasar rest frame are based on $z_{em}=2.095$ reported by the SDSS. The vertical flux scale in the figure applies to the SDSS spectrum. The other spectra are scaled vertically to match the SDSS fluxes at continuum wavelengths. The SDSS and MMT spectra are shown after smoothing for easier display. This smoothing has no effect on the appearance of broad features like the outflow lines and the broad emission lines (BELs).  \\begin{figure*} \\includegraphics[scale=0.65]{plot5.eps} \\caption{Spectra of J105400.40+034801.2  at rest-frame wavelengths (bottom scale) and observed (top) are represented by the black, blue and red curves for the three observations indicated. The broad \\siiv\\ and \\civ\\ outflow lines are marked by horizontal bars above the data. The locations of BELs are marked across the top. The steep rise at short wavelengths is due to strong \\nv\\ and Ly$\\alpha$ emission. All of the strong narrow absorption lines (not labeled) belong to an unrelated DLA system at $z_{abs}=2.0683$. The vertical flux scale applies to the SDSS spectrum. The other spectra are scaled vertically to match. The formal 1$\\sigma$ errors are shown by dotted curves near the bottom.} \\end{figure*}  Broad high-velocity outflow lines are clearly present in \\siiv\\ \\lam\\lam 1394,1403 and \\civ\\ \\lam\\lam 1548,1551 in the 2006.96 and 2008.48 spectra. They appeared in $\\leq$1.54 yr (rest frame) between the 2002.18 and 2006.96 observations, and then they weakened again 0.47 years later in the 2008.48 measurement. Meanwhile, the BEL profiles and equivalent widths (including Ly$\\alpha$ and \\ciii ] \\lam 1909 not shown in Fig. 1) varied by $\\la$5\\%. The discovery spectrum of J105400.40+034801.2 obtained in 1988.15 \\citep{Clowes94} is much noisier and has a lower resolution ($R\\sim 440$) than the data shown in Figure 1. Outflow lines like the ones measured in 2006.96 would probably not be detectable in those data. However, the 1988.15 spectrum does rule out the presence of broad absorption features with $\\sim$2 or more times the strength of the lines in 2006.96. No other outflow lines in the broad outflow system are detected in any of the data.  The 2006.96 and 2008.48 spectra both have a depression at the position of \\nv\\ \\lam\\lam 1239,1243 that is consistent with the broad absorption in \\civ\\ and \\siiv . However, we cannot confirm the reality of this feature because the spectra at those wavelengths are severely contaminated by Ly$\\alpha$ forest lines.  To measure the outflow lines quantitatively, we define a ``pseudo\"-continuum across the true quasar continuum and the tops of the BELs. This pseudo-continuum matches the 2008.48 spectrum except at the wavelengths of the broad absorption, where it follows the 2002.18 data. The result is shown in the top panel of Figure 2. Notice that the pseudo-continuum ignores the bump in the 2002.18 SDSS spectrum around 3900 \\AA\\ (1260 \\AA\\ rest). We attribute this bump to an anomaly in the SDSS flux calibration. Our experience with many similar SDSS spectra indicates that anomalies like this do sometimes occur near the blue edge of the SDSS wavelength coverage. In any case, this feature is not related to the \\siiv\\ outflow line because it has no analog near the \\civ\\ feature.  \\begin{figure} \\includegraphics[scale=0.49]{plot7.eps} \\caption{Top panel: Spectra of J105400.40+034801.2 in the same format as Figure 1, but with the SDSS data now drawn as a green curve and with fits to the pseudo-continuum and the outflow absorption lines shown by the smooth black curves. Bottom panel: Normalized spectra from 2006.96 (blue) and 2008.48 (red) together with the normalized line fits (black).} \\end{figure}  The bottom panel in Figure 2 shows the 2006.96 and 2008.48 spectra normalized by the pseudo-continuum. We fit the broad \\siiv\\ and \\civ\\ absorption lines in these normalized spectra using gaussian optical depth profiles. The \\siiv\\ and \\civ\\ transitions are doublets with velocity separations of $\\sim$500 \\kms\\ and $\\sim$1930 \\kms , respectively. Our fits include one gaussian per doublet member, with the redshifts and doppler velocities tied together. We assume that the doublet members have equal strengths based on the evidence for line saturation described in \\S4.1. In our fitting procedure, this means forcing the optical depths within the doublets to be equal even though their actual ratios based on oscillator strengths should be 2:1. The fits assume implicitly that the absorbing gas fully covers the quasar emission regions. We will discuss the effects of partial covering and line saturation in \\S4.1 below.  Figure 2 shows the final fits to the broad outflow lines. Table 1 lists some of the derived results, including the rest equivalent width (REW) of the entire blended doublet, and the absorption redshift ($z_{abs}$), the velocity shift ($v$ measured from $z_{em} = 2.095$), the full width at half minimum (FWHM) of the individual doublet lines, and the logarithmic ionic column density ($\\log N$) derived from the weaker doublet members. Note that the column densities are lower limits because of probable saturation effects (\\S4.1).  \\begin{table} \\centering \\begin{minipage}{77mm} \\caption{Outflow Line Measurements} \\begin{tabular}{@{}lccccc@{}} \\hline Line & REW & $z_{abs}$ & $v$ & FHWM & $\\log N$ \\\\ & (\\AA) &  & (\\kms ) & (\\kms ) & (cm$^{-2}$) \\\\ \\hline \\multicolumn{6}{c}{------------------------ 2006.96 (MMT) ------------------------}\\\\ Si IV & 6.21 & 1.833 & 26,430 & 4830 & $\\geq$14.84 \\\\ C IV & 4.83 & 1.835 & 26,270 & 3660 & $\\geq$15.08 \\\\  \\multicolumn{6}{c}{------------------------ 2008.48 (Keck) ------------------------}\\\\ Si IV & 4.07 & 1.838 & 25,900 & 4330 & $\\geq$14.66 \\\\ C IV & 2.54 & 1.830 & 26,810 & 3160 & $\\geq$14.80\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}  The broad outflow lines in 2006.96 had velocity shifts $v\\sim 26,300$ \\kms\\ and deblended FWHMs near $\\sim$4000 \\kms . The actual measured width of the \\civ\\ blend was FWHM $\\sim$ 4320 \\kms . These features are strong and broad enough to be considered BALs. However, their large velocity shifts lead to a balnicity index of zero, which excludes them from the standard BAL definition \\citep{Weymann91}. Their appearance in 2006.96 corresponds to strengthening by at least a factor of $\\sim$5 compared to the 2002.18 observation. Then they became roughly $\\sim$40\\% weaker and 10\\% to 15\\% narrower by 2008.48.  The uncertainties in our measured results are difficult to assess. The \\civ\\ data should be more reliable than \\siiv\\ because the spectra at the \\civ\\ wavelengths are less noisy and there is less blending with underlying BELs. There is no obvious explanation for the different FWHMs and velocity shifts reported for \\civ\\ and \\siiv\\ lines in Table 1 in terms of our fits to the lines or the pseudo-continuum. These differences might be real, or they might be indications of the true measurement uncertainties, mostly in the \\siiv\\ line. ",
        "conclusions": "We have described the first known case of an emerging broad-line outflow in a luminous high-redshift quasar, J105400.40+034801.2. The outflow is evident in \\siiv\\ and \\civ\\ absorption lines having FWHM $\\sim 4000$ \\kms\\ and velocity shifts $v\\sim 26,300$ \\kms . These features are not technically BALs because their large velocity shifts lead to a balnicity index of zero. Nonetheless, they represent a powerful outflow with $\\log N_H ({\\rm cm}^{-2})\\ga 21.2$ and $n_H\\ga 2\\times 10^5$ \\cmn . The outflow lines appeared between two observations 1.54 yr apart in the quasar rest frame. A third observation 0.47 yr later shows the lines becoming weaker again by $\\sim$40\\%. We attribute this behavior to a substantial flow structure crossing our line of sight to the continuum source.  Overall, the outflow lines discussed here are similar to high-velocity mini-BALs observed in other quasars \\citep{Paola08}. These features are characteristically narrower, weaker and more variable than the classic strong BALs. One possible explanation for these outflow lines is that they form along the ragged edge of the main BAL outflow, where the flow structures could be naturally smaller and more volatile. If BAL outflows are typically clumpy or filamentary \\citep{Hall07}, then perhaps the mini-BALs represent individual clumps or filaments along the edges of these flows. Another possibility is that mini-BALs and other similar outflow lines represent the sputtering beginning or end stages of a BAL evolutionary phase. Continued monitoring of quasars like J105400.40+034801.2, e.g., with multi-wavelength measurements to get better constraints on the ionization and column densities, will help to test these theoretical ideas."
    },
    "0808/0808.4095_arXiv.txt": {
        "abstract": "The various scenarios proposed for the origin of the multiple, helium-enriched populations in massive globular clusters are critically compared to the relevant constraining observations.  Among accretion of helium-rich material by pre-existing stars, star formation out of ejecta from massive AGB stars or from fast rotating massive stars, and pollution by Population III stars, only the AGB option appears to be viable. Accretion or star formation out of outflowing disks would result in a spread of helium abundances, thus failing to produce the distinct, chemically homogeneous sub-populations such as those in the clusters $\\omega$ Cen and NGC 2808. Pollution by Population III stars would fail to produce sub-populations selectively enriched in helium, but maintaining the same abundance of heavy elements.  Still, it is argued that for the AGB option to work two conditions should be satisfied: i) AGB stars experiencing the hot bottom burning process (i.e., those more massive than $\\sim 3\\;\\msun$) should rapidly eject their envelope upon arrival on the AGB, thus experiencing just a few third dredge-up episodes, and ii) clusters with multiple, helium enriched populations should be the remnants of much more massive systems, such as nucleated dwarf galaxies, as indeed widely assumed. ",
        "introduction": "The recent discovery of discrete, multiple stellar populations within several among the most massive globular clusters (GC) in the Milky Way (e.g., Bedin et al. 2004: Piotto et al. 2007) has brought new interest and excitement on GC research. Further excitement was added by the realization that some of these stellar populations are selectively enriched in helium to very high values ($Y\\sim 0.38$), without such enrichment being accompanied by a corresponding increase in the heavy element abundance of the expected size, if at all (e.g. Norris 2004; Piotto et al. 2005, 2007).  Thus, two main closely interlaced questions arise: how did such multiple populations form? and, which kind of stars have selectively produced the fresh helium, without contributing much heavy elements? Therefore, understanding the origin of multiple populations in GCs will need to make major steps towards a better knowledge of a variety of processes such as GC formation, star formation, as well as of some long standing issues in stellar evolution, e.g., the asymptotic giant branch (AGB) phase or the effect of rotation in massive stars.  Moreover, it is quite possible that some of the massive GCs with multiple populations are the compact remnants of nucleated dwarf galaxies (Bekki \\& Norris 2006), an aspect that widens even further the interest for this subject. All this together makes multiple populations in GCs an attractive, interdisciplinary subject of astrophysical investigation.  In this paper the main observational evidences are used to constrain the proposed possible scenarios for the origin of the multiple populations in GCs, and their associated chemical enrichment processes. In Section 2 the main relevant observational facts are briefly reviewed, and in Section 3 AGB stars and fast rotating massive stars are discussed as possible helium producers, along with Population III stars and the possible role of deep mixing during the red giant branch (RGB) phase of low mass stars.  In Section 4 various proposed scenarios for the origin of the multiple populations are confronted with the relevant observational facts, arguing that only massive AGB stars appear to remain viable as helium producers. A general discussion and the main conclusions are presented in Section 5. ",
        "conclusions": "In the previous sections it has been argued that only the AGB/S-AGB scenario remains viable to account for the helium-enriched sub-populations.  Still, with a serious difficulty to overcome, plus some minor ones. If our current understanding of the helium enrichment in intermediate mass stars is not grossly incorrect, then the mass of the first stellar population in a cluster such as $\\omega$ Cen had to be several times larger than the present mass of the cluster. Following Bekki \\& Norris (2006) this difficulty can be solved if clusters such as $\\omega$ Cen and NGC 2808 are the remnant nucleus of a nucleated dwarf galaxy that was torn apart by the tidal field of the Galaxy. Hence the parent first population providing the necessary raw material for the successive generation(s) would have been much more massive than these clusters are today. Some circumstantial evidence for this now widely entertained scenario is the finding that M54, one of the most massive GCs in the Galaxy, is indeed associated with the Sagittarius dwarf, albeit there is no evidence for multiple populations {\\it within} this cluster (Siegel et al. 2007). The nucleated dwarf NGC 205 may be another example relevant to this scenario: with its tidal stream towards M31 (McConnachie et al. 2004) it may represent an early stage of the process suggested by Bekki \\& Norris. Its nucleus is dominated by old stellar populations, and yet it is quite bright in the WFPC2 ultraviolet F225W and F185W passbands (Cappellari et al. 1999). If the UV light comes from an extended HB (such as that of $\\omega$ Cen and NGC 2808) it may also harbor helium-enriched populations, and would represent a possible testbed for the nucleated dwarf scenario of Bekki \\& Norris (2006).  Of course, for the nucleated dwarf scenario to work, successive stellar populations must have a much lower probability of being tidally stripped compared to the first population, otherwise the mass discrepancy would remain. This can only be achieved if the successive starbursts are far more centrally concentrated compared to the first one, i.e., the AGB/S-AGB ejecta from the first generation should collapse to the very bottom of the potential well before leading to star formation. In this connection, it is quite reassuring that the helium-enriched main sequence stars in $\\omega$ Cen are indeed markedly more centrally concentrated than the others (cf. Section 2.1.1).  Besides reproducing the required helium enhancement, a successful theory for the origin of multiple stellar populations in massive GCs should also account for the observed abundance and distribution of CNO and other intermediate mass elements. Several authors (e.g., Karakas et al. 2006; Romano et al. 2007; Choi \\& Yi 2008) have shown that yields of existing AGB models fail to satisfy this constraint, as along with helium stellar ejecta would be enriched also in carbon from the 3DU. The question is therefore as to whether one should conclude that the fresh helium does not originate from AGB stars, or that the used theoretical yields are not correct. Given the current uncertainties affecting the massive AGB models (especially due to the treatment of mass loss and convective overshooting) it seems worth keeping a pragmatic approach, hence taking existing AGB models with special caution. The schematic AGB evolution described in Section 3.1 is quite plausible given our ignorance of mass loss in bright AGB/S-AGB stars with HBB, and does not conflict with existing observations. Actually, it allows keeping AGB/S-AGB stars as viable helium producers for the multiple stellar generations in globular clusters.  For the present scenario to work it is essential that the AGB ejecta accumulate for a fairly long time ($\\sim 10^8$ yr) without any significant star formation before suddenly a major fraction of the interstellar medium is turned into stars by a burst. This may not be such an {\\it ad hoc} assumption, given that star formation in sporadic bursts appears to be the norm for dwarf galaxies (Gerola, Sneden \\& Sulman 1980), as also demonstrated by the discrete multiple generations in dwarfs such as Carina (e.g., Monelli et al. 2003).  In the case of NGC 2808 there had to be a second and a third stellar generation. If the scenario presented in this paper is basically correct, then the most helium rich secondary sub-population would have been the first to form out of the most massive AGB/S-AGB ejecta, which are the most helium rich. Then the ISM was replenished again by the ejecta of the less massive AGB stars from the first generation, plus perhaps a contribution by the secondary generation. If so, then the generation identified by the {\\it middle} of the three main sequences in this cluster was the last to form.  Worth briefly mentioning are those observational studies that may help testing the plausibility of the AGB/S-AGB scenario proposed in this paper. Some of these tests concern the adopted AGB/S-AGB evolution and nuclear yields, in particular concerning $\\sim 3$ to $\\sim 10\\;\\msun$ stars. It would be interesting to test whether the surface helium and nitrogen abundance are enhanced in $\\sim 6$ to $\\sim 10\\;\\msun$ main sequence stars, possibly by the meridional circulation process advocated by Maeder \\& Meynet (2006), albeit high rotational velocities may hamper accurate abundance determinations, and massive stars a metal poor as stars in $\\omega$ Cen and NGC 2808 are not within reach. Direct observations of bolometrically very luminous AGB/S-AGB stars in the Magellanic Clouds with very high mass loss rates should help understanding the crucial, final evolutionary stages of intermediate mass stars.  Moreover, high resolution spectroscopy of very large samples (several hundreds) of stars in the various sub-populations in globular clusters should help identify unequivocal chemical signatures of the {\\it donors} of the materials out of which these sub-populations have formed. These kind of studies are now possible thanks to the high multiplex multiobject spectrographs at 8--10m class telescopes, such as e.g., FLAMES at the VLT (Sollima et al. 2005; Carretta et al. 2006; Villanova et al. 2007).  Finally, the multifrequency study of nucleated dwarfs in and around the Local Group may help testing whether the most massive globular clusters may have originated from the tidal stripping of these objects.  In conclusion, excluding scenarios that {\\it qualitatively} conflict with observations, such as accretion, fast rotating massive stars or Population III stars, turn out to be much easier than proving others that qualitatively appear to work, but may have {\\it quantitative} difficulties. This is the case for the AGB/S-AGB scenario advocated in this paper, where the predicted helium abundance in the secondary populations admittedly falls a little short of the highest values suggested in the literature ($Y=0.38-0.40$). In this respect, one should consider than helium abundance estimates are affected by uncertainties that must be of the order of a few 0.01, hence no macroscopic discrepancy appear to exist. Still, it would help if, besides the 2DU, other processes contribute a little additional fresh helium in $\\sim 3$ to $\\sim 10\\,\\msun$ stars. Additional helium may come from the HBB process operating during the AGB/S-AGB phase, and/or from meridional circulations during the main sequence phase of the progenitors of AGB/S-AGB stars."
    },
    "0808/0808.2093_arXiv.txt": {
        "abstract": " ",
        "introduction": "Brightest cluster member (BCM) galaxies have historically held a special place in both the theoretical and observational study of galaxy evolution and formation since they are believed to be directly connected with the conditions in the cluster that they reside in. Consider, for instance, a BCM that is located at the bottom of the gravitational well of a galaxy cluster: there, a process such as cannibalism (e.g.\\ Ostriker \\& Tremaine 1975) would be at maximum efficiency there and we would be likely to observe BCMs with multiple cores (e.g.\\ Laine et al.\\ 2003; Yamada et al.\\ 2002; Oegerle \\& Hill 2001; Dubinski 1998; Lauer 1988; Hoessel \\& Schneider 1985). However, if a more hierarchical paradigm is correct, then the BCM should have formed through multiple minor merger events in a sub-group that was originally outside the cluster centre. It would have then subsequently merged with the rest of the cluster (Merritt 1985), yielding several observational road-signs such as high peculiar velocities for the BCM and cluster substructure (see Woudt et al.\\ 2008; Pimbblet et al.\\ 2006; Oergle \\& Hill 2001; Pinkney et al.\\ 1996; Quintana et al.\\ 1996; Valentijn \\& Casertano 1988).  Cases of dumbbell galaxies as BCMs (e.g.\\ Gregorini et al.\\ 1994; 1992) are doubly interesting for these reasons -- these galaxies should be an indicator of the pre-virialization scenario and should therefore be accompanied by significant substructure indicative of recent cluster merger activity. Indeed, Quintana et al.\\ (1996) present evidence for a dumbbell system where the two components of the BCM \\emph{each} belong to major sub-groups that are undergoing a merger (see also de Souza \\& Quintana 1990).  What is unclear is whether dumbbell BCM clusters in general are special? Does the presence of a dumbbell BCM point toward observable cluster activity such as subclustering (Quintana et al.\\ 1996) and large BCM peculiar velocities (Pimbblet et al.\\ 2006) as would be found during a merger event?  If so, then what stage is the merger event at? In order to make a pass at answering these questions, this work assembles a modest sample of dumbbell BCMs from Gregorini et al.\\ (1992) and a control sample derived from the 2 degree field galaxy redshift survey (2dFGRS; Colless et al.\\ 2001).  The format of this paper is as follows. In Section~2 we fully describe our derived dumbbell sample and control sample. In Section~3, we ask whether clusters with dumbbell BCMs are more likely to have substructuring than any other clusters and then in Section~4, we examine the incidence of significant peculiar velocities in our samples are different. We summarize our findings in Section~5. Throughout this work, we adopt a cosmological concordance model with values of $H_0 = 70$ kms$^{-1}$Mpc$^{-1}$, $\\Omega_M = 0.3$ and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "Even though the dumbbell sample is a collection of redshifts from many disparate sources, we believe that we have assembled a sufficient quantity of (bright) cluster members to adequately map out each of the clusters to a similar level found in 2dFGRS (cf.\\ Tables~1 \\&~2). Although the sample is likely not an optimal (homogeneously selected) one, we are nonetheless confident that the results presented would not change significantly with the addition of further observations.  Indeed, subtracting a small percentage of galaxies from the better sampled clusters produces no significant change in the measured peculiar velocities, or substructuring.  Our results show that dumbbell BCM clusters are no more or less likely than other clusters to possess (3:1) subclustering (cf.\\ Oegerle \\& Hill 2001). However, they are more likely to have at least one BCM component with a significant peculiar velocity.  We suggest the unifying interpretation of these observations is that our dumbbell BCM clusters may be in various states of post-merger activity.  Those with both subclustering and peculiar velocities (e.g.\\ Abell~3391) are probably in the early stages of mixing: the bulk of the galaxies are separable in to sub-groups with a dumbbell component belonging to one (or even each) sub-grouping (cf.\\ Quintana et al.\\ 1996). As the merger progresses, we may expect the less massive galaxies to homogenize first eventually leaving only a large peculiar velocity in one or both dumbbell components as the singular signpost to recent cluster activity, before the individual components of the dumbbell BCM itself relax with respect to each other and the rest of the merged cluster (Abell~2911). Alternatively, we note that intermediate clusters such as Abell~3570 may simply be the result of ellipticals preferentially being on radial orbits (e.g.\\ Ramirez \\& de Souza 1998; see also Hwang \\& Lee 2008) and hence observationally indistinguishable. The dumbbell itself may be a final transient phase before ultimately merging into a more massive cD-like galaxy (e.g.\\ Tremaine 1990; West 1994). These results and the emerging picture of the evolution of dumbbells with respect to cluster substructure are also consistent with other observational evidence that the timescale for clusters to accrete new galaxies is much shorter than the timescale required for the central galaxy to merge with the accreted galaxies (e.g. Cooray \\& Milosavljevi{\\'c} 2005; Brough et al.\\ 2008).   \\begin{figure*}[h] \\begin{center} \\vspace*{-1.5in} \\hspace*{-1in} \\includegraphics[scale=1, angle=0, width=8.in]{multiplot1.ps} \\vspace*{-1.5in} \\caption{Results of applying the DS test to the dumbbell sample. In these bubble plots, each circle's radius is scaled proportional to the Dressler \\& Shectman (1988) $\\delta$ statistic (i.e.\\ proportional to the deviation of a localized group of galaxies mean velocity away from the whole cluster's mean velocity). Substruture is therefore indicated by spatially close and relatively big overlapping circles.}\\label{fig:multi1} \\end{center} \\end{figure*}  \\begin{figure*}[h] \\begin{center} \\vspace*{-1.5in} \\hspace*{-1in} \\includegraphics[scale=1, angle=0, width=8.in]{multiplot2.ps} \\vspace*{-4.5in} \\setcounter{figure}{1} \\caption{continued.}\\label{fig:multi2} \\end{center} \\end{figure*}"
    },
    "0808/0808.2746_arXiv.txt": {
        "abstract": "For a mass-selected sample of 66544 galaxies with photometric redshifts ($z_{phot}$) from the Cosmic Evolution Survey (COSMOS), we examine the evolution of star formation activity as a function of stellar mass in galaxies. We estimate the cosmic star formation rates (SFR) over the range $ 0.2 < z_{phot} < 1.2$, using the rest-frame 2800\\,\\AA \\, flux (corrected for extinction). We find the mean SFR to be a strong function of the galactic stellar mass at any given redshift, with massive systems (log$(M/M_\\odot) > 10.5$) contributing less (by a factor of $\\sim$ 5) to the total star formation rate density (SFRD).  Combining data from the COSMOS and Gemini Deep Deep Survey (GDDS), we extend the $SFRD-z$ relation as a function of stellar mass to $z \\sim 2$. For massive galaxies, we find a steep increase in the $SFRD-z$ relation to $z\\sim 2$; for the less massive systems, the SFRD which also increases from z $= 0$ to 1, levels off at $z\\sim 1$. This implies that the massive systems have had their major star formation activity at earlier epochs ($z> 2$) than the lower mass galaxies.  We study changes in the SFRDs as a function of both redshift and stellar mass for galaxies of different spectral types. We find that the slope of the {\\it SFRD-z} relation for different spectral type of galaxies is a strong function of their stellar mass. For low and intermediate mass systems, the main contribution to the cosmic SFRD comes from the star-forming galaxies while, for more massive systems, the evolved galaxies are the most dominant population.   \\end {abstract} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.0605_arXiv.txt": {
        "abstract": "The probability density function (PDF) of the gas density in turbulent supersonic flows is investigated with high-resolution numerical simulations. In a systematic study, we compare the density statistics of compressible turbulence driven by the usually adopted solenoidal forcing (divergence-free) and by compressive forcing (curl-free). Our results are in agreement with studies using solenoidal forcing. However, compressive forcing yields a significantly broader density distribution with standard deviation $\\sim\\!3$ times larger at the same rms Mach number. The standard deviation-Mach number relation used in analytical models of star formation is reviewed and a modification of the existing expression is proposed, which takes into account the ratio of solenoidal and compressive modes of the turbulence forcing. ",
        "introduction": "The pioneering works by \\citet[][PNJ97 below]{PadoanNordlundJones1997} and \\citet[][PV98 below]{PassotVazquez1998} have shown that the standard deviation (stddev) $\\sigma_\\rho$, i.e., the width or dispersion of the linear density PDF $p_\\rho$ grows proportional to the rms Mach number $\\mathcal{M}$ of the turbulent flow, however, with proportionality constants in disagreement by a factor of $2$. The exact dependence of the PDF's stddev on the rms Mach number is a key ingredient for analytical models of star formation. For instance, \\citet{PadoanNordlund2002} and \\citet{HennebelleChabrier2008} relate the density PDF to the core mass function (CMF) and stellar initial mass function (IMF). \\citet{Tassis2007} uses the density PDF on galactic scales to reproduce the Kennicutt-Schmidt relation. \\citet{Elmegreen2008} suggests that the star formation efficiency is a function of the density PDF.  In the present work, we solve the discrepancy between PNJ97 and PV98 by showing that the stddev of the PDF is not only a function of the rms Mach number, but is also very sensitive to the relative importance of solenoidal (divergence-free) and compressive (curl-free) modes of the turbulence forcing, leading to variations in the stddev up to factors of $\\sim\\!3$ for the same rms Mach number. The main result of the present work is that the conclusions of PNJ97 and PV98 can be harmonized, if we take into account that PNJ97 have analyzed purely solenoidal forcing, whereas PV98 have studied purely compressive forcing. This apparent difference has not been considered analytically or numerically before.  We begin by explaining our numerical method in \\S\\ref{sec:methods}. \\S\\ref{sec:results} shows that our results are consistent with previous studies using solenoidal forcing, whereas compressive forcing yields a PDF with stddev $\\sim\\!3$ times larger in 3-dimensional turbulent flows. We present a heuristic model explaining this difference, which is based on the ratio of solenoidal to compressive modes of the forcing to estimate the proportionality constant in the stddev-Mach relation. \\S\\ref{sec:conclusions} summarizes our conclusions.  \\begin{figure*}[t] \\centerline{ \\includegraphics[width=0.49\\linewidth]{f1a.eps} \\includegraphics[width=0.49\\linewidth]{f1b.eps} } \\caption{Column density field in units of the mean column density for solenoidal forcing (\\emph{left}), and compressive forcing (\\emph{right}) at a randomly picked time in the regime of statistically stationary turbulence. Both maps show 4 orders of magnitude in column density with the same scaling for direct comparison of solenoidal and compressive forcing at rms Mach number $\\sim\\!5.5$.} \\label{fig:snapshots} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions} Performing high-resolution hydrodynamical simulations of driven isothermal compressible turbulence, we have found that the way of forcing the turbulence has a strong effect on density statistics. We have compared the usually adopted solenoidal forcing (divergence-free) and compressive forcing (curl-free). The most important result is that in 3D, compressive forcing yields a density PDF with stddev $\\sim\\!3$ times larger compared to solenoidal forcing for the same rms Mach number.  As found by PNJ97 and PV98, the stddev $\\sigma_\\rho$ is increasing directly proportional to the rms Mach number, however, they find different proportionality constants. Taking into account the ratio of solenoidal to compressive modes $\\zeta$ of the forcing resolves the disagreement between PNJ97 and PV98. We suggest that the proportionality constant $b$ in the stddev-Mach number relations~(\\ref{eq:sigrhoofmach}) and~(\\ref{eq:sigsofmach}) can be determined by a heuristic model summarized in eq.~(\\ref{eq:b}). In the case of compressive forcing ($\\zeta=0$), the proportionality constant $b\\!\\sim\\!1$ irrespective of the dimensionality of the simulation. Solenoidal forcing on the other hand yields $b\\!\\sim\\!1/3$ in 3D and $b\\!\\sim\\!1/2$ in 2D.  We mention the impact of our results on analytical models linking the statistics of supersonic turbulence to the CMF/IMF. \\citet{PadoanNordlund2002} and \\citet{HennebelleChabrier2008} rely on integrals over the density PDF to get a handle on the mass of objects above a certain density threshold. Since the width of the PDF is so sensitive to the mixture of solenoidal and compressive modes $\\zeta$, we also expect a strong influence on the derived CMF/IMF. Indeed, the larger dispersion obtained from compressive forcing leads to better agreement (Hennebelle 2008, private communication) of the analytic expression with the observed IMF in \\citet[][]{HennebelleChabrier2008}.  Given that many proposed sources of interstellar turbulence \\citep[e.g.,][]{ElmegreenScalo2004,MacLowKlessen2004} are likely to directly excite compressive modes (e.g., galactic spiral density shocks, large scale gravitational contraction, supernova explosions, protostellar jets, winds and outflows), it is reasonable to expect that turbulence in the ISM is driven by a mixture of solenoidal and compressive modes, possibly with compressive modes being more important than solenoidal modes."
    },
    "0808/0808.0757_arXiv.txt": {
        "abstract": "{} {We investigate the dependence of penumbral microjets inclination on the position within penumbra.} {The high cadence observations taken on 10 November 2006 with the Hinode satellite through the \\ion{Ca}{ii}~H and G--band filters were analysed to determine the inclination of penumbral microjets. The results were then compared with the inclination of the magnetic field determined through the inversion of the spectropolarimetric observations of the same region.} {The penumbral microjet inclination is increasing towards the outer edge of the penumbra. The results suggest that the penumbral microjet follows the opening magnetic field lines of a vertical flux tube that creates the sunspot.} {} ",
        "introduction": "\\citet[][hereafter Paper~I]{Katsukawa:2007} reported on the existence of small jet-like features that are observed at penumbral chromospheric layers. The authors used the term penumbral microjets (PJ) and this nomenclature is also used in this Letter. These new penumbral phenomena were found using the observations taken with the Hinode satellite through the broadband \\ion{Ca}{ii}~H filter.  As summarised in Paper~I, the PJs are highly transient events with lifetimes up to two minutes and lengths of a few thousand kilometres. The width of PJs is around 400~km for the largest ones where the smallest events are at the resolution limit of the observations. It can be expected that there are even smaller undetected penumbral microjets.  As explained in Paper~I, the three-dimensional configuration of the PJs can be estimated from the different visibilities of these events depending on the position on the solar disc. The intensity of microjets is comparable to the intensity of underlaying penumbral filaments, and the PJs can hardly be identified on observations taken close to the disc centre (if the running-difference or the high-pass filter are not applied). It implies that the azimuthal orientation of PJs is the same as for penumbral filaments, i.e., radial. In the case of observations taken farther from the disc centre, the microjets become visible due to the difference in their inclination compared to the inclination of penumbral filaments that are nearly horizontal. In Paper~I the authors estimated the inclinations (angle between the local normal line and the microjet) to be mostly between 40$^\\circ$ and 60$^\\circ$.  Although the observations with a cadence of four seconds can be achieved (with a limited field of view and only for a short time), the observations of penumbra with a 20~sec cadence are currently the best available measurements that can be used to study the properties of PJs. Taking the lifetimes of PJs into account, the available high-cadence observations are not fast enough to study the photospheric counterparts of the PJs onsets because the original formation area can become dark in 20~seconds. As already pointed out in Paper~I, there are some indications that the PJs are related to the penumbral bright grains observed in the photosphere. However, this topic is not discussed in this Letter and we concentrate on a detailed analysis of the penumbral microjets inclination and its dependence on the position within the penumbra. The results are then compared with the inclination of the photospheric magnetic field. ",
        "conclusions": "We identified 209 penumbral microjets in almost two-hour long observations of the penumbra in AR10923 using the \\ion{Ca}{ii}~H images taken with 30~s cadence. In combination with simultaneous G--band observations, the inclinations of these microjets (angle between the PJ and local normal line) are determined along with their approximate position in the penumbra. The results show a clear increase in the PJ inclination toward the outer penumbral edge. We find on average inclinations around 35$^\\circ$ at the umbra/penumbra boundary and 70$^\\circ$ at the penumbra/quiet sun boundary.  The found radial variation in the PJ inclination resembles the change in magnetic field inclination at higher photospheric layers across the penumbra. We find the difference of 10$^\\circ$ between the inclinations of magnetic field lines and penumbral microjets, with the former more vertical. This difference can be explained easily if we suppose that the PJs follow the magnetic field lines that are opening with height. There are a few observed PJs that show the change in inclination with height in the atmosphere and support this hypothesis.  The observed difference also implies that on average we detect the PJs at higher atmospheric layers and only rarely at the initial stage when the PJs inclination angles should be close to (or same as) the magnetic field inclination at higher photospheric layers. The scatter of individually determined values of PJs inclination in Fig.~\\ref{plot} can be influenced by the different time spans between the onset and the determination of the PJs inclination.  Although our results cannot clarify the mechanism responsible for the formation of the PJs, they imply that the PJs are guided by magnetic field lines that are fanning out with height. Higher cadence observations of \\ion{Ca}{ii}~H and G--band filtergrams and simultaneous SP measurements are needed to study this problem in detail."
    },
    "0808/0808.0561_arXiv.txt": {
        "abstract": "Using the ASTE 10 m submillimeter telescope and the 1.4 m Infrared Survey Facility (IRSF), we performed an extensive outflow survey in the Orion Molecular Cloud -2 and -3 region. Our survey, which includes 41 potential star-forming sites, has been newly compiled using multi-wavelength data based on millimeter- and submillimeter-continuum observations as well as radio continuum observations. From the CO (3--2) observations performed with the ASTE 10 m telescope, we detected 14 CO molecular outflows, seven of which were newly identified. This higher detection rate, as compared to previous CO (1--0) results in the same region, suggests that CO (3--2) may be a better outflow tracer. Physical properties of these outflows and their possible driving sources were derived. Derived parameters were compared with those of CO outflows in low- and high-mass star-forming regions. We show that the CO outflow momentum correlates with the bolometric luminosity of the driving source and with the envelope mass, regardless of the mass of the driving sources. In addition to these CO outflows, seven sources having NIR features suggestive of outflows were also identified. ",
        "introduction": "CO molecular outflows are ubiquitous phenomena in low- to high-mass star-forming regions (e.g. Bachiller \\& Tafalla 1999 for reviews, Zhang et al. 2001, Arce et al.  2007 for reviews). This phenomenon is considered to be directly associated with the main accretion phase in protostars, and is frequently used to identify protostars in dense cores. Hence the CO outflows provide the primary information concerning the mass accretion and ejection processes during the protostellar evolution. Previous single-dish surveys revealed the physical nature of CO outflows toward hundreds of low- to high-mass young stellar objects (e.g., Cabrit \\& Bertout 1992, Shepherd \\& Churchwell 1996, Bontemps et al. 1996, Beuther et al. 2002, Zhang et al. 2001, 2005). An extensive outflow survey of low-mass objects by Bontemps et al. (1996) found that the momentum flux of the CO outflows is positively correlated with the bolometric luminosity and also positively correlated with the mass of the associated parent cores. Recent high-mass studies also show the above correlations continuing to the high-mass region over several orders of magnitude in mass and bolometric luminosity (e.g., Curchwell 1997, Beuther et al. 2002, Zhang et al. 2005). On the other hand, the systematic studies focused on outflows associated with intermediate-mass protostars ($2~{\\leq}~M_{\\ast}~{\\leq}~8~M_{\\odot}$) are still limited, although several candidates have been investigated in detail to date (e.g., Fuente et al. 2005, Takahashi et al. 2006, Beltran et al. 2006, 2007). Studies of the intermediate-mass outflows are important for revealing the evolution of intermediate-mass stars. But more importantly, such studies address the question of wether the same basic accretion and outflow processes operate continuously from low- to high- mass stars.  In this paper, we present the results of our extensive outflow survey toward one of the nearest intermediate-mass star-forming regions: The Orion Molecular Cloud -2/3 (OMC-2/3; $d$=450 pc; Genzel \\& Stutzki 1989). The goals of this project are to (i)search for outflows in OMC-2/3 with high sensitivity and high spatial resolution, (ii) combine the detected outflow results and the centimeter- to mid-infrared continuum information to identify a driving source of the outflows, and (iii) reveal the structures and the physical properties of each outflow. In the OMC-2/3 region, 38 millimeter- and submillimeter dust continuum sources were detected (Chini et al. 1997, Lis et al. 1998, Nielbock et al. 2003). Free-free jets were also detected at 3.6 cm toward 11 of the embedded sources, and eight of 11 sources are associated with the dust continuum sources (Reipurth et al. 1999). Shock-excited H$_2$ knots, which probably trace shock fronts within each flow, are associated with these embedded sources (Yu et al. 1997, 2000, and Stanke et al. 2002). Aso et al. (2000) and Williams et al. (2003) made mapping observations in the CO (1--0) emission and identified nine CO outflows. They found that the outflows in OMC-2/3 are mostly spatially compact (i.e., possess short dynamical time scales of a few $\\times~10^4$ yr), and energetic (${\\sim}$ several ${\\times}~10^{-4}$ M$_{\\odot} $km s$^{-1}$). However, previous outflow identifications have ambiguities, because the CO (1--0) emission is often significantly affected by the contamination from the ambient molecular cloud. This made it difficult to identify small and faint outflows and differentiate them from the ambient molecular gas. Furthermore, pre- and proto-stellar candidates in OMC-2/3 are lying along the filament within a few $\\times$ 0.1 pc. The crowded core distributions also made it difficult to identify which is the driving source of each outflow.  We report results of our high-sensitivity molecular outflow survey in the submillimeter CO (3--2) emission with the ASTE telescope. The ASTE (Atacama Submillimeter Telescope Experiment) telescope equipped with a 345 GHz receiver is one of the most powerful facilities with which to seek molecular outflows over a wide field of views. The ASTE telescope in the ``On-The-Fly'' mapping mode enabled us to observe the entire region of OMC-2/3 with a high sensitivity limit (i.e., corresponding mass of the sensitivity limit is $\\sim$10$^{-4}~M_{\\odot}$). The CO (3--2) emission exclusively traces higher temperature gas ($\\geq$ 33 K). Our CO (3--2) results are compared with previous CO (1--0) results. We also show deep $JHK_s$ images taken by the Simultaneous Three-Color InfrarRed Imager for Unbiased Survey camera (SIRIUS) on Infrared Survey Facility (IRSF), which provide us with the distributions of faint reflection nebulae, chains of knots, and jet-like features. Furthermore, we also used the archived 24 $\\mu$m data obtained by the infrared camera MIPS (Multiband Imaging Photometer) on the infrared space telescope Spitzer, which are sensitive to hot dusty components with a temperature of ${\\sim}$ 150 K, and probably trace 100 AU scale innermost envelopes/circumstellar disks associated with the central heating sources.  The details of the observations of the submillimeter CO (3--2) emission and the near infrared $JHKs$-bands are presented in Section 2. In Section 3, we show the results of the outflow survey. Data analysis of the CO (3--2) outflows is shown in Section 4 and properties of the outflow driving sources are described in Section 5. In Section 6, we discuss the physical properties of the CO outflows in OMC-2/3 and compare them with outflows associated with low- to high-mass YSOs in previous studies. Finally, Section 7 summarizes the paper. ",
        "conclusions": "\\subsection{Comparison with previous CO (1--0) observations} In addition to our CO (3--2) observations, outflow observations in the OMC-2/3 region were also conducted by Aso et al. (2000) with the Nobeyama 45 m telescope, and Williams et al. (2003) with the FCRAO 14 m telescope and the BIMA array. These previous observations were done in CO (1--0). In this section, we compare our results with those obtained in the previous observations.  Out of 14 outflows (10 are categorized as CLEAR and 4 are categorized as PROBABLE) we identified, the one associated with CSO 33 located in the most southern part of the OMC-2/3 region was not included in the previous outflow observations. Out of the rest, eight outflows associated with MMS 2, MMS 5, MMS 7, MMS 9, VLA 9, FIR 1$b$, FIR 2 and FIR 3 were also detected in both of the previous observations although one of them, associated with VLA 9, was considered to be associated with MMS 10 in the previous observations (see section 3.3.3; see also below). The physical parameters of these eight outflows derived from our CO (3--2) data (see Table 4) are roughly in agreement with those derived from the previous CO (1--0) data. On the other hand, five outflows associated with SIMBA $a$, SIMBA $c$, VLA 13, FIR 6$b$, and FIR 6$c$ have been detected only in our CO (3--2) observations.  Why were these five outflows not detected in the previous CO (1--0) observations? One reason is the complexity of the region; some parts of the OMC-2 region form young stars as groups, and as a result, it is difficult for us to identify each outflow. For example, even though previous observations detected high velocity red-shifted CO (1--0) emission near VLA 13, it was not identified as an outflow associated with VLA 13. In this study, it was identified as an outflow associated with VLA 13 because of the near- and mid-IR data as well as the CO (3--2) data. A similar case is the outflow associated with VLA 9; previous studies misidentified this outflow as that associated with MMS 10 because of the lack of near- and mid-infrared data.  Another possibility is that high velocity emission is weaker in the CO (1--0). In the case of outflows associated with SIMBA $a$ and SIMBA $c$, no high velocity emission was clearly detected in CO (1--0) in the previous observations. These outflows are weak even in the CO (3--2) emission. Similarly, outflows associated with FIR6 $b$ and FIR6 $c$ are significantly weak in CO (1--0) even though they are quite strong in CO (3--2).  Are outflow emissions systematically weaker in CO (1--0) as compared with CO (3--2)? In order to examine this, we compared the intensity of CO (1--0) and (3--2) emissions toward the outflow lobes. For this comparison, we used new CO (1--0) outflow maps of the OMC-2/3 region, recently taken using the Nobeyama 45 m telescope with the OTF mode (Kawabe et al.) in addition to our CO (3--2) maps taken by ASTE. Since the CO (1--0) maps have a higher angular resolution (22$''$) than the CO (3--2) maps (26$''$), the CO (1--0) maps were convolved with a beam of 26$''$. For the outflows significantly detected in both CO (1--0) and (3--2), CO (3--2)/CO (1--0) intensity ratios were estimated to be 1-3, suggesting that CO (3--2) emission is at least comparable to, or brighter than CO (1--0) for the outflows in the OMC-2/3 region.  It is interesting to consider what kind of physical conditions could make the CO (3--2) emission a few times stronger than the CO (1--0) emission. With the Large Velocity Gradient (LVG) calculations, the CO (3--2)/CO(1--0) ratio would be more than unity when H$_2$ density range of 1.5$\\times$10$^3$ to 2$\\times$10$^4$ cm$^{-3}$ and the kinematic temperature of 40 K. The gas kinetic temperature at some of outflow lobes actually estimated to be more than 40 K from CO (3-2) line profiles. Hence, these molecular outflows we observed in the OMC-2/3 region probably have high density gas of $\\sim$10$^4$ cm$^{-3}$. Such high density and high temperature gas were found in the molecular outflow associated with L1157, one of the best-studied molecular outflows (Hirano \\& Taniguchi 2001).   \\subsection{Outflow Properties: OMC-2/3 Region versus Other Regions}  In the following section, we compare the physical nature of the outflows in the OMC-2/3 region with those found through previous systematic outflow searches in other star-forming regions. In order for us to minimize the ambiguity of the outflow identification, we limited our outflow sample in the OMC-2/3 region to those categorized as CLEAR or PROBABLE in the sections 6.2.1 and those categorized as CLEAR in the section 6.2.2.  \\subsubsection{Detection Rate of the outflows} Among the 41 potential star-forming sites listed in Table 1, 14 sites were found to be associated with CO outflows categorized as CLEAR or PROBABLE. When we simply calculate the CO outflow detection rate using this result, the rate is 34 \\%. On the other hand, previous outflow searches in low-mass star-forming regions (e.g., Bontemps et al. 1996, Hogerheijde et al. 1998) and in high-mass star-forming regions (e.g., Shepherd \\& Churchwell 1996, Zhang et al. 2001) showed that the detection rate is more than 80 \\%. This might suggest that the CO outflow detection rate in the OMC-2/3 region is significantly lower than in the other star-forming regions. We should note, however, that 23 of the 41 potential star-forming sites (49 \\%) do not have Spitzer 24 $\\mu$m source; these sites are the so-called prestellar cores, where no infrared embedded sources have been found. In contrast, the previous outflow searches mentioned above were biased toward IRAS sources. When we limit our regions to those with 24 $\\mu$m sources, then the CO outflow detection rate becomes 67 \\%. This is comparable to, but not as high as those of the previous searches. This is probably because; (i) a part of the ambient molecular cloud in the OMC-2/3 region shows CO emissions with higher velocities, which are difficult to distinguish from molecular outflows; (ii) some of the 24 $\\mu$m sources seem to be more evolved with less surrounding material.   \\subsubsection{Relation between outflow parameters and driving source properties} Previous studies have discussed the driving mechanisms and the physical properties of CO outflows (e.g., Lada 1985, Cabrit et al. 1992, Bontemps et al. 1996). Their main conclusions are as follows: (i) CO molecular outflows driven by a highly collimated flow (i.e., primary jet) are observed as entrained surrounding molecular gas, (ii) Energy of each CO outflow is mainly determined by the luminosity of the central driving source, which is probably related to the processes of mass accretion onto the central stars. In this subsection, we investigate the relationship between the physical parameters of the molecular outflows and of their driving sources.  In Figure \\ref{f11}, we plot the momentum flux of the outflow ($F_{\\rm{CO}}$) as a function of the bolometric luminosity of the outflow driving source ($L_{\\rm{bol}}$) for 10 sources from our sample. For comparison, Figure \\ref{f11} also shows the data of 101 CO outflows associated with young stellar objects (YSOs) in low- and high-mass star-forming regions, which were previously observed in CO with single-dish telescopes (Bontemps et al. 1996 and Hogerheijde et al. 1998 for low-mass star-forming regions, Beuther et al. 2002, Zhang et al. 2005 for high-mass star-forming regions).  Figure \\ref{f11} clearly shows a positive correlation between $F_{\\rm{CO}}$ and $L_{\\rm{bol}}$. The correlation is valid in a wide range of $L_{\\rm{bol}}$ (0.1 ${\\leq}L_{\\rm{bol}}{\\leq}~10^5{L_{\\odot}}$) and $F_{\\rm{CO}}$ ($10^{-6}{\\leq}{F_{\\rm{CO}}}{\\leq}0.1~M_{\\odot}$ km s$^{-1}$ yr$^{-1}$). Sources in low-mass star-forming regions have lower $F_{\\rm{CO}}$ and $L_{\\rm{bol}}$, while those in high-mass star-forming regions have higher $F_{\\rm{CO}}$ and $L_{\\rm{bol}}$. This conclusion has already been reported by Beuther et al. (2002). The sources in the OMC 2/3 region are located in this plot between the sources in low- and high-mass star-forming regions, and bridge a gap between them. The sources in the OMC-2/3 region have intermediate $F_{\\rm{CO}}$ and $L_{\\rm{bol}}$ in Figure \\ref{f11}, which naturally suggests that they are most probably intermediate-mass YSOs.  It is interesting to consider why both $L_{\\rm{bol}}$ and $F_{\\rm{CO}}$ can be good indicators of the mass of the YSOs. For low- and intermediate-mass YSOs, their bolometric luminosities are most probably dominated by the accretion luminosity ($L_{\\rm{acc}}{\\equiv}GM_{\\ast}{\\dot{M}_{\\rm{acc}}}/R_{\\ast}$). If we assume that the stellar radius ($R_{\\ast}$) does not significantly change during protostellar evolution, the relation suggests that a higher $L_{\\rm{bol}}$ indicates either a higher mass accretion rate or a higher central stellar mass, or both. Since a higher mass accretion rate results in a higher stellar mass as long as the mass accretion time scale is similar among low- and intermediate- mass protostars, it is fair to assume that $L_{\\rm{bol}}$ is an indicator for the mass of YSO for low- and intermediate-mass YSOs. For high-mass YSOs, on the other hand, it is not always true that $L_{\\rm{bol}}$ is dominated by the accretion luminosity since some of them reach the main sequence when they are still embedded in molecular clouds. If this is the case, $L_{\\rm{bol}}$ is basically the same as the stellar luminosity, indicating that $L_{\\rm{bol}}$ is an indicator of the stellar mass.  How about $F_{\\rm{CO}}$? It is not very clear how $F_{\\rm{CO}}$ can be an indicator of the mass of YSOs. Bontemps et al. (1996) suggested that $F_{\\rm{CO}}$ has a linear relationship with the rate of mass accretion onto the central star for low-mass YSOs. This is because the most plausible energy source for a jet/wind, which is considered to drive a CO outflow, is the gravitational energy released by infall and/or accretion onto the central protostar. If this is valid for low- to high-mass YSOs, $F_{\\rm{CO}}$ is indirectly related to the central stellar mass which is more or less determined by the mass accretion rate.  In Figure \\ref{f12}, we plot $F_{\\rm{co}}$ as a function of the envelope mass ($M_{\\rm{env}}$). In the same way as in Figure \\ref{f11}, we plot sources in the low- and high-mass star-forming regions (Bontemps et al. 1996, Hogerheijde et al. 1998, and Beuther et al. 2002) as well as those in the OMC-2/3 region. Figure \\ref{f12} clearly shows a good correlation between $F_{\\rm{co}}$ and $M_{\\rm{env}}$. Note that Bontemps et al. (1996) also showed the same correlation for low-mass YSOs, whereas Figure \\ref{f12} includes not only low-mass Class 0 and Class I sources, but also intermediate- and high-mass YSOs. Figure \\ref{f12} shows that more massive YSOs are associated with higher envelope masses, suggesting that $M_{\\rm{env}}$ is a good indicator of the YSO mass. This is probably not surprising because the core mass function often shows a relation similar to the initial mass function in various star-forming regions (e.g., Testi \\& Sargent, 1998, Motte et al. 1998, Nutter \\& Ward-Thompson 2007).  As discussed above, our sample sources are considered to be intermediate-mass YSOs. In the OMC2/3 region, however, low-mass YSOs should form as well (see Section 3.4). Why are no low-mass YSOs included in our sample sources? One probable reason is that our sample sources are selected based on 1.3 mm, 350 $\\mu$m, and 3.6 cm observations, as described in Section 3.1. The mass sensitivity (3 $\\sigma$) of the dusty envelope detected in the 1.3 mm observations is roughly 1 $M_{\\odot}$. For comparison, a typical value of  the envelope mass associated with low-mass protostars are 0.01 to 1 $M_{\\odot}$ (e.g., Bontemps et al. 1996, Hogerheijde et al. 1998; see Figure 12). Therefore, most low-mass candidates may have been missed in our sample sources. In other words, our observations were biased toward the intermediate-mass proto-stellar cores.  In addition to the discussion whether the outflow momentum flux depends on the central protostellar mass, it is also important to consider the evolutional effect on the outflow momentum flux. Bontemps et al. (1996) found a decline in the CO momentum flux from the Class 0 to Class I phase, suggesting that the outflow power declines with age. This result would be related to the decline of the mass accretion rate during the protostellar evolution. Our observational results were also examined from this point of view. In Figures 11 and 12, Class I sources in our sample were plotted as open diamonds while Class 0-like sources in our sample were plotted as filled diamonds. Figures 11 and 12 show no clear difference in the distribution of the momentum flux between Class 0 and Class I source in our samples. Why was no evolutionary trend seen in our samples? One possible reason is that the number of our sample sources may be too small to see such an evolutionary trend.  We should, however, note that there might be one example that shows a possible evolutionary trend. The source is SIMBA $a$, one of the Class 0-like sources in our sample. This source has particularly small momentum flux in comparison with the other sources in our sample, and seems not to follow the $M_{\\rm{env}}-F_{\\rm{CO}}$ correlation in Figure 12. Even though $M_{\\rm{env}}$ of SIMBA $a$ is as massive as those of our other samples, its $F_{\\rm{CO}}, 3.6 {\\times}10^{-6}~M_{\\odot}~km~s^{-1}~{yr^{-1}}$, is 1-2 orders of magnitude smaller than others. This source might be in the earliest stage of the protostellar evolution, in which the mass of the envelope is large whereas the accretion/outflow activity is still less."
    },
    "0808/0808.1421_arXiv.txt": {
        "abstract": "\\indent We present all-particle primary cosmic-ray energy spectrum in the $3\\cdot10^6-2\\cdot10^8$ GeV energy range obtained by a multi-parametric event-by-event evaluation of the primary energy. The results are obtained on the basis of an expanded EAS data set detected at mountain level ($700$ g/cm$^2$) by the GAMMA experiment. The energy evaluation method has been developed using the EAS simulation with the SIBYLL interaction model taking into account the response of GAMMA detectors and reconstruction uncertainties of EAS parameters. Nearly unbiased ($<5\\%$) energy estimations regardless of a primary nuclear mass with an accuracy of about $15-10\\%$ in the $3\\cdot10^6-2\\cdot10^8$ GeV energy range respectively are attained.\\\\ \\indent An irregularity (`bump') in the spectrum is observed at primary energies of $\\sim7.4\\cdot10^7$ GeV.  This bump exceeds a smooth power-law fit to the data by about 4~standard deviations. Not rejecting stochastic nature of the bump completely, we examined the systematic uncertainties of our methods and conclude that they cannot be responsible for the observed feature. ",
        "introduction": "\\indent Study of the fine structure in the primary energy spectrum is one of the most important tasks in the very high energy cosmic ray experiments \\cite{EW1}. Commonly accepted values of the all-particle energy spectrum indexes of $-2.7$ and $-3.1$ before and after the knee are an average and may not reflect the real behavior of the spectrum particularly after the knee. It is necessary to pay special attention to the energy region of $(1-10)\\cdot10^7$ GeV, where experimental results have been very limited up to now. Irregularities of the energy spectrum in this region were observed a long time ago. They can be seen from energy spectrum obtained more than 20 years ago with AKENO experiment \\cite{AKENO} as well as in later works of the GAMMA \\cite{GPRIM} and TUNKA \\cite{TUNKA} experiments. At the same time the large statistical errors did not allow to discuss the reasons of these irregularities. \\\\ \\indent On the other hand results of many experiments on the study of EAS charge particle spectra, the behavior of the age parameter and muon component characteristics point out that the primary mass composition at energies above the knee becomes significantly heavier. Based on these indications, additional investigations of the fine structure of the primary energy spectrum at $(1-10)\\cdot10^7$ GeV have an obvious interest.\\\\ \\indent There are two ways to obtain the primary energy spectra using detected extensive air showers (EAS). The first way is a statistical method, which unfolds the primary energy spectra from the corresponding integral equation set based on a detected EAS data set and the model of the EAS development in the atmosphere \\cite{KAS05,KAS07,G5a,G6,G6a}. The second method is based on an event-by-event evaluation \\cite{AKENO,G5b,G6b,SC,G05} of the primary energy of the detected EAS with parameters $\\mathbf{q}\\equiv q(N_e,N_{\\mu},N_h,s,\\theta)$ using parametric $E=f(\\mathbf{q})$ \\cite{AKENO,G5b,G6b,G05} or non-parametric \\cite{SC} energy estimator previously determined on the basis of shower simulations in the framework of a given model of EAS development.\\\\ \\indent Here, applying a new event-by-event parametric energy evaluation $E=f(\\mathbf{q})$, the all-particle energy spectrum in the knee region is obtained on the basis of the data set obtained with the GAMMA EAS array \\cite{G5a,G6,G6a,G6b} and a simulated EAS database obtained using the SIBYLL \\cite{SIBYLL} interaction model. Preliminary results have been presented in \\cite{G5b,G6b}. ",
        "conclusions": "% \\indent The multi-parametric event-by-event method (Sections 3,4) provides the high accuracy for the energy evaluation of primary cosmic ray nuclei $\\sigma(E)\\simeq10-15\\%$ regardless of the nuclei mass (biases$<5\\%$) in the $3-200$ PeV energy region. Using this method the all-particle energy spectrum in the knee region and above has been obtained (Fig.~9, Table~3) using the EAS database from the GAMMA facility. The results are obtained for the SIBYLL2.1 interaction model.\\\\ \\indent The all-particle energy spectrum in the range of statistical and systematic errors agrees with the same spectra obtained using the EAS inverse approach \\cite{G5a,KAS07,G6} in the $3-200$ PeV energy range.\\\\ \\indent The high accuracy of energy evaluations and small statistical errors point out at the existence of an irregularity (`bump') in the $60-80$ PeV primary energy region.\\\\ \\indent The bump can be described by 2-component model (parametrization (11)), of primary cosmic ray origin, where additional (pulsar) $Fe$ component are included with very flat power law energy spectrum ($\\gamma_{1p}\\sim1\\pm0.5$) before cut-off energy $E_{c,Fe}$ (Fig.~13, Table~4). At the same time, the EAS inverse problem solutions for energy spectra of pulsar component forces the solutions for the slopes of Galactic component above the knee to be steeper (Table~4), that creates a problem of underestimation of all-particle energy spectrum in the range of HiRes \\cite{HiResMIA} and Fly's Eye \\cite{FlySt} data at $E>200$ PeV. From this viewpoint the underestimation ($10-15\\%$) of the all-particle energy spectrum (bold solid line in Fig.~12) in the range of $E>200$ PeV can be compensated by the expected extragalactic component \\cite{Gaisser}.\\\\ \\indent Though we cannot reject the stochastic nature of the bump completely, our examination of the systematic uncertainties of the applied method lets us believe that they cannot be responsible for the observed feature. The indications from other experiments mentioned in this paper provide the argument for the further study of this interesting energy region."
    },
    "0808/0808.2034_arXiv.txt": {
        "abstract": "Recent results of the spectral and timing analysis of X-ray pulsars in hard X-rays with the INTEGRAL observatory are reviewed. The evolution of the cyclotron line energy with the source luminosity was studied in detail for the first time for several sources. It was shown that for V0332+53 this dependence is linear, but for 4U0115+63 and A0535+262 it is more complicated. There are some evidences of the ''reverse'' evolution for GX301-2 and Her X-1, and no evolution was found for Vela X-1, Cen X-3, etc. A strong dependence of the pulse fraction on the energy and source luminosity was revealed and studied in detail. A prominent feature in the pulse fraction dependence on the energy was discovered near the cyclotron frequency for several bright sources. The obtained results are compared with results of observations in standard X-rays and briefly discussed in terms of current models; some preliminary explanations are proposed. ",
        "introduction": "According to the simple theory of accretion onto a rapidly rotating neutron star with a strong magnetic field, the matter from the normal star is stopped near the Alfven surface by the magnetic field pressure, trapped into it and moved along the field lines to small ringing regions on the neutron star surface, emitting X-rays. Due to the dipole structure of the neutron star magnetic field two hot spots will be formed on the surface. In such a case, for a rotating neutron star, the observer will detect on the source light curve pulses of different forms depending on the physical and geometrical conditions as near the neutron star surface, where these pulses are formed, as on the line of sight. The corresponding pulse fraction depends on the configuration of the emitted regions, relative position of the dipole to the observer and the energy.  As it was shown in several early papers \\cite{wangwel81,white83}, observed pulse profiles are essentially differed for different sources; moreover, they demonstrate the wide variety of the forms depending on the energy and can shift, in some cases, up to 180\\deg on the pulse phase. There is also a strong pulse profile dependence on the source luminosity, which can be connected with changes of the beam function from the pencil-beam to the fan-beam ones with increasing luminosity from $<10^{37}$ to $several\\times10^{38}$ ergs s$^{-1}$ \\cite{bs1976}. The presence of the strong ($10^{11}-10^{13}$ G) magnetic field near the emitting regions on the surface of the neutron star can be an origin of additional peculiarities in the observed properties of X-ray pulsars, e.g. the presence of the cyclotron resonance absorption features in the X-ray pulsar spectra, whose energy is directly connected with the value of the magnetic field. Moreover, properties of the accreting plasma are significantly changed at the cyclotron frequency, that can leads to changes in the emission beam function \\cite{mesnag81}. The corresponding changes of the pulse profile near the cyclotron energy were detected from several sources (see, e.g., \\cite{tsy06}).  Obviously, such strong changes in the beam function, conditions and geometry of the emission formation regions should lead to strong dependence of the pulse fraction on the energy and source luminosity. This fact was recognized as early as 80-90's, but only with launches of the orbital observatories RXTE and INTEGRAL, which have high timing and energy resolution (especially at high energies, where the observed emission isn't subject of the influence of photoabsorption and depends only on the system geometry and physical conditions in the formation regions) has it been possible to carry out systematic investigations in this field. In particular, Tsygankov et al. \\cite{tsy07} showed that the pulse fraction of the X-ray pulsar 4U0115+634 is decreased with the source luminosity increase and increased with the energy, having maximums near harmonics of the cyclotron absorption line. The pulse fraction increase with the energy was also found for several other X-ray pulsars (see, e.g.  \\cite{fer07,barn08}). Using INTEGRAL data, Tsygankov \\& Lutovinov \\cite{tsylut08} studied in detail the pulse fraction dependence on the energy and luminosity for ten bright X-ray pulsars in hard X-rays.  Here we summarize and briefly review current results of observations of X-ray pulsars with the INTEGRAL observatory, drawing the main attention to the cyclotron line energy, pulse profiles and pulse fraction dependences on the luminosity and energy band. ",
        "conclusions": "\\begin{itemize}  \\item Most accurate results of the spectral and timing analysis of several dozens X-ray pulsars in hard X-rays are obtained at the moment with the INTEGRAL and RXTE observatories.  \\item The evolution of the cyclotron energy with the source luminosity was studied in detail for the first time. It was shown that for V0332+53 this dependence is linear (the cyclotron line energy is increased with the luminosity decrease), but for 4U0115+63 and A0535+262 it is more complicated; the ''reverse'' behaviour is detected for Her X-1 and GX301-2.  \\item The strong dependence of the pulse fraction on the energy and source luminosity is revealed and study in detail.  \\item The prominent feature in the dependence of the pulse fraction on energy band was revealed near the cyclotron frequency for several bright sources.  \\item There are preliminary ideas how to explain the above, but the additional theoretical studies and simulations are needed for resolving the puzzles.  \\end{itemize}   \\begin{theacknowledgments}  We thank J.Poutanen, M.Revnivtsev and V.Suleimanov for the useful discussion during meetings in the ISSI (Bern), whose support and hospitality we also acknowledge. This work was supported by the Russian Foundation for Basic Research (Projects No. 07-02-01051 and 08-08-13734), the Presidium of the Russian Academy of Sciences (The Origin and Evolution of Stars and Galaxies Program), and the Program of the President of the Russian Federation (Project No. NSh-5579.2008.2).   \\end{theacknowledgments}"
    },
    "0808/0808.0341_arXiv.txt": {
        "abstract": "Several anomalies appear to be present in the large-angle cosmic microwave background (CMB) anisotropy maps of WMAP.  One of these is a lack of large-scale power.  Because the data otherwise match standard models extremely well, it is natural to consider perturbations of the standard model as possible explanations. We show that, as long as the source of the perturbation is statistically independent of the source of the primary CMB anisotropy, no such model can explain this large-scale power deficit. On the contrary, any such perturbation always {\\it reduces} the probability of obtaining any given low value of large-scale power. We rigorously prove this result when the lack of large-scale power is quantified with a quadratic statistic, such as the quadrupole moment.  When a statistic based on the integrated square of the correlation function is used instead, we present strong numerical evidence in support of the result. The result applies to models in which the geometry of spacetime is perturbed (e.g., an ellipsoidal Universe) as well as explanations involving local contaminants, undiagnosed foregrounds, or systematic errors. Because the large-scale power deficit is arguably the most significant of the observed anomalies, explanations that worsen this discrepancy should be regarded with great skepticism, even if they help in explaining other anomalies such as multipole alignments. ",
        "introduction": "Observations of cosmic microwave background (CMB) anisotropy, particularly the data from WMAP \\cite{wmap1yr1,wmap1yr2,wmap1,wmap5yrbasic}, have revolutionized cosmology.  These observations are a major contributor to the emergence of a cosmological ``standard model'' of a Universe dominated by dark energy and cold dark matter, with a nearly scale-invariant spectrum of Gaussian adiabatic perturbations \\cite{wmap5yrparams,wmap5yrinterp}.  The overall consistency of the CMB data with this model is quite remarkable, but there appear to be some anomalies on the largest angular scales, such as a lack of large-scale power \\cite{wmap1yr2,copi2,dOCTZH}, alignment of low-order multipoles \\cite{schwarz,copi1,dOCTZH,hajian}, and hemispheric asymmetries \\cite{eriksen2004,freeman}.  The significance of and explanations for these puzzles are hotly debated.  In particular, it is difficult to know how to interpret a posteriori statistical significances: when a statistic is invented to quantify an anomaly that has already been noticed, the low $p$-values for that statistic cannot be taken at face value. Nonetheless, the number and nature of the anomalies (in particular, the fact that several seem to pick out the same directions on the sky) seem to suggest that there may be something to explain in the data. In this paper, we will tentatively assume that there is a need for an explanation and consider what that explanation might be.  Since the standard model is in general highly consistent with the CMB and a wide variety of other observations, it is natural to look for explanations of these puzzles that consist of perturbations added onto the standard model.  Such explanations can be based on nonstandard cosmologies, such as ellipsoidal models \\cite{ellipsoid}, large-scale magnetic fields \\cite{barrow}, and theories based on Bianchi VIIh spacetimes with rotation and shear \\cite{ghosh,bianchi}. They can also involve phenomena on much smaller scales (e.g., \\cite{inoue, inoue2,abramo,cooray}), perhaps even within the Solar System \\cite{frisch}. Any uniagnosed foreground contaminant would fall into the class of explanations we consider, as would many systematic errors.  All of these models can be described by assuming that the observed CMB sky is the sum of two terms: \\beq T_{\\rm obs}({\\bf r})=T_0({\\bf r})+T_c({\\bf r}), \\eeq where $T_0$ is a Gaussian CMB sky with a power spectrum given by the standard model and $T_c$ is a contaminant. The contaminant can be a fixed function of sky position $\\bf r$ or a realization of a random process.  In the latter case, we assume nothing about the statistics of this process except that it is independent of the Gaussian random process that produced $T_0$. We wish to consider the possibility that such a model can explain some or all of the large-angle anomalies.  In this article, we will present strong evidence that on the contrary all such models actually exacerbate one of the anomalies, namely the observed lack of power in the large-angular-scale CMB anisotropy. This anomaly is formally highly statistically significant, and as we will argue below it is one for which the problems of a posteriori statistics are not particularly severe.  It is therefore arguably the most in need of explanation of all of the large-angle CMB puzzles.  We conclude, therefore, that this entire category of possible explanations should be regarded with great skepticism.  In particular, the absence of large-scale power in the WMAP data is in fact a strong argument {\\it against} the existence of undiagnosed foreground contamination,  as well as systematic errors that would produce an additive contaminant to the observed sky maps.  We can quantify the lack of power in the large-angle CMB by considering either the power in low-order multipoles (especially the quadrupole) or a statistic based on the two-point angular correlation function (see Fig.~\\ref{fig:corrfunc}). In either case, the $p$-values (that is, the probabilities of getting as low a value of the chosen statistic as the one in the actual data) are low; in fact, for some choices of statistic, they are less than 0.1\\% \\cite{copi2} (but see \\cite{efstathiou2} for a constrasting analysis). By definition, for an alternative theory to explain this anomaly, it would have to generate larger $p$-values.  We will show in this paper that all proposed models of the form described above in fact {\\it reduce} the $p$-values based on these statistics. Therefore, although such models might alleviate some of the other large-angle anomalies, they worsen this one.   \\begin{figure} \\includegraphics[width=3.5in]{corrplot.eps} \\caption{The two-point correlation function for the WMAP data.  The solid curve shows the correlation function for the three-year WMAP internal linear combination (ILC) data at resolution $N_{\\rm side}=32$, computed with the SpICE software \\cite{spice}.  The dashed curves show the 95\\% confidence range in a set of simulations.  The simulations virtually never produce correlation functions that are as close to zero at large angles as the real data.} \\label{fig:corrfunc} \\end{figure}   At one level, this is not surprising.  For the models considered here, in which the observations are the sum of two statistically independent terms, the observed power spectrum is simply the sum of the standard-model spectrum and the spectrum of the contaminant.  Addition of the contaminant therefore biases all multipoles up, including the quadrupole.  This is merely a statement about mean-square values, however, and does not tell us about the probability distribution of the multipoles. It is logically possible that a (non-Gaussian) contaminant, even as it biases the mean-square quadrupole up, widens the probability distribution for the quadrupole in such a way as to enhance the probability of getting low values. Indeed, any proposal to explain the lack of large-scale power through a perturbation to the standard model must be proposing such an effect, since this is what it would mean to ``explain'' the discrepancy.  For example, suggestions have been made that the low quadrupole might be explained by an extended local foreground \\cite{abramo}, by dust-filled local voids\\footnote{A suggestion is made in the cited work that the hypothesis that the contaminant is uncorrelated with the primary signal may not apply.  If this is true, then the arguments in the present paper would not apply to this model.  It is not clear to us that a strong correlation of the proposed form exists in the model under consideration, and as far as we know no detailed calculation of this effect has been performed.} \\cite{inoue,inoue2}, or by an ``ellipsoidal'' universe that expands at different rates in different directions \\cite{ellipsoid}. Each such explanation assumes that a chance anticorrelation between the contaminant and the intrinsic CMB anisotropy has occurred.  In order for this to count as an explanation, however, such an anticorrelation must be sufficiently probable that it raises the probability of finding the observed lack of power. Although this is a logical possibility, we will argue below that it in fact never occurs, whether the lack of power is quantified via the quadrupole moment or the correlation function. For some specific cases, such as the quadrupole moment in an ellipsoidal universe, previous work \\cite{ellipsoidiswrong} has already established this; in this paper we prove it in general. In summary, such models cannot explain the lack of large-scale power, and in fact always ``anti-explain'' it by reducing the already-low probability.  Section \\ref{sec:quadratic} proves this general result in the case where the lack of power is quantified via a quadratic estimator such as the mean-square quadrupole moment.  Section \\ref{sec:correlation} presents strong numerical evidence that the result is also true in the case of a statistic based on the two-point correlation function. Section \\ref{sec:discussion} contains a brief discussion of the results, and an appendix proves a key mathematical result needed in section \\ref{sec:quadratic}.  \\begin{figure*}[t] \\includegraphics[width=3.5in]{cl2ellipticitynocut.eps} \\includegraphics[width=3.5in]{cl2ellipticitykp0.eps} \\caption{ Cumulative probability distributions for the quadrupole power $C_2$ in an ellipsoidal Universe, in dimensionless $\\Delta T/T$ units.  In the left panel, data from the entire sky were used, and in the right panel the WMAP Kp0 cut was applied.  The solid curve shows a standard LCDM model with no eccentricity.  From bottom to top, the other three curves correspond to eccentricities $5\\times 10^{-3},6.2\\times 10^{-3},7.4\\times 10^{-3}$. The horizontal line shows the quadrupole found in the actual data. } \\label{fig:C2} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  We have considered a broad class of cosmological models, obtained by adding a contaminant to the standard best-fit inflation-based model. The only assumption we have made about the contaminant is that it is statistically independent of the cosmological signal.  We have argued that all such models exacerbate rather than alleviating the lack of large-scale power in the WMAP data.  We hve proven this result to be true when the lack of power is quantified by the quadrupole moment and have presented strong numerical evidence in support of it when the two-point correlation function is used.  Since the latter in particular is discrepant at a highly significant level already, any theory that worsens this discrepancy should be regarded with great skepticism.   \\begin{figure} \\includegraphics[width=3.5in]{sprobpatternkp0.eps} \\caption{Cumulative probability distributions for models in which a fixed contaminant of the form shown in Fig.~\\ref{fig:testpattern} is added to the sky.  This pattern corresponds to the most negative eigenvalue of the matrix ${\\bf H}$ and so might be expected to increase the probability of finding a low value of $S_{1/2}$.  The curves shown in the figure are for varying amplitudes of the contaminant, with root-mean-square pixel values of 0 (solid) 2\\,$\\mu$K (dotted), 4\\,$\\mu$K (dashed), 8\\,$\\mu$K (dot-dashed). No value of the amplitude causes the probability of getting values as low as those in the real data to increase noticeably.} \\label{fig:sprobpattern} \\end{figure}  In addition to exotic cosmologies such as models with a global ellipsoidal anisotropy, the class of models considered herein includes more mundane possibilities such as undiagnosed foregrounds and many systematic errors.  In particular, since several of the observed anomalies seem to ``pick out'' the ecliptic plane as a preferred direction, some attention has focused on a local foreground as a possible explanation.  The calculations presented here argue against such models.  In any particular model with a contaminant, of course, it is possible that a chance cancellation between the contaminant and the intrinsic CMB anisotropy can occur, leading to the observed lack of large-scale power.  What we have shown is that such a chance cancellation is always unlikely, and in particular that it is always more unlikely than the lack of large-scale power occurring based on cosmic variance alone, without a contaminant.  The question of how seriously to take the various large-angle CMB anomalies, including the lack of large-scale power as well as the various other puzzles, has been much debated \\cite{efstathiou}.  In particular, because they are all based on a posteriori statistics ({\\it i.e.}, on statistical significances calculated after the anomalies had already been noted), the quoted significances cannot be taken at face value.  Arguably, however, the large-scale power deficit suffers less from this problem of a posteriori statistics. After all, for virtually the entire existence of the field of CMB anisotropy studies, the two-point correlation function has been regarded as one of the most natural statistics to use in quantifying the level of structure in CMB maps as a function of angular scale. For instance, upper limits on CMB anisotropy in the pre-COBE era were usually presented as limits on the correlation function.  Although the particular statistic $S_{1/2}$ is an a posteriori invention, it merely quantifies the mean-square level of this function, which was already regarded a priori as a natural function to compute.  Although one can certainly dispute the extent to which the significance of the lack of large-angle correlations is an artifact of the particular choice of statistic (for instance, see \\cite{gaztanaga}, who do not use the $S_{1/2}$ statistic and find less significant discrepancies), nonetheless we believe that, of all the observed anomalies, the large-scale power deficit is one of the most in need of explanation.  Anomalies in a data set naturally prompt thoughts of systematic errors or contaminants in the data.  Perhaps counterintuitively, this particular anomaly provides a strong argument against such possibilities.  In particular, a foreground that was not removed from the data (due to having a spectrum indistinguishable from the CMB, for example) would fall precisely into the category considered herein.  Note, however, that if the foreground removal procedure itself removes part of the cosmological signal, the resulting error would not fall into the category considered herein. In particular, the ILC method does project out some of the intrinsic CMB signal and so in principle does reduce the amount of large-scale power. This effect is calculable and has been found to be negligible, however.  There are of course a wide variety of possible explanations for the anomalies that do not fall into the category considered here.  For example, simply modifying the primordial power spectrum at large scales naturally alleviates the problem of a lack of large-scale CMB power (e.g., \\cite{cline,contaldi}). Some models with nontrivial topology (e.g., \\cite{starobinsky,stevens,angelica,bonehead,luminet}) also have this effect, although such models have other problems \\cite{circles}. The framework of spontaneous isotropy breaking \\cite{gordon} also provides a class of models that are not based on simply adding a perturbation to the standard cosmology.  Models such as these (and many others) may provide an explanation for the puzzles in the large-angle CMB."
    },
    "0808/0808.0494_arXiv.txt": {
        "abstract": "Recently, rapid observational and theoretical progress has established that black holes (BHs) play a decisive role in the formation and evolution of individual galaxies as well as galaxy groups and clusters. In particular, there is compelling evidence that BHs vigorously interact with their surroundings in the central regions of galaxy clusters, indicating that any realistic model of cluster formation needs to account for these processes. This is also suggested by the failure of previous generations of hydrodynamical simulations without BH physics to simultaneously account for the paucity of strong cooling flows in clusters, the slope and amplitude of the observed cluster scaling relations, and the high-luminosity cut-off of central cluster galaxies. Here we use high-resolution cosmological simulations of a large cluster and group sample to study how BHs affect their host systems. We focus on two specific properties, the halo gas fraction and the X--ray luminosity-temperature scaling relation, both of which are notoriously difficult to reproduce in self-consistent hydrodynamical simulations. We show that BH feedback can solve both of these issues, bringing them in excellent agreement with observations, without alluding to the `cooling only' solution that produces unphysically bright central galaxies. By comparing a large sample of simulated AGN-heated clusters with observations, our new simulation technique should make it possible to reliably calibrate observational biases in cluster surveys, thereby enabling various high-precision cosmological studies of the dark matter and dark energy content of the universe. ",
        "introduction": "Numerous observational and theoretical studies of galaxy clusters and groups show that the astrophysical processes relevant for determining their properties are still poorly understood. Hydrodynamical simulations of cluster formation which only include radiative cooling processes \\citep[e.g.][]{Lewis2000, Muanwong2001, Yoshida2002} suffer from excessive overcooling within the densest cluster regions, where gas cooling times are short. This translates into a very large fraction of cold gas and consequently a large amount of stars that would form out of it, reaching at least $40-50\\%$ of the baryonic cluster content. While the associated removal of low-entropy gas from cluster centers breaks the self-similarity of the cluster scaling relations in a way that resembles observations \\citep{Bryan2000, Dave2002, Nagai2007}, this `cooling-only' scenario is generally discounted because observed clusters do not show such strong cooling flows and enormously bright central galaxies. Instead, it is widely believed that some non-gravitational energy input must strongly affect the physics of the intracluster medium (ICM).   There has been considerable effort \\citep[see][and references therein]{Borgani2004} to include feedback mechanisms associated with star formation in hydrodynamical simulations in order to reduce excessive overcooling in cluster cores. However, so far simulation models have failed to simultaneously reproduce the masses and colors of central galaxies, the observed temperature and metallicity profiles, as well as the observed X--ray luminosity-temperature ($L_{\\rm X}-T$) relation. In particular, the X--ray luminosities have been found to be substantially larger than the observed values \\citep{Borgani2004} on the group scale. Furthermore the observed baryon fractions in clusters and groups are typically smaller than in simulations, which for `standard' physics invariably predict a value close to the universal cosmic baryon fraction \\citep[e.g.][]{Frenk1999,OShea2005,Kravtsov2005}.   Several attempts have been made to tackle these problems with so-called `preheating' schemes \\citep[e.g.][]{Bialek2001,Borgani2005}, that are meant to mimic more energetic astrophysical processes than the relatively inefficient supernovae feedback. However, the preheating scenarios are physically not well motivated and rely on an ad-hoc choice of the amount and epoch of energy injection into the ICM.  On the other hand, it is becoming increasingly clear that active galactic nuclei (AGN) play an important role for understanding the properties of clusters and groups \\citep[see ][for a recent review]{McNamara2007}. Also analytic work \\citep[e.g.][]{Valageas1999, Bower2001, Cavaliere2002} suggests that including self-consistent feedback from AGN in simulations might resolve the discrepancies by providing a heating mechanism that offsets cooling, lowers star formation rates and removes gas from the centers of poor clusters and groups, thereby reducing their X--ray luminosities to values compatible with the observed $L_{\\rm X}-T$ relation.  In this work we investigate whether AGN feedback can indeed solve these problems by performing hydrodynamical simulations of galaxy clusters and groups that employ a state-of-the-art model \\citep{Sijacki2007} for BH growth and associated feedback processes. ",
        "conclusions": "\\label{SecConclusions}  We have performed very high-resolution numerical simulations of a mass-selected sample of galaxy clusters and groups. In order to investigate whether AGN feedback makes simulations and X--ray observations of cluster and groups compatible, we have carried out two kinds of simulations for each halo, namely (1) hydrodynamical simulations with a treatment of radiative cooling, star formation, supernovae feedback, and heating by a UV background, and (2) simulations that additionally employ a model for black hole growth and associated feedback processes. Our main findings are: \\begin{itemize} \\item AGN feedback significantly lowers the gas mass fractions in poor clusters and groups, even though fewer baryons are turned into stars at the same time. This is because the AGN heating drives gas from halo centers to their outskirts. In massive clusters, on the other hand, it mainly lowers the central gas density and substantially reduces the amount of stars formed. Overall, both the gas and stellar fractions of the whole sample of our simulated groups and clusters are in a much better agreement with observations when AGN are included than in simulations without AGN. \\item AGN feedback significantly reduces the X--ray luminosities of poor clusters and groups, while the X--ray temperature within $r_{\\rm 500}$ stays roughly the same or is even slightly reduced. This results in a steepening of the $L_{\\rm X}-T$ relation on the group scale. \\item The $L_{\\rm X}-T$ relation obtained from simulations with AGN feedback is in excellent agreement with observations at all mass scales. \\end{itemize}  We find it extremely encouraging that this simple model for BH growth and feedback is capable of bringing the analyzed properties of galaxy clusters and groups into much better agreement with observations. This not only resolves a long-standing problem in their hydrodynamical modeling, but also opens up new exciting possibilities for using numerical simulations to investigate the properties of clusters and groups, and their co-evolution with central supermassive BHs. Also, it considerably brightens the prospects to use simulations of cluster formation to accurately calibrate and correct systematic effects in future X-ray and Sunyaev-Zel'dovich cluster surveys. This is essential for the exploitation of the full potential of clusters as cosmological probes to accurately constrain the expansion history of the Universe."
    },
    "0808/0808.3668_arXiv.txt": {
        "abstract": "We present results on a survey to find extremely dust-reddened Type-1 Quasars. Combining the FIRST radio survey, the 2MASS Infrared Survey and the Sloan Digital Sky Survey, we have selected a candidate list of 122 potential red quasars. With more than 80\\% spectroscopically identified objects, well over 50\\% are classified as dust-reddened Type 1 quasars, whose reddenings (E(B-V)) range from approximately 0.1 to 1.5 magnitudes. They lie well off the color selection windows usually used to detect quasars and many fall within the stellar locus, which would have made it impossible to find these objects with traditional color selection techniques. The reddenings found are much more consistent with obscuration happening in the host galaxy rather than stemming from the dust torus. We find an unusually high fraction of Broad Absorption Line (BAL) quasars at high redshift, all but one of them belonging to the Low Ionization BAL (LoBAL) class and many also showing absorption the metastable FeII line (FeLoBAL). The discovery of further examples of dust-reddened LoBAL quasars provides more support for the hypothesis that BAL quasars (at least LoBAL quasars) represent an early stage in the lifetime of the quasar. The fact that we see such a high fraction of BALs could indicate that the quasar is in a young phase in which quasar feedback from the BAL winds is suppressing star formation in the host galaxy. ",
        "introduction": "\\defcitealias{f2m07}{F2M} \\defcitealias{f2m04}{F2M}  Since the identification of 3C 273 \\citep{3c273}, our understanding of quasars has evolved tremendously. There exists now a relatively robust quasar model in which different features of a quasar are explained invoking an orientation-model \\citep{antonucci,urry}. The orientation model describes the quasar by an accretion disk at the center and broad emission line clouds directly above and below the disk. The broad-line clouds can be shielded from our line of sight by a dusty donut-shaped torus coplanar with the disk. One explains whether or not an object is observed to have broad lines in the rest-frame optical wavelengths with the viewing angle of the AGN. This model has also been successful in explaining the shape of the Extragalactic X-ray background spectrum, which peaks at $\\sim$ 30keV \\citep{treister, gilli}. There must be a large population of obscured quasars contributing to it to explain its spectrum peaking at such high energies.  The discovery of quasars through color selection techniques based on the typically blue spectrum of quasars has been extremely successful. It has resulted in a large output of optical spectra such as from the Sloan Digital Sky Survey (SDSS, \\cite{sc05}), but also in quasar identifications in deep multiwavelength surveys (e.g. GOODS \\citep{cdfs,goodsir} and NOAODWFS \\citep{noaodwfs}).  By 2008 there are well over 100,000 quasars known, with another 100,000 likely candidates \\citep{richards04}. However, the question of completeness of the current samples of quasars remains. Recent surveys have shown that optically-selected quasars comprise less than half of the total quasar population \\citep{martinez05,stern,alonso}. Optical quasar surveys tend to miss dust-reddened quasars that have been found either with infrared surveys \\citep{cutri,lacy,f2m04,f2m07}, radio surveys \\citep{postman}, hard X-ray surveys \\citep{polletta} or surveys for high-ionization narrow-line objects \\citep{zakamska}. Furthermore, the distinction of unobscured / obscured has been proven to not be necessarily equivalent to Type 1 / Type 2, as many dust-reddened and obscured quasars have been found with broad lines or featureless spectra \\citep{alonso,martinez06}. Also, due to polarized scattering material in the polar directions the Type 1 / Type 2 spectral features as an indicator of orientation may be ineffective; quasars with high dust-covering factors will be always appear as Type 1 objects optically, regardless or orientation \\citep{scatter}. In particular, broad absorption line quasar (BAL quasar) samples have redder colors than typical quasars \\citep{sp92,reich} and higher X-ray column densities \\citep{gal} also suggesting that a large fraction of this population is missing from optically selected samples.  This discrepancy in type of obscuration has been observationally recorded, especially in objects with different X-ray/mid-IR properties. In some AGN which display Type 2 column densities in the X-rays, the mid-infrared dust abundances are not correlated with those columns \\citep{rigby}. Furthermore, there are many high luminosity objects in the mid-infrared with AGN signatures, but with a prominent Si absorption feature at 10$\\mu$m, that do not exhibit strong absorption in the X-rays \\citep{hao}.  Current numerical simulations suggest that AGN spend a large phase of their lifetime enshrouded by dust from the merging host galaxies \\citep{hopkins05,li07}. This scenario has already been suggested in earlier studies, based on mid- and far-infrared observational evidence \\citep{sanders}. AGN in a growing phase are thought of as relatively luminous and heavily absorbed. Eventually the intense radiation field from the central engine destroys the surrounding obscuring material and the AGN shines through unabsorbed \\citep{dimatteo}. These studies suggest that a non-negligible fraction of AGN are obscured by dust in the young merging host galaxy rather than a torus. In this picture, the quasar is also responsible for quenching some amount of the massive star formation occurring in the bulge of the host galaxy. Although there is a strong debate over how much influence this ``quasar feedback'' has on the host galaxy, these scenarios often invoke strong quasar winds that persist well outside the galaxy. BALs, the quasar phenomenon that we associate with very strong quasar winds, do not seem to have an evolutionary nature, but rather are explained by an ``orientation hypothesis'' \\citep{weymann,gal07}. However, the special cases of LoBALs (low ionization BALs, showing broad absorption troughs in the MgII line) and especially FeLoBALS (with absorption troughs in the metastable FeII line) are often invoked as classes of BALs that might be explained by a ``youth hypothesis'' in which BALs are at an early stage of their evolution in their quasar lifetime \\citep{voit,bobbal,Trump}. Early studies of LoBALs show them to have larger dust covering fractions than standard BAL QSOs and ``normal'' QSOs \\citep{boroson}. Also for the FeLoBALs, radiative transfer spectral synthesis models show them to have large covering fractions for their outflows \\citep{casebeer} and MIPS photometry shows them to be very luminous in the mid-IR, probably due to star formation \\citep{farrah}. Unfortunately, so far there is not conclusive evidence for PAH emission signifying star formation in BAL quasars, but a comprehensive unbiased sample, especially of LoBALs, is missing \\citep{shi06,shi07}.  Finding dust-obscured quasars in the optical has proven to be a difficult task, as dust suppresses mostly UV and optical light. Since we want to study quasars at high redshift where the star formation rate and quasar fraction peaks, the task is even more difficult, as the optical regime now covers mostly rest-frame UV light that is suppressed by several magnitudes from the dust. Nevertheless (\\cite{f2m04,f2m07}; hereafter F2M) conducted a survey to find these dust-reddened and -obscured quasars using the 2MASS near infrared survey \\citep{2mass} and the FIRST radio survey \\citep{first} and following up spectroscopically. This F2M survey largely concluded that UVX- and optical color-selected sample in the optical miss $\\ge$ 10\\% of the quasar population and that the red quasar population constitutes $>$ 20\\% of the total quasar population.  In this paper we describe a follow up to the FIRST-2MASS survey described in \\citetalias{f2m07} to find dust-reddened quasars using the FIRST radio survey \\citep{first} and the 2MASS Infrared Survey \\citep{2mass} but with the additional photometry of the SDSS. It expands the original survey area, but only explores the objects with the reddest r'-K colors. Throughout this paper we use a flat universe with $H_0$ = 70 km$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}$ = 0.7 cosmology. ",
        "conclusions": "We have expanded the original FIRST-2MASS red quasar survey to include the whole extent of the 2MASS point source catalog but only include objects which have a very red color (r'-K $>$ 5). This yielded 122 red quasar candidates. Spectroscopic follow-up on 100 of the sources revealed that well over 50\\% are red Type-1 QSOs, consistent with the original FIRST-2MASS Red Quasar survey.  As described in Section \\ref{specobs} the deduced reddening is most likely due to dust absorption. The question remains where the obscuring dust is located. It could be in regions of the torus, where sublimation does not affect it, or it could be in a cold outside absorber, i.e. the host galaxy. Some clues as to the nature of the extinction in these objects came from our HST/ACS study of a sample of 13 F2M quasars (those in the F2MS sample are labeled in Table \\ref{observe}). The images showed impressive evidence for a recent merger in an unusually high percentage (85\\%) of the red quasars. Also, there was a weak correlation between the reddening and the amount of interaction displayed in the host galaxy measured by the Gini coefficient. These dust-reddened quasars may be young, recently ignited objects, where quasar feedback has yet to quench the star formation activity \\citep{redqso-hst}. Ongoing Spitzer IRS and MIPS observations of the HST/ACS sample should help to test this hypothesis, as we will be able to measure PAH strengths and 70$\\mu$m dust emission likely associated with young star forming host galaxies. The Spitzer IRS spectra will also describe the absorbing dust properties better by analyzing the 9.7 $\\mu$m Silicate absorption feature and relating that to our deduced reddening.  There is still debate over how biased our surveys are. Radio-selected samples tend to have redder colors. However, IR selected samples tend to have larger fractions of BALQSOs \\citep{dai}. There has been some debate over whether selection biases in the optical make one underestimate the fraction of BALQSOs \\citep{krolik}. Essentially none of the BALs discovered with the F2MS survey would have been discovered with the SDSS color selection techniques (except for serendipitous spectra or searches for high redshift quasars). Hence, our deduced fraction of BALs is too high (63\\%) and even the conservative fraction of (40\\%) deduced by \\cite{dai} is quite large. It becoms difficult to explain these nature objects as a pure orientation effect as the wind opening angle would be very large. However, since we are biased in only looking at dust-reddened objects, the dust could be present in the BAL clouds themselves and we would find naturally only find such a large number of BALs in objects that are reddened. Also, in the orientation picture, the highly collimated BAL winds could be grazing the dust torus, meaning that they are naturally redder than usual. These caveats aside, we believe that with such large dust reddenings, the evidence for mergers at lower redshift for similarly selected objects, and the large intrinsic fraction lead us to believe that BALQSOs are in an early evolutionary phase in the lifetime of a quasar. Their powerful winds might be the quenchers of star formation in the host galaxy, often invoked as quasar feedback (e.g. \\cite{sr98,kauffmann,li07}).  The ``youth hypothesis'' for LoBALs fits very well with the picture of dust-reddened quasars being an evolutionary step in the lifetime of a quasar. The results from the hydrodynamic simulations predict that there is some amount of quasar feedback seen as a kind of ``blowout'' just after the phase where the merger has coalesced and the quasar is starting to ignite (albeit, enshrouded in dust) \\citep{hop07}. Whether this quasar feedback is actually done by a quasar wind that produces BAL-Type spectra is not clear yet though there is some evidence that LoBAL winds can extend into the host galaxy \\citep{dekool}. Recent poststarburst galaxies also show evidence for massive outflows which must have originated in a quasar when the starburst was still young \\citep{tremonti}."
    },
    "0808/0808.3314_arXiv.txt": {
        "abstract": "{ {We present a possible scenario for the ejection of a superluminal component in the jet of the Broad Line Radio Galaxy \\object{3C~111} in early 1996. VLBI observations at 15 GHz discovered the presence of two jet features on scales smaller than one parsec. The first component evolves downstream, whereas the second one fades out after 1 parsec.} {We propose the injection of a perturbation of dense material followed by a decrease in the injection rate of material in the jet as a plausible explanation.} {This scenario is supported by 1D relativistic hydrodynamic and emission simulations. The perturbation is modeled as an increase in the jet density, without modifying the original Lorentz factor in the initial conditions.} {We show that an increase of the Lorentz factor in the material of the perturbation fails to reproduce the observed evolution of this flare. We are able to estimate the lifetime of the ejection event in 3C~111 to be $36\\pm7$~days.} {} ",
        "introduction": "\\label{intro} Flaring events at radio frequencies are known to take place in Active Galactic Nuclei (AGN), usually followed by the observation of new radio features in the parsec-scale jets \\citep[e.g.,][]{sa02}. It has been shown that the ejection of those features, or components, is related to dips in the X-ray emission from the active nucleus in the case of 3C~120 \\citep{Mar02}, and perhaps also in 3C~111 \\citep{Mar06}. The dips in X-rays precede the observations of new radio-components. The decrease in X-ray emission may be caused by the loss of the inner regions of the disc. In this scenario, a fraction of the accreted material is injected in the jet and a new component is later observed in VLBI images, after the material becomes detectable at the observing frequencies, as it evolves downstream. The components are interpreted as the shocks produced by the ejection of denser and/or faster plasma in the flaring event from the accretion disc \\citep{mar85}. The conditions for triggering the ejection of the material in those radio features are still unknown. In the case of microquasars, it has been proposed that the stronger components are ejected right before the passage of the source from the X-ray hard/low state, associated with higher radio brightness, to the soft/high state, associated with lower radio emission -- a decrease in the injection of emitting particles in the jet \\citep[][ and references therein]{fb04}.  If we interpret jet components as shocks propagating in a supersonic flow, their origin must be related to an increase of pressure and/or velocity in an injected perturbation with respect to the steady initial flow. Hydrodynamical simulations \\citep[][ A03 hereafter]{Alo03} have shown that such jet perturbations produce a forward and a reverse structure, which would be expected to be observed as a fast front and a slower back component.  In the jet of the Broad Line Radio Galaxy (BLRG) 3C~111 ($z=0.049$, $1\\,\\rm{mas}\\simeq 1\\,\\rm{pc}$), a very strong flaring event in early 1996 gave rise to the ejection of two jet features observed at 15 GHz with the Very Long Baseline Array (labeled as components E and F -- see Fig.~\\ref{fig:0} and \\citealt*{Kad08}, hereafter \\citetalias{Kad08}). Both component trajectories can be back-extrapolated to similar ejection epochs within 3 months (around 1996.10). However, they show different speeds and the time evolution of their brightness is different (see Fig.~\\ref{fig:0}): the inner component F is initially brighter (1996.82 and 1997.19) and fades out very rapidly (1997.66 and 1998.18), while the leading component E shows a slower decrease in flux density. After 1999 (see \\citetalias{Kad08}), F disappeared and E evolves, accelerating and generating trailing components in its wake \\citep{Agu01}. The differences betweeen E and F cannot be attributed to different Doppler factors of the components, as these are very similar ($D_E\\sim3.2$, $D_F\\sim3.1$, following \\citetalias{Kad08}). The possibility that component F represents a second injection after component E is highly improbable on the basis of its velocity and its brightness evolution, which cannot be linked to the propagation in the wake of the latter. In this letter, we investigate, in a qualitative way, the possibility that these components are the front and rear region of a single perturbation. A quantitative comparison between the simulations and the observations would require detailed knowledge about the nature of the flow and is out of the scope of this work. The numbers and error estimates used in this letter that are related to the observations are taken or derived from \\citetalias{Kad08}.  \\begin{figure}[!t] \\centering \\includegraphics[clip,angle=0,width=0.17\\textheight]{fig1a} \\includegraphics[clip,angle=0,width=0.18\\textheight]{fig1b} \\caption{Core distance and flux density evolution with time of components E and F in 3C~111, based on the results from \\citetalias{Kad08}.} \\label{fig:0} \\end{figure} ",
        "conclusions": "We are able to describe the evolution of the leading components of a major ejection in the BLRG 3C~111 as a perturbation of dense (overpressured) material followed by a dilute medium. The qualitative explanation given in the previous section supports the picture of the ejection of perturbations of dense material giving rise to radio components. Our model would remain valid if the amplitude or time-duration of the perturbation is changed. In addition, we postulate that these ejections could be followed by a decrease in the injection rate of bulk flow particles. This avoids the formation of a reverse shock, which would lead to a qualitatively different observational result. Observational support for the inclusion of this tenuous material in the simulations can be found in the prominent emission gap following behind the E/F complex in 3C~111 (see \\citetalias{Kad08}). New ejection of emitting material is detected on the time scale of more than 2 years, corresponding to a gap width of up to 2\\,mas in 1999 (cf. \\citetalias{Kad08}). The observations of microquasar jets show that, in general, major ejections are followed by a decrease in radio brightness \\citep{fb04}. If we interpret this as a decrease in the injection rate in the jet, while the system relaxes back to the initial steady state or generates the conditions for the injection of a new perturbation, it could be a process similar to that explained here. However, the strong decrease in radio emission observed in microquasar jets is not observed in AGN jets. This setup, in which we put in relation the processes taking place in the jet and the accretion disk, relies also on the results of multiwavelength observational campaigns by \\cite{Mar02} and \\cite{Mar06}. The latter work showed a relationship between dips in the X-ray emission from the accretion disk and the ejection of radio components in 3C~120, and possibly in 3C~111.  We can place an upper limit for the lifetime of the ejection event in its passage through the radio core (opaque and compact emitting region at the origin of the radio jet) using only observational data. We consider that: a) the perturbation started to cross the radio core in the observed jet in 1996.10 \\citepalias{Kad08}, b) the receding rarefaction ``eroding'' component F moves with the observed velocity $v_F=0.91\\,c$ (Fig.\\ref{fig:0}), and c) the last epoch at which component F is observed (1998.18) is taken as the time at which the receding rarefaction has completely eroded this component -- which is justified since the flux of component F is one order of magnitude smaller than that of component E in this epoch \\citepalias{Kad08}. With these assumptions we can calculate the time that the last material of the perturbation needs to catch up the receding rarefaction, and from that, we can estimate the time lapse of the crossing through the core. We use the velocity of component E as an upper limit for the velocity of the fluid. The result tells us that the crossing of gas through the core had to end as soon as $(1.98\\pm0.02)$~years before epoch 1998.18, this is, $(0.10\\pm0.02)$~years after 1996.10. This means $\\Delta t \\sim (36\\pm7)$~days between the passage of the first material and that of the last portion of gas in the perturbation. This lapse of time should be smaller if the gas catching up the rarefaction is slower than the velocity of the head. This represents a first order estimate, which only depends on the ratio between the velocity of the material in the perturbation and that of the rarefaction wave, and provides an upper limit for the duration of such an event. We could consider this lifetime of the ejection as an upper limit for its triggering event in the black-hole/inner-accretion-disc system, as we are not taking into account the collimation and acceleration processes, which should be kinematically important in the most compact scales. In the case of microquasar jets, the radio flares that follow the dips in X-rays --and are related to the ejection of components-- have been estimated to last between seconds and tenths of minutes \\citep{mr99}, i.e., $10^4-10^6$ times smaller time-scales than for 3C~111, whereas the ratio of the black-hole masses is of the order of $10^8$ \\citep{mr99,gr06}. If we take into account that our result represents an upper limit for the timescale of the ejection event, this suggests that the time-scale factor between quasars and microquasars may not come directly from the ratio of black hole masses \\citep[see also ][]{Mar06b}.  Strong radio flares associated with the ejection of radio-components should be carefully followed up, with a sufficiently dense time sampling, in 3C~111 and other sources for which we can achieve similar or better linear resolution with the VLBI technique. This would help in performing analyses like that presented here and to test our conclusions. The stretching of the size of any relativistic structure propagating through a jet must largely favour the detection of such double structures in the jets of nearby AGN. Databases provided by monitoring programmes such as the MOJAVE/2~cm~VLBA \\citep{K04,LH05} survey are very valuable in this context. Future work includes monitoring in radio/X-ray campaigns of different AGN sources and further numerical calculations including multi-dimensional RHD, RMHD and emission simulations."
    },
    "0808/0808.1899_arXiv.txt": {
        "abstract": "{The radio sky is poorly sampled for rapidly varying transients because of the narrow field-of-view of most imaging radio telescopes at cm and shorter wavelengths. The emergence of sensitive long wavelength observations with intrinsically larger fields-of-view are changing this situation, as partly illustrated by our ongoing meter-wavelength monitoring observations and archival studies of the Galactic Center. In this search, we discovered a transient, bursting, radio source in the direction of the Galactic Center, GCRT J1745$-$3009, with extremely unusual properties. Its flux and rapid variability imply a brightness temperature $>10^{12}$ K if it is at a distance $>70$ pc, implying that it is a coherent emitter. I will discuss the discovery of the source and the subsequent re-detections, as well as searches for counterparts at other wavelengths, and several proposed models. }  \\FullConference{Bursts, Pulses and Flickering: Wide-field monitoring of the dynamic radio sky June 12-15 2007\\\\ Kerastari, Tripolis, Greece}  \\begin{document} ",
        "introduction": "It is an interesting fact that relatively few transient radio sources are known, particularly when one considers only radio-selected sources. That is, most that are found are triggered by transient phenomena at other wavelengths such as supernova, GRBs, magnetar flares, and microquasar outbursts. In fact, a great deal of radio astronomy makes the tacit assumption that the sky is unchanging. Common examples include sky surveys and multi-configuration imaging.  Why is this the case? Certainly the answer is partly intrinsic.  Many bright radio sources come from extended regions such as supernova remnants and radio lobes that physically cannot vary rapidly. However, most radio observations have been performed and analyzed in a way that is highly insensitive to transient source detection. There are several reasons for this.  Many observations have moved up to high frequencies in search of high angular resolution and since the primary beam width is inversely proportional to frequency, this rapidly reduces the solid angle coverage of an observation. To make matters worse, most observers have a specific target in mind and only image a small field of view, ignoring most of the primary beam. Finally, limited observing time at the major interferometers means that most fields are revisited only rarely resulting in extremely sparse sampling for any transient source searches.  Pushing to low frequencies in search of a naturally larger primary beam requires sufficient angular resolution to mitigate confusion and provide good sensitivity. This requires long interferometer baselines (> 5 km), and has been historically very challenging because ionospheric and 3-D effects require computationally-intensive algorithms to enable sensitive wide-field imaging. Only recently has the computational power and software become available to enable ionospheric calibration techniques and faceted wide-field imaging algorithms to be used routinely \\cite{kle+07}. Nevertheless, there are many potential classes of transient radio sources \\cite{clm04} that could probe some of the most interesting questions in astrophysics.  This situation contrasts strongly with that in the X-ray band, where the full sky is routinely monitored at a range of timescales.  The most prominent example is the All-Sky Monitor (ASM) on the Rossi X-ray Timing Explorer (RXTE), which has surveyed the full sky many times per day for over a decade.  This has resulted in a great understanding of the population of bursting and transient sources on timescales from hours to years.  Many of these transient X-ray sources involve highly accelerated electrons, often in the form of jets, that also emit in the radio.  In addition there could be other populations of sources that produce radio transients without strong X-ray emission. It seems likely that the paucity of radio transients arises predominately from a lack of sensitive, wide-field observations, a situation that will dramatically change with the advent of many new large-etendue telescopes over the next few years. ",
        "conclusions": ""
    },
    "0808/0808.3778_arXiv.txt": {
        "abstract": "To make predictions for an eternally inflating ``multiverse,'' one must adopt a procedure for regulating its divergent spacetime volume.  Recently, a new test of such spacetime measures has emerged: normal observers --- who evolve in pocket universes cooling from hot big bang conditions --- must not be vastly outnumbered by ``Boltzmann brains'' --- freak observers that pop in and out of existence as a result of rare quantum fluctuations. If the Boltzmann brains prevail, then a randomly chosen observer would be overwhelmingly likely to be surrounded by an empty world, where all but vacuum energy has redshifted away, rather than the rich structure that we observe. Using the scale-factor cutoff measure, we calculate the ratio of Boltzmann brains to normal observers.  We find the ratio to be finite, and give an expression for it in terms of Boltzmann brain nucleation rates and vacuum decay rates.  We discuss the conditions that these rates must obey for the ratio to be acceptable, and we discuss estimates of the rates under a variety of assumptions. ",
        "introduction": "The simplest interpretation of the observed accelerating expansion of the universe is that it is driven by a constant vacuum energy density $\\rho_\\Lambda$, which is about three times greater than the present density of nonrelativistic matter. While ordinary matter becomes more dilute as the universe expands, the vacuum energy density remains the same, and in another ten billion years or so the universe will be completely dominated by vacuum energy. The subsequent evolution of the universe is accurately described as de Sitter space.  It was shown by Gibbons and Hawking~\\cite{GH} that an observer in de Sitter space would detect thermal radiation with a characteristic temperature $T_{\\rm dS}=H_\\Lambda/2\\pi$, where \\beq H_\\Lambda=\\sqrt{{8 \\pi \\over 3} G \\rho_\\Lambda} \\eeq is the de Sitter Hubble expansion rate. For the observed value of $\\rho_\\Lambda$, the de Sitter temperature is extremely low, $T_{\\rm dS}= 2.3\\times 10^{-30}$~K.  Nevertheless, complex structures will occasionally emerge from the vacuum as quantum fluctuations, at a small but nonzero rate per unit spacetime volume. An intelligent observer, like a human, could be one such structure. Or, short of a complete observer, a disembodied brain may fluctuate into existence, with a pattern of neuron firings creating a perception of being on Earth and, for example, observing the cosmic microwave background radiation.  Such freak observers are collectively referred to as ``Boltzmann brains''~\\cite{Rees1,Albrecht}. Of course, the nucleation rate $\\Gamma_{\\rm BB}$ of Boltzmann brains is extremely small, its magnitude depending on how one defines a Boltzmann brain.  The important point, however, is that $\\Gamma_{\\rm BB}$ is always nonzero.  De Sitter space is eternal to the future.  Thus, if the accelerating expansion of the universe is truly driven by the energy density of a stable vacuum state, then Boltzmann brains will eventually outnumber normal observers, no matter how small the value of $\\Gamma_{\\rm BB}$~\\citenosort{Susskind,Page1,Page2,Page06,BF06} might be.  To define the problem more precisely, we use the term ``normal observers'' to refer to those that evolve as a result of non-equilibrium processes that occur in the wake of the hot big bang.  If our universe is approaching a stable de Sitter spacetime, then the total number of normal observers that will ever exist in a fixed comoving volume of the universe is finite. On the other hand, the cumulative number of Boltzmann brains grows without bound over time, growing roughly as the volume, proportional to $e^{3H_\\Lambda t}$. When extracting the predictions of this theory, such an infinite preponderance of Boltzmann brains cannot be ignored.  For example, suppose that some normal observer, at some moment in her lifetime, tries to make a prediction about her next observation.  According to the theory there would be an infinite number of Boltzmann brains, distributed throughout the spacetime, that would happen to share exactly all her memories and thought processes at that moment. Since all her knowledge is shared with this set of Boltzmann brains, for all she knows she could equally likely be any member of the set.  The probability that she is a normal observer is then arbitrarily small, and all predictions would be based on the proposition that she is a Boltzmann brain. The theory would predict, therefore, that the next observations that she will make, if she survives to make any at all, will be totally incoherent, with no logical relationship to the world that she thought she knew. (While it is of course true that some Boltzmann brains might experience coherent observations, for example by living in a Boltzmann solar system, it is easy to show that Boltzmann brains with such dressing would be vastly outnumbered by Boltzmann brains without any coherent environment.) Thus, the continued orderliness of the world that we observe is distinctly at odds with the predictions of a Boltzmann-brain-dominated cosmology.\\footnote{Here we are taking a completely mechanistic view of the brain, treating it essentially as a highly sophisticated computer. Thus, the normal observer and the Boltzmann brains can be thought of as a set of logically equivalent computers running the same program with the same data, and hence they behave identically until they are affected by further input, which might be different.  Since the computer program cannot determine whether it is running inside the brain of one of the normal observers or one of the Boltzmann brains, any intelligent probabilistic prediction that the program makes about the next observation would be based on the assumption that it is equally likely to be running on any member of that set.}  This problem was recently addressed by Page~\\cite{Page06}, who concluded that the least unattractive way to produce more normal observers than Boltzmann brains is to require that our vacuum should be rather unstable.  More specifically, it should decay within a few Hubble times of vacuum energy domination; that is, in 20 billion years or so.  In the context of inflationary cosmology, however, this problem acquires a new twist. Inflation is generically eternal, with the physical volume of false-vacuum inflating regions increasing exponentially with time and ``pocket universes'' like ours constantly nucleating out of the false vacuum.  In an eternally inflating multiverse, the numbers of normal observers and Boltzmann brains produced over the course of eternal inflation are both infinite.  They can be meaningfully compared only after one adopts some prescription to regulate the infinities.  The problem of regulating the infinities in an eternally inflating multiverse is known as the measure problem~\\cite{MPreviews}, and has been under discussion for some time.  It is crucially important in discussing predictions for any kind of observation.  Most of the discussion, including the discussion in this paper, has been confined to the classical approximation.  While one might hope that someday there will be an answer to this question based on a fundamental principle~\\cite{principle}, most of the work on this subject has focussed on proposing plausible measures and exploring their properties.  Indeed, a number of measures have been proposed~\\cite{Linde:1986fd,LM,LLM,GBLL,AV94,AV95,Linde06,Linde07, GTV,pockets,GSPVW,ELM,diamond,censor,Vanchurin07,Winitzki08,LVW}, and some of them have already been disqualified, as they make predictions that conflict with observations.  In particular, if one uses the proper-time cutoff measure \\cite{Linde:1986fd,LM,LLM,GBLL,AV94}, one encounters the ``youngness paradox,'' predicting that humans should have evolved at a very early cosmic time, when the conditions for life were rather hostile~\\cite{Tegmark:2004qd}. The youngness problem, as well as the Boltzmann brain problem, can be avoided in the stationary measure \\cite{Linde07,LVW}, which is an improved version of the proper-time cutoff measure. However, the stationary measure, as well as the pocket-based measure, is afflicted with a runaway problem, suggesting that we should observe extreme values (either very small or very large) of the primordial density contrast $Q$~\\cite{FHW,QGV} and the gravitational constant $G$~\\cite{GS}, while these parameters appear to sit comfortably in the middle of their respective anthropic ranges~\\citenosort{anthQ,GS}. Some suggestions to get around this issue have been described in Refs.~\\cite{GarciaBellido:1994ci,QGV,Linde:2005yw,HWY}. In addition, the pocket-based measure seems to suffer from the Boltzmann brain problem. The comoving coordinate measure \\cite{Starobinsky:1986fx,Linde:1986fd} and the causal-patch measures~\\cite{diamond,censor} are free from these problems, but have an unattractive feature of depending sensitively on the initial state of the multiverse.  This does not seem to mix well with the attractor nature of eternal inflation: the asymptotic late-time evolution of an eternally inflating universe is independent of the starting point, so it seems appealing for the measure to maintain this property.  Since the scale-factor cutoff measure\\footnote{This measure is sometimes referred to as the volume-weighted scale-factor cutoff measure, but we will define it below in terms of the counting of events in spacetime, so the concept of weighting will not be relevant.  The term ``volume-weighted'' is relevant when a measure is described as a prescription for defining the probability distribution for the value of a field.  In Ref.~\\cite{Linde06}, this measure is called the ``pseudo-comoving volume-weighted measure.''} \\cite{LM,LLM,GBLL,AV95,Linde06,Clifton} has been shown to be free of all of the above issues~\\cite{DSGSV}, we consider it to be a promising candidate for the measure of the multiverse.  As we have indicated, the relative abundance of normal observers and Boltzmann brains depends on the choice of measure over the multiverse.  This means the predicted ratio of Boltzmann brains to normal observers can be used as yet another criterion to evaluate a prescription to regulate the diverging volume of the multiverse: regulators that predict normal observers are greatly outnumbered by Boltzmann brains should be ruled out.  This criterion has been studied in the context of several multiverse measures, including a causal patch measure~\\cite{BF06}, several measures associated with globally defined time coordinates~\\citenosort{Linde06,Page3,Linde07,BFY,LVW}, and the pocket-based measure~\\cite{AVBB}.  In this work, we apply this criterion to the scale-factor cutoff measure, extending the investigation that was initiated in Ref.~\\cite{Linde06}.  We show that the scale-factor cutoff measure gives a finite ratio of Boltzmann brains to normal observers; if certain assumptions about the landscape are valid, the ratio can be small.\\footnote{In a paper that appeared simultaneously with version 1 of this paper, Raphael Bousso, Ben Freivogel, and I-Sheng Yang independently analyzed the Boltzmann brain problem for the scale-factor cutoff measure~\\cite{BFY2}.}  The remainder of this paper is organized as follows.  In Section~\\ref{sec:sfcutoff}, we provide a brief description of the scale-factor cutoff and describe salient features of the multiverse under the lens of this measure.  In Section~\\ref{sec:BBabundance} we calculate the ratio of Boltzmann brains to normal observers in terms of multiverse volume fractions and transition rates.  The volume fractions are discussed in Section~\\ref{sec:minilandscapes}, in the context of toy landscapes, and the section ends with a general description of the conditions necessary to avoid Boltzmann brain domination. The rate of Boltzmann brain production and the rate of vacuum decay play central roles in our calculations, and these are estimated in Section~\\ref{sec:nucleationvsdecay}.  Concluding remarks are provided in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  If the observed accelerating expansion of the universe is driven by constant vacuum energy density and if our universe does not decay in the next 20 billion years or so, then it seems cosmology must explain why we are ``normal observers'' --- who evolve from non-equilibrium processes in the wake of the big bang --- as opposed to ``Boltzmann brains'' --- freak observers that arise as a result of rare quantum fluctuations~\\cite{Rees1,Albrecht,Susskind,Page06,BF06}.  Put in experimental terms, cosmology must explain why we observe structure formation in a residual cosmic microwave background, as opposed to the empty, vacuum-energy dominated environment in which almost all Boltzmann brains nucleate.  As vacuum-energy expansion is eternal to the future, the number of Boltzmann brains in an initially-finite comoving volume is infinite. However, if there exists a landscape of vacua, then rare transitions to other vacua populate a diverging number of universes in this comoving volume, creating an infinite number of normal observers.  To weigh the relative number of Boltzmann brains to normal observers requires a spacetime measure to regulate the infinities.  Recently, the scale-factor cutoff measure was shown to possess a number of desirable attributes, including avoiding the youngness paradox~\\cite{Tegmark:2004qd} and the $Q$ (and $G$) catastrophe~\\cite{FHW,QGV,GS}, while predicting the cosmological constant to be measured in a range including the observed value, and excluding values more than about a factor of ten larger and smaller than this~\\cite{DSGSV}.  The scale-factor cutoff does not itself select for a longer duration of slow-roll inflation, raising the possibility that a significant fraction of observers like us measure cosmic curvature significantly above the value expected from cosmic variance~\\cite{DSGSV2}.  In this paper, we have calculated the ratio of the total number of Boltzmann brains to the number of normal observers, using the scale-factor cutoff.  The general conditions under which Boltzmann brain domination is avoided were discussed in Subsection~\\ref{ssec:genconds}, where we described several alternative criteria that can be used to ensure safety from Boltzmann brains.  We also explored a set of assumptions that allow one to state conditions that are both necessary and sufficient to avoid Boltzmann brain domination. One relatively simple way to ensure safety from Boltzmann brains is to require two conditions: (1) in any vacuum, the Boltzmann brain nucleation rate must be less than the decay rate of that vacuum, and (2) for any anthropic vacuum $j$ with a decay rate $\\kappa_j \\approx q$, and for any non-anthropic vacuum $j$, one must construct a sequence of transitions from $j$ to an anthropic vacuum; if the sequence includes suppressed upward jumps, then the Boltzmann brain nucleation rate in vacuum $j$ must be less than the decay rate of vacuum $j$ times the product of all the suppressed branching ratios $B_{\\rm up}$ that appear in the sequence. The condition (2) might not be too difficult to satisfy, since it will generically involve only states with very low vacuum energy densities, which are likely to be nearly supersymmetric and therefore unlikely to support the complex structures needed for Boltzmann brains or normal observers. Condition (2) can also be satisfied if there is no unique dominant vacuum, but instead a dominant vacuum system that consists of a set of nearly degenerate states, some of which are anthropic, which undergo rapid transitions to each other, but only slow transitions to other states. The condition (1) is perhaps more difficult to satisfy. Although nearly-supersymmetric string vacua can in principle be long-lived~\\cite{Cvetic:1996vr,Ceresole:2001wi,Behrndt:2001mx,Louis:2006wq, KKLT,CDGKL}, with decay rates possibly much smaller than the Boltzmann brain nucleation rate, recent investigations suggest that other decay channels may evade this problem~\\cite{Kachru:2002gs,Frey:2003dm,FreivogelLippert}. However, the decay processes studied in ~\\cite{Cvetic:1996vr,Ceresole:2001wi,Behrndt:2001mx,Louis:2006wq, KKLT,CDGKL,Kachru:2002gs,Frey:2003dm,FreivogelLippert} do not describe some of the situations which are possible in the string theory landscape, and the strongest constraints on the decay rate obtained in \\cite{FreivogelLippert} are still insufficient to guarantee that the vacuum decay rate is always smaller than the fastest estimate of the Boltzmann brain production rate, Eq.~(\\ref{computer}).  One must emphasize that we are discussing a rapidly developing field of knowledge. Our estimates of the Boltzmann brain production rate are exponentially sensitive to our understanding of what exactly the Boltzmann brain is. Similarly, the estimates of the decay rate in the landscape became possible only five years ago, and this subject certainly is going to evolve. Therefore we will mention here two logical possibilities which may emerge as a result of the further investigation of these issues.  If further investigation will demonstrate that the Boltzmann brain production rate is always smaller than the vacuum decay rate in the landscape, the probability measure that we are investigating in this paper will be shown not to suffer from the Boltzmann brain problem. Conversely, if one believes that this measure is correct, the fastest Boltzmann brain production rate will give us a rather strong lower bound on the decay rate of the metastable vacua in the landscape. We expect that similar conclusions with respect to the Boltzmann brain problem should be valid for the causal-patch measures~\\cite{diamond,censor}.  On the other hand, if we do not find a sufficiently convincing theoretical reason to believe that the vacuum decay rate in all vacua in the landscape is always greater than the fastest Boltzmann brain production rate, this would motivate the consideration of other probability measures where the Boltzmann brain problem can be solved even if the probability of their production is not strongly suppressed.  In any case, our present understanding of the Boltzmann brain problem does not rule out the scale-factor cutoff measure, but the situation remains uncertain."
    },
    "0808/0808.2642_arXiv.txt": {
        "abstract": "\\def\\car{{\\hbox{}\\par\\parindent=0pt}} \\def\\Msun{M_{\\odot}} \\def\\Mstar{M_*} \\def\\re{r_e} \\def\\mre{\\Mstar/\\re} \\def\\omit#1{} \\def\\yrinv{yr$^{-1}$} \\def\\msunkpc{$\\Msun$ kpc$^{-2}$}  We present an analysis of galaxies in the CDF-South. We find a tight relation to $z=3$ between color and size at a given mass, with red galaxies being small, and blue galaxies being large. We show that the relation is driven by stellar surface density or inferred velocity dispersion: galaxies with high surface density are red and have low specific star formation rates, and galaxies with low surface density are blue and have high specific star formation rates. Surface density and inferred velocity dispersion are better correlated with specific star formation  rate and color than stellar mass. Hence stellar mass by itself is not a good predictor of the star formation history of galaxies. In general, galaxies at a given surface density have higher specific star formation rates at higher redshift. Specifically, galaxies with a surface density of $1-3\\ 10^9$ $M_\\odot$ $kpc^{-2}$ are ``red and dead'' at low redshift, approximately 50\\% are forming stars at $z=1$, and almost all are forming stars by $z=2$. This provides direct additional evidence for the late evolution of galaxies onto the red sequence. The sizes of galaxies at a given mass evolve like $1/(1+z)^{0.59 \\pm 0.10}$. Hence galaxies undergo significant upsizing in their history. The size evolution is fastest for the highest mass galaxies, and quiescent galaxies. The persistence of the structural relations from $z=0$ to $z=2.5$, and the upsizing of galaxies imply that a relation analogous to the Hubble sequence exists out to $z=2.5$, and possibly beyond. The star forming galaxies at $z\\ge 1.5$ are quite different from star forming galaxies at $z=0$, as they have likely very high gas fractions, and star formation time scales comparable to the orbital time. ",
        "introduction": "\\def\\car{{\\hbox{}\\par\\parindent=0pt}} \\def\\Msun{M_{\\odot}} \\def\\Mstar{M_*} \\def\\re{r_e} \\def\\mre{\\Mstar/\\re} \\def\\omit#1{} \\def\\yrinv{yr$^{-1}$} \\def\\msunkpc{$\\Msun$ kpc$^{-2}$}  \\def\\figprezero{\\figzero} \\def\\figpreone{\\figone} \\def\\figpretwo{\\figtwo} \\def\\figprethree{\\figthree} \\def\\figprefour{\\figfour} \\def\\figprefive{\\figfive} \\def\\figpresix{\\figsix} \\def\\figpreseven{\\figseven} \\def\\figpreeight{\\figeight} \\def\\figprenine{\\fignine} \\def\\figpreten{\\figten} \\def\\figpreeleven{\\figeleven} \\def\\figpretwelve{\\figtwelve} \\def\\figpreAone{\\figAone}  \\def\\figsubzero{} \\def\\figsubone{} \\def\\figsubtwo{} \\def\\figsubthree{} \\def\\figsubfour{} \\def\\figsubfive{} \\def\\figsubsix{} \\def\\figsubseven{} \\def\\figsubeight{} \\def\\figsubnine{} \\def\\figsubten{} \\def\\figsubeleven{} \\def\\figsubtwelve{} \\def\\figsubAone{}    \\def\\figone{ \\begin{figure*}[t] \\plotone{fig2.eps} \\caption{\\small Comparison of the star formation rates estimated from the UV + 24$\\mu$m MIPS fluxes, and the star formation rates estimated by SED fits. There is generally a good correspondence at redshifts below 2.5, with mild offsets smaller than 30\\%. For the remainder of the paper, we use the UV + 24$\\mu$m MIPS star formation rates. The typical formal error in the estimated star formation rates is shown in the lower right corner. The systematic uncertainties are a factor of 2 or more. \\label{figone}} \\end{figure*} }  \\gdef\\figurenum#1{} \\def\\figzero{ \\begin{figure}[t] \\epsscale{0.9} \\plotone{fig1.eps} \\epsscale{1.0} \\caption{\\small The redshift distribution of the final CDFS sample. The galaxies are  selected in the observed $K$ band, which lies redward of the rest-frame Balmer-4000\\AA\\ break for this sample. \\label{figone}} \\end{figure} }   \\def\\figtwo{ \\begin{figure*}[t] \\epsscale{1.2} \\plotone{fig3.eps} \\epsscale{1.0} \\caption{\\small Top row: the restframe $u-g$ color against stellar mass, for the low redshift sample (from the SDSS), and the redshifts 1, 2 and 3. A large spread in restframe $u-g$ colors exists for masses between $10^{10}$ and $10^{11}$, for all redshifts. The absence of red, low mass galaxies at $z\\ge 2$ is due to selection effects. The thick curves in the higher redshift panels delineate the  area where the samples are less than 75 \\% complete. The red galaxies to the upper left of the curve  are missing as they are too faint in the observed near-IR. Bottom row: The relation between  size  and  stellar mass for our redshift intervals. A large spread in sizes exists for masses between $10^{10}$ and $10^{11}$, for all redshifts. The thin dashed line indicates the mass-size relation from Shen et al. (2003) for early-type galaxies at $z=0$. It is clear from the figure that the scatter in color and size is very large for masses between 10$^{10}$ and 10$^{11}$ $\\Msun$. As in subsequent plots, the  typical formal error is shown in the lower right corner. \\label{figone}} \\end{figure*} }    \\def\\figthree{ \\begin{figure*}[t] \\epsscale{1.2} \\plotone{fig4.eps} \\epsscale{1.00} \\caption{\\small The correlation between color and size, at fixed stellar mass, for our sample. The top row shows the relation for the SDSS sample, the row below for galaxies with $0.5<z < 1.5$, the 3rd row for galaxies with $1.5< z < 2.5$, the bottom row for galaxies with $2.5 < z < 3.5$. The top rows show a strong correlation between color and size in the mass bins, with large galaxies being blue, and small galaxies being red. This demonstrates that the scatter in color and size at a given mass is driven by a third parameter. A correlation is still present at $z\\approx 2$, especially in the mass bin $2.5\\ 10^{10} \\Msun < M_* < 4\\ 10^{10}$. At high redshift ($z>2.5$) the sample is too small to establish whether or not a relation exists. \\label{figone}} \\end{figure*} }  \\def\\figfour{ \\begin{figure*}[t] \\figurenum{4a} \\epsscale{1.2} \\plotone{fig5.eps} \\epsscale{1.0} \\caption{\\small Left panels: the correlation between color and surface density, at redshifts 0, 1, 2 and 3. At all redshifts, a correlation is found, with high surface density galaxies being red, and low surface density galaxies being blue. This indicates that surface density is one of the main driving parameters for galaxy evolution. The color is more strongly correlated with surface density, than with mass; and this holds out to the highest redshift. Right panels: the correlation between color and velocity dispersion, at the same redshifts. Again, a good correlation is found, with high velocity dispersion galaxies being red, and low velocity dispersion galaxies being blue. The velocity dispersions have been estimated from $M_*/r_e$. \\label{figone}} \\end{figure*} }     \\def\\figfive{ \\begin{figure*}[t] \\figurenum{5a} \\epsscale{1.2} \\plotone{fig6.eps} \\epsscale{1.0} \\caption{\\small Left panels: the correlation between specific star formation rate and surface density, at redshifts 0, 1, 2, and 3. The SDSS specific star formation rate is based on an analysis of emission lines, whereas the high redshift points are based on the 24 $\\mu$m emission. The arrows denote upper limits. As can be seen, the highest surface density galaxies have low specific star formation rates out to redshift of 2.5, showing that specific star formation rate is causing to a large degree the relation between color and surface density. The limiting specific star formation rates for the redshift of 3 galaxies are too high to allow the detection of a correlation. Right panels: The correlation between velocity dispersion and specific star formation rate. Again, a good correlation persists to $z=2.5$, and above that the specific star formation rates are not well enough determined. \\label{figone}} \\end{figure*} }      \\def\\figsix{ \\begin{figure*}[t] \\figurenum{6} \\epsscale{1.2} \\plotone{fig7.eps} \\epsscale{1.0} \\caption{\\small The evolution of the specific star formation rate in mass bins. In all mass bins, the specific star formation rates increase with redshift, and  the spread remains large. The dashed line indicates the specific star formation rate equal to $t_{hubble}^{-1}$. \\label{figone}} \\end{figure*} }   \\def\\figseven{ \\begin{figure*}[t] \\figurenum{7} \\epsscale{1.2} \\plotone{fig8.eps} \\epsscale{1.0} \\caption{ \\small The evolution of the specific star formation rate in bins of surface density. The specific star formation rates increase in all bins. In the lowest surface density bin, the specific star formation rates are always high, and very few quiescent galaxies exist. We define quiescent galaxies as those having specific star formation rates below $0.3 \\times t_{hubble}^{-1}$. The dashed line indicates a specific star formation rate of $t_{hubble}^{-1}$. In the intermediate surface density bin with $ 10^9 < SD < 10^{9.5}$ the specific star formation rates are very low at $z<0.2$, increase rapidly to high values at $1=0.8-1.5$, to have only high values at $z \\ge 1.5$. The highest surface density bins always have the lowest specific star formation rates. \\label{figone}} \\end{figure*} }   \\def\\figeight{ \\begin{figure*}[t] \\figurenum{8} \\epsscale{1.2} \\plotone{fig9.eps} \\epsscale{1.0} \\caption{\\small The median specific star formation rate as a function of stellar surface density, at 3 redshift intervals: local galaxies from SDSS, and galaxies at $0.5<z<1.5$, and $1.5<z<2.5$. The median specific star formation rate is fairly constant at low surface densities, to decrease rapidly at higher surface densities. We have defined the threshold density, where the specific star formation rate is a factor of 3 lower than the specific star formation rate of the plateau. \\label{figone}} \\end{figure*} }  \\def\\fignine{ \\begin{figure*}[t] \\figurenum{9} \\epsscale{1.2} \\plotone{fig10.eps} \\epsscale{1.0} \\caption{\\small a) The evolution of surface density threshold, above which the specific star formation rate is low. At higher redshifts, the threshold has been measured in bins of width $\\delta z = 1$. The threshold evolves fast, close to $(1+z)^{1.5}$, as shown by the curve. b). The evolution of the specific star formation rate below the surface density threshold. The specific star formation rate evolves very rapidly like $(1+z)^{3.8}$, as shown by the curve. The dotted line curve the specific star formation rate equal to $t_{hubble}^{-1}$. \\label{figone}} \\end{figure*} }  \\def\\figten{ \\begin{figure*}[t] \\figurenum{10} \\epsscale{1.2} \\plotone{fig11.eps} \\epsscale{1.0} \\caption{\\small The evolution of half-light radius with redshift, in narrow mass bins. The sizes steadily decreases with  redshift, both for high mass and low mass galaxies. The curves show fits using the average coefficient $r_e \\propto (1+z)^{-0.59}$. The dashed curves show the fit using $r_e \\propto H(z)^{-2/3}$, which is predicted from simple formation models. The dotted curves show the fits of $\\re \\propto 1/(1+z)^\\alpha$ to the individual bins. The index $\\alpha$ varies from 0.14$\\pm$0.07 to 0.81$\\pm$ 0.1, from the low mass to the high mass bin. The average index is 0.59 $\\pm$ 0.10. The triangles are quiescent galaxies with low specific star formation rates (as defined in the text). They evolve faster, like $\\re \\propto 1/(1+z)^{1.22\\pm 0.15}$. Overall, the sizes evolve significantly with redshift, the strongest for the high mass galaxies, and the quiescent galaxies. \\label{figone}} \\end{figure*} }  \\def\\figeleven{ \\begin{figure*}[t] \\figurenum{11} \\epsscale{1.0} \\plotone{fig12.eps} \\epsscale{1.0} \\caption{\\small Left panel: The ratio of star formation time scale to dynamical time scale. Whereas at low redshift, the dynamical timescale is generally more than 30 times longer than the star formation time scale , it decreases to values of 3-10 at a redshift of 1.5 and above for the strongly star forming galaxies. For these galaxies, the gas has barely time to settle in a disk. Right panel: The evolution of the ratio of gas mass to stellar mass with redshift. The gas mass has been estimated from the star formation rate, and the kennicutt star formation law. Galaxies with significant star formation are indicated with a large symbol (specific star formation $>$ 1/3 threshold specific star formation rate). Whereas at low redshift the typical gas fraction is around 10\\%, it increases steadily with redshift, to reach values of 0.3-1 at $z=1.5$ and above. The most strongly star forming galaxies at $z=1.5$ and above have gas masses comparable to their stellar masses. \\label{figone}} \\end{figure*} }  \\def\\figtwelve{ \\begin{figure*}[t] \\figurenum{12} \\epsscale{1.0} \\plotone{fig13.eps} \\epsscale{1.0} \\caption{\\small Simulated selection effects for our $z=2$ sample. The left panel shows the selection on  observed $H_\\alpha$ flux. The $H_\\alpha$ flux has been estimated from the star formation rate, and extinction corrected, assuming the extinction from the SED fit. Galaxies with H$\\alpha$ flux $>$ $10^{42}$ ergs sec$^{-1}$  and $V_{ab} < 25.5$ are indicated  with full squares, the galaxies with H$\\alpha$ flux $>$ $10^{42}$ ergs sec$^{-1}$ and $V_{ab} > 25.5$ are indicated with open triangles, and the other galaxies are indicated with open squares. The V band limit is used as most often Near-IR spectroscopy is done on galaxies with optical redshifts measured first. The right panel shows the selection on  UV-slope and UV continuum flux (right), simular to those used for Ly-break and BM-BX galaxies (e.g., Steidel et al. 2006). Filled squares are galaxies with $V_{ab} < 25.5 $ and $B-V < 1.2$. In both panels, we see that the selected galaxies are prefentially large compared to the sample as a whole. It is obviously difficult to study galaxies at the full range of sizes, given the mass. We note that our full sample is also likely to be incomplete at low masses and small sizes. \\label{figone}} \\end{figure*} }  \\def\\figAone{ \\begin{figure*}[t] \\epsscale{1.2} \\plotone{fig14.eps} \\epsscale{1.0} \\caption{\\small Simulation  of the selection effects. Top row: galaxies at $0.5 < z < 1.5$ have been shifted to higher redshift by $\\delta z =  1$, while keeping their intrinsic properties the same. The open symbols show the galaxies which would still be selected using the selection criteria in this paper. It is clear that they are not biased in the mass radius relation or in the relation between specific star formation rate and surface density. Bottom row: the same, but now shifting galaxies at $1.5 < z < 2.5$ by $\\delta z = 1$. Again, the resulting relations are the same. If anything, the selected galaxies are slightly larger than the original full distribution, caused by the fact that small galaxies are red, and therefor dropped earlier from the sample. The median radius increases by 15\\%. The simulations shows that the evolutionary effects observed in this paper are not caused by selection effects. \\label{figone}} \\end{figure*} }    The dramatic progress in observational capabilities of the last decade have enabled studies of galaxy evolution and formation to epochs which were completely inaccessible only 15 years ago. Nevertheless, despite this dramatic progress, many open questions remain. Some of these are due to very fundamental problems related to studies of galaxy formation. The most important problem is that it has become clear that galaxy evolution is a complex process. Galaxy mergers change the mass function continuously. Furthermore, star formation does probably not occur with nice, smooth exponentially declining star formation rates - starbursts  can change the luminosities of galaxies on a short timescale. As a result, studies of galaxy evolution and formation are of a statistical nature: we cannot directly establish which galaxies at $z=1$ would evolve into what galaxies at $z=0$, and we therefore have to analyze statistically the sample as a whole.   Modern analyses therefore employ large samples, and attempt to use stellar masses whereever possible, as these are less likely to evolve rapidly (e.g., Brinchman \\& Ellis 2001, Kauffmann et al. 2003b, Dickinson et al. 2003, Rudnick et al. 2003, Fontana et al. 2005, Borch et al. 2006, Faber et al. 2007). However, as stressed by Faber et al. (2007), the mass function does not show a strong mode or other feature. As a result, we have to characterize the mass function by the characteristic mass at which it turns over. The determination of this characteristic mass and the determination of the general shape of the mass function is hard.   A different type of information may come from studies of the structure of galaxies. By studying the sizes, colors, velocity dispersions, etc, one may be able to derive tight correlations between galaxy properties. These tight correlations can then be used to study galaxy evolution. For example, by measuring velocity dispersions, sizes, and luminosities of early-type galaxies one can establish correlation between mass-to-light ratio and mass and size of a galaxy (e.g., Djorgovski et al. 1988, Faber et al. 1987). The evolution of these correlations with redshift (e.g., Franx 1993, van Dokkum and Franx 1996, van der Wel et al. 2005, Treu et al. 2005, van Dokkum \\& van der Marel 2007) provide the evolution of the mass-to-light ratio with redshift, which is essential for a proper interpretation of the evolution of the luminosity function. In addition this evolution puts strong constraints on the evolution of the early-type galaxies per se, and especially the time at which they formed their stars. However, we have to add that this interpretation is not entirely insensitive to the statistical evolution of the class of early-types (e.g., van Dokkum and Franx 2001). Nevertheless, detailed studies of galaxy properties can provide important additional constraints on galaxy evolution, even when relatively small samples are studied compared to the full statistical analyses.  In this paper we follow the second approach, where we study a fairly limited sample of $K$-band selected galaxies in the CDFS (1155 galaxies  from Wuyts et al. 2008), augmented by a local comparison sample based on the Sloan Digital Sky Survey (Strauss et al. 2002). % We use size measurements to explore the relations between color, masses, star formation rates and size. We show that a good correlation exists between the sizes, masses, colors and  specific star formation rates of the galaxies from $z=0$ to $z=3$, and we analyze this correlation in this paper. The paper is built up in the following way: section 2 presents the data used in the analysis, both the high redshift sample, and the low redshift sample. Section 3 presents the correlation between mass, size, color  at redshifts $z=0$, 1, 2 and 3. In addition, the relation between mass, size and specific star formation rate is presented. Section 4 presents the evolution of the specific star formation rate as a  function of redshift, and section 5 presents the evolution of size. The results are discussed in section 6, and in section 7 we discuss the potential biases that may occur in current samples of high redshift galaxies and analyses. The results are summarized in section 8. ",
        "conclusions": "We have shown that massive galaxies from $z=0$ out to $z=3.5$ have a strong correlation between size and color and size and specific star formation rate, at a given mass. Galaxies with high specific star formation rates are large, galaxies with low specific star formation rates are small. At increasing redshifts, the overall specific star formation rates go up.   In general, the specific star formation rates correlate better with surface density, and inferred velocity dispersion, than with  mass. This suggest that surface density, or inferred velocity dispersion, is the driving parameter. We find that, as expected, specific star formation rates  at a given surface density increase with redshift. We identified a threshold surface density at each redshift interval: below the threshold the specific star formation rates are high with little variation, above the threshold density galaxies have low specific star formation rates. As expected, the threshold increases with redshift: high specific star formation occurs at higher and higher surface density with increasing redshift. As a result, we find that many  galaxies which are on the red sequence at $z=0$ are star forming at $z=1$.  Furthermore, the sizes of galaxies at a given mass decrease with redshift steadily from $z=0$ to $z=3$.  This overall evolution shows that galaxies grow inside out. As this growth also occurs for 'red and dead galaxies', it suggests that these galaxies keep evolving - and are never just 'passively' evolving. In short, all galaxies show 'up-sizing', and the most massive galaxies show the strongest evidence for it. There is no evidence in this sample that very massive galaxies do not evolve between $z=1$ and $z=0$ (Scarlata et al. 2007), but obviously larger area studies can address this issue better. The very small, red galaxies at $z=2-3$ have to evolve into larger galaxies by $z=0$ through merging, accretion, and star formation. Their small sizes may simply be due to the fact that they are on the tail of the distribution at $z=2-3$, where all sizes are smaller. Two processes are likely  responsible for their small sizes: the halos were smaller and denser, and additionally their gas fractions were higher, and dissipation during merging was stronger, as suggested by Khochfar \\& Silk (2006). The observed evolution  is also consistent with merger simulations (e.g., Hopkins et al. 2008).  The similarity in structural relations for galaxies from $z=0$ to $z=2.5$, and the fact that galaxies grow inside out suggests that one form or another of the Hubble sequence persists to $z=2.5$, and possibly beyond. The older stars likely dominate in the centers, and younger stars are likely distributed over a larger radius - similar to bulges and disks in spiral galaxies. Although we don't have the resolution to establish this directly for the galaxies, the evolution of the mass-size relation with redshift also strongly supports such inside-out growth of galaxies. The exact dynamical state of high redshift galaxies still needs to be determined.  The multicomponent nature of galaxies suggest that modeling these galaxies is much harder, as the different populations have different spatial distributions, and therefore different extinction, star formation history, etc. Analysis of simulations suggests that this can lead to under-estimates of the masses through SED fits, and over-estimates of the sizes (e.g., Wuyts et al. 2007a).  The galaxies with very high star formation rates at $z \\ge 1.5$ likely have very high ratios  of gas mass to stellar mass (approximately 1 or above). Furthermore, they are large ($\\approx 3 kpc$), and therefore different from  ULIRGs in the nearby universe, which have typical sizes of the star formation regions of $<< 1kpc$. However, these high redshift galaxies may not be simple cold disks: their star formation time scale is only a few dynamical times, and therefore the gas had barely time to settle in discs, if accretion is causing the very high star formation. Alternatively, the star formation may have been overestimated, and this would allow the gas more time to settle.  Obviously, these results call for many follow-up observations. First and foremost, the mass estimates must be improved, hopefully through dynamical mass estimators, either through near-ir spectroscopy, or spectroscopy with ALMA. Second, higher resolution Near-IR imaging can determine the structure of nature of the high redshift galaxies better. High resolution imaging with ALMA will be able to establish the distribution of the star formation across the galaxies, and spectroscopy will allow the determination of the gas content and gas distribution.  Third, it will be important to extend this work to even higher redshifts, where samples selected in the rest-frame optical are very rare, and structural analyses absent.  Fourth, the environment of the high redshift  galaxies needs to be determined. This will allow a study of the the relation between structure, star formation history and environment  at high redshift. At low redshifts, environment plays an important role in setting the star formation rate (e.g., Kauffmann et al. 2004), and a full understanding requires an extension to high redshift. Finally, a determination of evolution of the surface density function can help to confirm the simple picture in which the centers of massive galaxies formed first.  The results show that stellar surface density, or inferred velocity dispersion, is one of the main driving parameters of galaxy evolution. Studies of the correlation of other galaxy properties with these parameters would be extremely valuable: metallicities, observed circular velocities, but also correlation length, and environment. Such studies require surveys of much larger areas, and extensive spectroscopy. The velocity dispersions used here are estimated from $M/\\re$, and it remains to be verified whether the two are well correlated at all redshifts, for all galaxies. If not, it is crucial to determine which is the driving parameter."
    },
    "0808/0808.0192_arXiv.txt": {
        "abstract": "The Einstein radius plays a central role in lens studies as it characterises the strength of gravitational lensing. In particular, the distribution of Einstein radii near the upper cutoff should probe the probability distribution of the largest mass concentrations in the universe. Adopting a triaxial halo model, we compute expected distributions of large Einstein radii. To assess the cosmic variance, we generate a number of Monte-Carlo realisations of all-sky catalogues of massive clusters.  We find that the expected largest Einstein radius in the universe is sensitive to parameters characterising the cosmological model, especially $\\sigma_8$: for a source redshift of unity, they are $42{}^{+9}_{-7}$, $35{}^{+8}_{-6}$, and $54{}^{+12}_{-7}$ arcseconds (errors denote $1\\sigma$ cosmic variance), assuming best-fit cosmological parameters of the Wilkinson Microwave Anisotropy Probe five-year (WMAP5), three-year (WMAP3) and one-year (WMAP1) data, respectively. These values are broadly consistent with current observations given their incompleteness.  The mass of the largest lens cluster can be as small as $\\sim10^{15}M_\\odot$. For the same source redshift, we expect in all-sky $\\sim 35$ (WMAP5), $\\sim 15$ (WMAP3), and $\\sim 150$ (WMAP1) clusters that have Einstein radii larger than $20''$. For a larger source redshift of 7, the largest Einstein radii grow approximately twice as large. Whilst the values of the largest Einstein radii are almost unaffected by the level of the primordial  non-Gaussianity currently of interest, the  measurement of the abundance of moderately large lens clusters should probe non-Gaussianity competitively with cosmic microwave background experiments, but only if other cosmological parameters are well-measured. These semi-analytic predictions are based on a rather simple representation of clusters, and hence calibrating them with $N$-body simulations will help to improve the accuracy. We also find that these ``superlens'' clusters constitute a highly biased population. For instance, a substantial fraction of these superlens clusters have major axes preferentially aligned with the line-of-sight. As a consequence, the projected mass distributions of the clusters are rounder by an ellipticity of $\\sim 0.2$ and have $\\sim 40\\%-60\\%$ larger concentrations compared with typical clusters with similar redshifts and masses. We argue that the large concentration measured in A1689 is consistent with our model prediction at the $1.2\\sigma$ level. A combined analysis of several clusters will be needed to see whether or not the observed concentrations conflict with predictions of the flat $\\Lambda$-dominated cold dark matter model. ",
        "introduction": "\\label{sec:intro}  In the standard Cold Dark Matter (CDM) model, which for these purposes we shall assume includes the presence of a cosmological constant and a flat spatial geometry, structure grows hierarchically from small objects that merge together to form larger objects (hereafter F$\\Lambda$CDM). Strong gravitational lensing by massive clusters of galaxies is one of the most important tests of this model in the sense that it probes the rarest high density peaks in the universe. For instance, the CDM model predicts wide-angle lensing events, on scales as large as several tens arcseconds, due to massive clusters \\citep[e.g.,][]{turner84,narayan84,narayan88,wambsganss95}. This has been broadly verified by the discovery of many lensed background galaxies \\citep[e.g.,][]{lefevre94,luppino99,gladders03,zaritsky03,sand05,hennawi08} or quasars \\citep{inada03,inada06}. Quantitative comparisons of expected lensing rates in the F$\\Lambda$CDM model and observed numbers of lenses should serve as an important test of our understanding of the universe.  A possible simple test of the F$\\Lambda$CDM model is the statistics of Einstein radii, particularly near the upper cutoff. The Einstein radius is a characteristic scale of strong lensing and is related mainly to the aperture mass it encloses. Therefore it is expected that the largest Einstein radii in the universe probe the structure and abundance of the most massive clusters. This enables a test of the F$\\Lambda$CDM model at the very tail of the halo distribution. An advantage of this test is the simple and straightforward determination of the Einstein radius in observations and its correspondence to identify large lenses.  Lensing properties of massive clusters have mainly been studied using ray-tracing in $N$-body simulations \\citep[e.g.,][]{wambsganss95, wambsganss04,bartelmann95,bartelmann98,meneghetti03a,meneghetti05, dalal04,ho05,li05,horesh05,hennawi07a,hennawi07b,hilbert07,hilbert08}. Whilst the numerical approach allows one to take account of the full complexity of lens potentials, it is often computationally expensive to perform large box-size simulations which retain enough particles in each halo for strong lensing studies. In particular reliable predictions for the rarest lensing events in the universe require many realisations of such high-resolution Hubble-size simulations in order to estimate the effect of the cosmic variance. This is impractical with current computational capabilities.  The complementary semi-analytic approaches often invoked simple spherically symmetric mass profiles, for calculational reasons \\citep[e.g.,][]{maoz97,hamana97,molikawa99,cooray99,wyithe01,takahashi01, kochanek01,keeton01,li03,lopes04,huterer04,kuhlen04,chen04,chen05,oguri06}. A more advanced calculation adopted an ellipsoid for projected cluster mass distributions \\citep{meneghetti03a,fedeli07a,fedeli07b,fedeli08}. However, because of the triaxial nature of F$\\Lambda$CDM haloes \\citep[e.g.,][]{jing02,allgood06,shaw06,hayashi07} the lensing properties of individual clusters vary drastically as a function of viewing angle \\citep[e.g.,][]{dalal04,hennawi07b}, resulting in the significant increase of average lensing efficiencies due to halo triaxiality. This indicates that any analytic models of cluster lensing should take proper account of triaxiality for reliable theoretical predictions, as is done in several papers \\citep{oguri03,oguri04,minor08}.  In this paper, we take a semi-analytic approach to predict the largest Einstein radius in all-sky survey, based on the F$\\Lambda$CDM model. We invoke an analytic mass function of dark haloes to generate a catalogue of massive clusters with the Monte-Carlo method. The shape of each halo is assumed to be triaxial, and the projection along random directions is considered. This Monte-Carlo approach allows us to  evaluate the range of the largest Einstein radii due to cosmic variance. We also characterise such ``superlens'' clusters, i.e., clusters which produce widest-angle lensing, to see how ``unusual'' are these clusters.  These issues are well illustrated by detailed observations of the largest Einstein radius known to data, which may conflict with the F$\\Lambda$CDM model. The lensing data of A1689, one of the best-studied clusters to date, suggest that the mass profile is apparently more centrally concentrated \\citep{broadhurst05a,broadhurst08a} than the F$\\Lambda$CDM prediction \\citep[e.g.,][]{neto07,duffy08,maccio08},  although the exact degree of concentration is somewhat controversial \\citep[e.g.,][]{halkola06,medezinski07,limousin07,comerford07,umetsu08}. It has been argued that a part of discrepancy can be explained by halo triaxiality \\citep{oguri05a,gavazzi05,hennawi07b,corless07,corless08} or the projection of the secondary mass peak along the line-of-sight \\citep{andersson04,lokas06,king07}, suggesting the importance of careful statistical studies with the selection effect taken into account. Indeed, it should be pointed out that a weak lensing analysis of stacked clusters of lesser mass does not exhibit the high concentration problem \\citep{johnston08,mandelbaum08}.  We believe that our predictions will be helpful for interpreting surveys of distant ($z\\ga 6$) galaxies near critical curves of massive clusters \\citep{ellis01,hu02,kneib04,richard06,richard08,stark07,willis08,bowens08}. The survey area of this technique is simply proportional to the square of the Einstein radius, thus clusters with very large Einstein radii are thought to be the best sites to conduct this search. Our calculations should provide a useful guidance to discover such giant lens clusters.  This paper is organised as follows. We describe our theoretical model in Section~\\ref{sec:mc}. Predictions for the largest Einstein radius and the abundance of large lens clusters are shown in Section~\\ref{sec:largest} and Section~\\ref{sec:num}, respectively.  Section~\\ref{sec:ng} includes the effect of primordial non-Gaussianities. We discuss the results in Section~\\ref{sec:dis}, and give our conclusion in Section~\\ref{sec:conc}. Throughout the paper, we consider three cosmological parameter sets obtained from the {\\it Wilkinson Microwave Anisotropy Probe (WMAP)}, mainly to show how sensitive our results are to cosmological parameters. These are the best-fit parameter sets from the WMAP one-year data \\citep[WMAP1;][]{spergel03}, ($\\Omega_M$, $\\Omega_b$, $\\Omega_\\Lambda$, $h$, $n_s$, $\\sigma_8$)$=$(0.270, 0.046, 0.730, 0.72, 0.99, 0.9), WMAP three-year data \\citep[WMAP3;][]{spergel07}, (0.238, 0.042, 0.762, 0.732, 0.958, 0.761), and WMAP five-year data \\citep[WMAP5;][]{dunkley08}, (0.258, 0.044, 0.742, 0.719, 0.963, 0.796). The most important difference between these models is the matter density and the normalisation of matter fluctuations. Indeed, it has been shown that the smaller values of $\\Omega_M$ and $\\sigma_8$ in WMAP3 resulted in much smaller number of cluster-scale lenses compared with WMAP1 \\citep[e.g.,][]{li06,li07}. Unless otherwise specified, we adopt the WMAP5 cosmology as our fiducial cosmological model. ",
        "conclusions": "\\label{sec:conc}  We have calculated the expected distributions of large Einstein radii in all-sky (40,000~${\\rm deg^2}$) using a triaxial halo model. Our approach to generate all-sky mock catalogue of massive clusters and the properties of individual clusters with the Monte-Carlo method allows us to evaluate the cosmic variance for such statistics, and at the same time, to study biases in the population of clusters having such large Einstein radii.  The largest Einstein radius in all-sky for source redshift $z_s=1$ was predicted to be $42{}^{+9}_{-7}$ arcseconds for WMAP5, $35{}^{+8}_{-6}$ arcseconds for WMAP3, and $54{}^{+12}_{-7}$ arcseconds for WMAP1, where errors are $1\\sigma$ cosmic variance. The sensitivity to $\\sigma_8$ suggests that this statistic is a good measure of it. The Einstein radii are approximately twice as large for larger source redshift, $z_s=7$. In some realisations the largest lens cluster is the most massive cluster in the universe; in others smaller than $10^{15}M_\\odot$. We have found that the population of these ``superlens'' clusters are significantly biased: their major axes are almost always aligned with the line-of-sight, and their projected 2D mass distributions appear rounder (by $\\Delta e\\sim0.2$) and more concentrated ($\\sim 40-60\\%$ larger values of concentration parameters). These biases are stronger than those found in more common lens clusters with smaller Einstein radii \\citep{hennawi07b}. In particular we have pointed out that the high concentration observed in A1689 is consistent with our theoretical expectation at the $1.2\\sigma$ level. Thus the combined analysis of several clusters will be essential to address the claimed high concentration problem. Finally, we have studied the effect of primordial non-Gaussianities, and concluded that the abundance of relatively large lens clusters can in principle constrain primordial non-Gaussianities at a level comparable to the current CMB experiments ($|f_{\\rm NL}|\\sim 100$), if other cosmological parameters are fixed.  It will be very important to compare our analytic predictions with ray-tracing in $N$-body simulations. In particular, the large cosmological Millennium Simulation \\citep{springel05} has sufficient resolution to resolve the centres of massive dark haloes for strong lensing studies \\citep{hilbert07,hilbert08}. Although its small box size ($500h^{-1}{\\rm Mpc}$) does not allow predictions for all-sky distributions of Einstein radii, we can use these simulations to validate and calibrate our semi-analytical model predictions. The comparison of our results with those from $N$-body simulations will be presented in a forthcoming paper.  The statistics of large Einstein radii provide an important opportunity to test the standard F$\\Lambda$CDM paradigm, as it probes both the high-mass end of the cluster mass function and central mass distributions of massive clusters. The measurement of Einstein radii is fairly robust, and future all-sky samples will soon be available to perform this study."
    },
    "0808/0808.2018_arXiv.txt": {
        "abstract": "We calculate the  shear viscosity $\\eta \\approx \\eta_\\mathrm{e\\mu}+\\eta_\\mathrm{n}$ in a neutron star core composed of nucleons, electrons and muons ($\\eta_\\mathrm{e\\mu}$ being the electron-muon viscosity, mediated by collisions of electrons and muons with charged particles, and $\\eta_\\mathrm{n}$ the neutron viscosity, mediated by neutron-neutron and neutron-proton collisions). Deriving $\\eta_\\mathrm{e\\mu}$, we take into account the Landau damping in collisions of electrons and muons with charged particles via the exchange of transverse plasmons. It lowers $\\eta_\\mathrm{e\\mu}$ and leads to the non-standard temperature behavior $\\eta_\\mathrm{e\\mu}\\propto T^{-5/3}$. The viscosity $\\eta_{\\rm n}$ is calculated taking into account that in-medium effects modify nucleon effective masses in dense matter. Both viscosities, $\\eta_\\mathrm{e\\mu}$ and $\\eta_\\mathrm{n}$, can be important, and both are calculated including the effects of proton superfluidity. They are presented in the form valid for any equation of state of nucleon dense matter. We analyze the density and temperature dependence of $\\eta$ for different equations of state in neutron star cores, and compare $\\eta$ with the bulk viscosity in the core and with the shear viscosity in the crust. ",
        "introduction": "\\label{introduc}  Neutron stars are very compact. Their typical masses are $\\sim 1.4\\, M_\\odot$ (where $M_\\odot$ is the mass of the Sun), while their radii are as small as $\\sim 10$~km. As a result, a neutron star core contains matter, whose density $\\rho$ reaches several $\\rho_0$ ($\\rho_0 \\approx 2.8 \\times 10^{14}$ g~cm$^{-3}$ being the density of the standard saturated nuclear matter). The core is composed of uniform neutron-rich nuclear matter and extends from $\\rho \\approx 0.5\\,\\rho_0$ to the stellar center (where $\\rho$ can be as high as $10 \\rho_0$). It attracts special attention because of its poorly known composition and equation of state (EOS); e.g., Ref.\\ \\cite{hpy07}. From outside,  the core is surrounded by a thin ($\\sim 1$ km thick) and light (a few per cent by mass) crust composed of atomic nuclei, strongly degenerate electrons and (after the neutron drip at $\\rho \\gtrsim 4 \\times 10^{11}$ g~cm$^{-3}$) free neutrons.  In this paper, we study the shear viscosity of neutron star cores. It is an important transport property which affects the relaxation of hydrodynamic motions, particularly, a possible differential rotation within the star and stellar oscillations \\cite{cl87}.  The shear viscosity can be important for damping gravitational wave driven instabilities (for instance, r-modes; e.g., \\cite{anderssonetal00} and references therein). Its knowledge is required to analyze the efficiency of such instabilities for generating gravitational waves.   For simplicity, we consider the cores composed of strongly degenerate neutrons (n), protons (p), electrons (e), and muons ($\\mu$) -- $\\mathrm{npe\\mu}$-matter, neglecting a possible appearance of hyperons and/or exotic forms of matter (pion or kaon condensates or quarks or their mixtures) as predicted by some EOSs at $\\rho \\gtrsim 2\\,\\rho_0$; see, e.g., Ref.\\ \\cite{hpy07}. The electrons and muons constitute almost ideal gases. The muons are absent in the outermost part of the core. They appear at densities exceeding a threshold value $\\rho_\\mu \\sim \\rho_0$ \\cite{st83} at which the electron chemical potential reaches the muon rest-mass energy ($\\mu_{\\rm e}=m_\\mu c^2 \\approx 207\\,m_{\\rm e}c^2$). The electrons are ultra-relativistic, while the muons are non-relativistic just after the threshold but become relativistic at higher $\\rho$. In contrast to electrons and muons, nucleons constitute a strongly interacting Fermi liquid where protons are essentially non-relativistic, while neutrons become mildly relativistic at $\\rho \\gtrsim 2\\,\\rho_0$. The neutrons and protons can be in superfluid state (e.g., Ref.\\ \\cite{ls01}).  The main contribution to the shear viscosity $\\eta$ in a neutron star core comes from electrons and muons (lightest and most mobile particles) and neutrons (most abundant particles), \\begin{equation} \\eta=\\eta_\\mathrm{e\\mu}+\\eta_\\mathrm{n}. \\label{eta_emun} \\end{equation} The viscosity $\\eta_\\mathrm{e\\mu}$ of electrons and muons is mainly limited by collisions of electrons and muons between themselves and with other charged particles (protons, in our case) via electromagnetic forces. In contrast, the neutron contribution $\\eta_\\mathrm{n}$ is limited by neutron-neutron and neutron-proton collisions mediated by strong interactions. As a result, $\\eta_\\mathrm{e\\mu}$ and $\\eta_\\mathrm{n}$ are nearly independent (belong to different -- electromagnetic and nuclear -- sectors) and can be calculated separately \\cite{fi79}.  In applications, one often employs the viscosity $\\eta_\\mathrm{e\\mu}$ calculated by Flowers and Itoh \\cite{fi79} for non-superfluid matter. Recently, Andersson et al.\\ \\cite{acg05} have estimated $\\eta_\\mathrm{e\\mu}$ for superfluid matter. However, these studies neglect an enhancement of collisions of relativistic charged particles due to the exchange of transverse plasmons. The significance of this effect was demonstrated by Heiselberg and Pethick \\cite{hp93} in their study of transport properties of ultra-relativistic quark matter. Recently we (Shternin and Yakovlev \\cite{sy07} -- hereafter  SY07) have reconsidered the electron-muon thermal conductivity $\\kappa_\\mathrm{e\\mu}$ taking into account the exchange of transverse plasmons. This effect can reduce $\\kappa_\\mathrm{e\\mu}$ by several orders of magnitude.  Here we reanalyze $\\eta_\\mathrm{e\\mu}$ in the same manner. We closely follow SY07 and omit technical details. In addition, we reconsider $\\eta_\\mathrm{n}$, which is a more difficult task involving nucleon-nucleon collisions. The viscosity $\\eta_\\mathrm{n}$  was calculated by Flowers and Itoh \\cite{fi79} for one EOS of non-superfluid matter assuming in-vacuum nucleon-nucleon scattering. These results were fitted by Cutler and Lindblom \\cite{cl87} by a simple analytical expression which is widely used; according to Ref.\\ \\cite{fi79}, $\\eta_\\mathrm{n}>\\eta_\\mathrm{e\\mu}$. Recently Benhar and Valli \\cite{bv07} have calculated $\\eta_\\mathrm{n}$ for pure neutron matter in a self-consistent manner using the same nucleon interaction potential to derive $\\eta_\\mathrm{n}$ and construct the EOS (also for one EOS). We calculate $\\eta_\\mathrm{n}$ in a more general way than Flowers and Itoh \\cite{fi79}. Our approach is similar to that used by Baiko, Haensel and Yakovlev \\cite{bhy01} (hereafter BHY01) for evaluating the thermal conductivity of neutrons. In addition, we employ recent developments \\cite{zhangetal07} in calculations of nucleon-nucleon scattering cross sections in nuclear matter. As in BHY01, we take into account superfluidity of protons. Again, we closely follow the derivation of BHY01 and omit the details.  After calculating $\\eta_\\mathrm{e\\mu}$ and $\\eta_\\mathrm{n}$, we analyze the shear viscosity in neutron star cores with different EOSs. ",
        "conclusions": "We have calculated the shear viscosity in a neutron star core as a sum of the electron-muon viscosity $\\eta_\\mathrm{e\\mu}$ and the neutron viscosity $\\eta_\\mathrm{n}$. Calculating the viscosity $\\eta_\\mathrm{e\\mu}$, which is mediated by collisions of charged particles, we have taken into account the exchange of transverse plasmons (that has not been done before). Our results include also the effects of proton superfluidity. They are universal, presented in the form of analytic fit expressions convenient for implementing into computer codes for any EOS of nucleon matter in neutron star cores.  Our main conclusions are: \\begin{enumerate} \\item The exchange of transverse plasmons strongly reduces $\\eta_\\mathrm{e\\mu}$ for all temperatures and densities of interest in a non-superfluid core. A low temperatures, we have $\\eta_\\mathrm{e\\mu}\\propto T^{-5/3}$.  \\item The  viscosity $\\eta_\\mathrm{e\\mu}$ generally dominates over $\\eta_\\mathrm{n}$, although $\\eta_\\mathrm{n}$ can exceed $\\eta_\\mathrm{e\\mu}$ at  $T\\lesssim 10^7$~K and $\\rho\\lesssim 4\\times 10^{14}$~g~cm$^{-3}$ (for $m_{\\rm n}^*\\approx m_{\\rm p}^*\\approx 0.8m_{\\rm n}$).  \\item The viscosity  $\\eta_{\\rm n}$ strongly depends on the nucleon effective masses. Typically, it is more than one order of magnitude lower, than that calculated in Ref.\\ \\cite{fi79} and parametrized in Ref.\\ \\cite{cl87}.  \\item Strong proton superfluidity significantly increases $\\eta_\\mathrm{e\\mu}$ and restores its Fermi-liquid temperature dependence, $\\eta_\\mathrm{e\\mu}\\propto T^{-2}$. In this regime, $\\eta_\\mathrm{e\\mu}$ exceeds $\\eta_\\mathrm{n}$.  \\item The shear viscosity $\\eta$ is comparable with the bulk viscosity $\\zeta$ at $T \\sim 10^8$~K (for a star vibrating at a frequency $\\omega \\sim 10^4$~s$^{-1}$) and dominates at lower $T$. Superfluidity increases the importance of $\\eta$ in comparison with $\\zeta$. \\end{enumerate}  Our results can be used in simulations of neutron star hydrodynamics, in particular, to analyze the damping of internal differential rotation, stellar oscillations, gravitational wave driving instabilities.  Our results can be improved further in many respects. It would be most important to study the shear viscosity problem in the presence of neutron and proton superfluidity in the frame of multifluid hydrodynamics as discussed in Sec.\\ \\ref{superflu}.  Nevertheless, even our restricted standard one-fluid formulation is incomplete. Our calculations of $\\eta_{\\rm n}$ can be improved by taking into account the medium effects on the matrix elements of nucleon-nucleon scattering. However, we rely on the results of Ref.\\ \\cite{zhangetal07} that these medium effects are  weaker than the effects of nucleon effective masses (which we include explicitly). An account for the medium effects on the matrix elements would complicate the expressions for $\\eta_{\\rm n}$ (making them non-universal).  We have also neglected the effects of strong magnetic field which can modify the shear viscosity. For not too high magnetic fields, $B \\lesssim 10^{13}$~G, which do not affect the plasma polarization functions (e.g., Ref.\\ \\cite{sh08b}), the generalization of the present results to the magnetic case is straightforward. For stronger fields, the polarization tensor becomes anisotropic and the viscosity problem is very complicated.  The present results are in line with our studies of kinetic properties of relativistic plasma taking into account the exchange of transverse plasmons. These effects were studied by Heiselberg and Pethick \\cite{hp93} for ultra-relativistic quark plasma. They should be included in all calculations of kinetic properties of relativistic plasmas, particularly in neutron stars.  For the neutron star crust, the effect was studied in \\cite{sy06} (thermal conductivity) and \\cite{sh08a} (shear viscosity). For neutron star cores, it was analyzed in \\cite{sy07} (thermal conductivity), \\cite{sh08b} (electrical conductivity), and \\cite{jgp05} (neutrino pair emission in electron-electron collisions)."
    },
    "0808/0808.2526_arXiv.txt": {
        "abstract": "In the standard  paradigm, satellite galaxies are believed  to  be  associated  with  the population  of  dark  matter subhalos.   The assumption  usually  made is  that the  relationship between satellite galaxies and  sub-halos is similar to that between central  galaxies  and  host  halos.   In this  paper,  we  use  the conditional  stellar  mass  functions  of {\\it  satellite  galaxies} obtained from a large galaxy group catalogue together with models of the  subhalo mass  functions  to explore  the  consequences of  such assumption in  connection to the stellar mass  function of satellite galaxies and the fraction and fate of stripped stars from satellites in galaxy groups and clusters  of different masses.  The majority of the  stripped stars  in massive  halos are  predicted to  end  up as intra-cluster stars, and the  predicted amounts of the intra-cluster component as a function of  the velocity dispersion of galaxy system match  well the observational  results obtained  by Gonzalez  et al. (2007).   The fraction of  the mass  in the  stripped stars  to that remain bound  in the central  and satellite galaxies is  the highest ($\\sim  40\\%$  of the  total  stellar  mass)  in halos  with  masses $M_h\\sim  10^{14}\\msunh$.   If  all   these  stars  end  up  in  the intra-cluster component (Max), or  a maximum amount of these stars is accreted into the central galaxy  (Min), then the maximum fraction of the total  stars in  the whole universe that is in the diffused  intra-cluster component is $\\sim 19\\%$,  and the minimum is $\\sim  5\\%$. In the case of `Max', the stellar mass  of the intra-cluster component in massive halos with $M_h \\sim  10^{15} \\msunh$ is  roughly 6 times as large  as that of the  central galaxy.  This  factor decreases to $\\sim 2$, $1$ and $0.1$ in halos with $M_h \\sim 10^{14}$, $10^{13}$, and  $10^{12}  \\msunh$, respectively.   The  total  amount of  stars stripped  from satellite galaxies  is insufficient  to build  up the central galaxies in halos  with masses $\\la 10^{12.5}\\msunh$, and so the  quenching of  star formation  must occur  in halos  with higher masses.   In  semi-analytical models  and  simulations  that do  not resolve the diffused component, caution must be exercised when using the  observed  stellar   mass/luminosity  function  of  galaxies  to constrain the star-formation, feedback  and merger processes in dark matter halos. ",
        "introduction": "\\label{sec_intro}  Recent years  have seen a dramatic  impetus to link  galaxies to their dark  matter  halos.   In  particular,  the  development  of  powerful statistical  tools  such  as  the  halo  model,  the  halo  occupation distribution  (HOD)  and the  conditional  luminosity function  (CLF), combined with the  availability of large redshift surveys  such as the two-Degree Field  Galaxy Redshift Survey (2dFGRS;  Colless \\etal 2001) and  the  Sloan Digital  Sky  Survey  (SDSS;  York \\etal  2000),  have resulted  in  reliable descriptions  of  how  galaxies with  different properties are distributed over  halos of different masses (e.g.  Jing \\etal  1998; Peacock  \\& Smith  2000; Berlind  \\& Weinberg  2002; Yang \\etal 2003; van  den Bosch \\etal 2003, 2007;  Zheng \\etal 2005; Tinker \\etal 2005; Cooray 2006; Brown \\etal 2008; Cacciato \\etal 2008).  An  important aspect  of this  galaxy-dark matter  connection  is that there  are two different  kinds of  galaxies: central  galaxies, which reside at rest at the center  of their dark matter halo, and satellite galaxies, which orbit  the halo associated with a  central galaxy.  It is  generally believed  that satellite  galaxies themselves  reside in dark matter subhalos.  Prior to being accreted into their current host halo, satellite  galaxies were central galaxies,  and their associated subhalos were host halos (throughout  this paper we use the term `host halo' to  refer to a dark matter  halo which is not  a subhalo).  This implies a direct link between the occupation statistics of dark matter halos, and their merger/accretion histories.  While orbiting the host  halo, subhalos and their associated satellite galaxies  are  subjected  to   dynamical  friction  which  causes  the substructure to loose  its momentum and to sink  towards the center of the  host halo. In  the mean  time, the  substructure is  subjected to tidal forces which  try to dissolve it.  In  particular, tidal heating and stripping cause  the system to loose mass, and  may even result in the complete disruption of a  subhalo and its satellite galaxy.  There are thus three possible fates for the stars in satellite galaxies: (i) they  remain bound  to a  surviving  satellite galaxy,  (ii) they  are accreted by  the central galaxy, or  (iii) they are  stripped from the satellite galaxy, giving rise to  a stellar halo, which are stars that orbit  the host halo  but that  are not  gravitationally bound  to any particular galaxy.  Throughout this paper  we will refer to  the stars that  belong  to  (ii)   and  (iii)  combined  as  the  `non-surviving population'.  Obviously, the  satellites that are disrupted contribute their entire  stellar content to  the central galaxy or  stellar halo, while   the   survived   satellite   galaxies  may   also   contribute significantly due to tidal stripping.  The idea that satellite galaxies are stripped and/or disrupted is a   standard  prediction   of  hierarchical   models   of  structure formation. Numerous studies have addressed how this phenomenon gives rise to stellar  halos in systems ranging from  spiral galaxies like our own  Milky Way (e.g.  Searle  \\& Zinn 1978;  Jonston \\etal 1996; 2001; Robertson \\etal 2005; Font \\etal 2006) to rich galaxy clusters (e.g. Gallagher \\& Ostriker  1972; Merritt 1983; Mihos 2004; Willman \\etal 2004; Lin \\& Mohr 2004; Conroy \\etal 2007; Purcell \\etal 2007; 2008;  Henriques  \\etal 2008).  There  is  also ample  observational support for the  existence of stellar halos formed  out of disrupted satellite galaxies.   In particular, in  recent years it  has become clear that the stellar halo of  the Milky Way reveals a large amount of substructure  in the form  of stellar streams (Helmi  \\etal 1999; Yanny \\etal 2003; Bell \\etal  2008), In some cases these streams can be  unambiguously associated with  their original  stellar structure (Ibata,  Gilmore \\&  Irwin 1994;  Odenkirchen \\etal  2002).  Similar streams have  also been  detected in our  neighbor galaxy  M31 (e.g. Ferguson  \\etal 2002).   Unfortunately, due  to their  extremely low surface brightnesses,  it is very difficult to  detect stellar halos in   more   distant  galaxies   (but   see   Zibetti  \\etal   2004). Consequently,  our  knowledge  of  stellar halos  around  individual galaxies  is  fairly  limited.   However,  groups  and  clusters  of galaxies also contain a significant stellar halo component, which is usually      referred      to      as     `intra-cluster      stars' (ICS)\\footnote{Throughout this paper we will use the terms `ICS' and `stellar  halo' without  distinction.}.  In  fact,  data indicates that the  fraction of  stars associated with  such an  ICS component increases with increasing halo mass (Gonzalez, Zabludoff \\& Zaritsky 2007): in  massive clusters the mass  of the stellar halo  can be as large  as ten  times the  stellar  mass of  the central  (brightest) cluster  galaxy (Gonzalez,  Zabludoff  \\& Zaritsky  2005; Seigar  et al. 2007).  With the use of high-resolution numerical simulations, it has recently become possible to accurately determine the properties (mass function, spatial  distribution,   orbits,  density  profiles,   spins)  of  the population of  dark matter  subhalos (e.g.  Gao  \\etal 2004;  De Lucia \\etal 2004; Tormen \\etal 2004;  van den Bosch \\etal 2005; Weller \\etal 2005; Diemand \\etal 2007; Giocoli \\etal 2008).  If subhalos are indeed associated with satellite galaxies, these properties should be related to  the   luminosity  function,  spatial   distribution,  orbits,  and structural properties of satellite galaxies.  Unfortunately, the exact link  between  satellite galaxies  and  dark  matter  subhalos is  not trivial.  Since the stellar component of a satellite is expected to be more  tightly   bound  than  its  surrounding  dark   matter  (due  to dissipation  during the formation  process), the  dark matter  is more easily stripped, thus causing the ratio between stellar mass and total mass to increase  with time.  Consequently, it is  to be expected that the occupation statistics of subhalos as a function of their (current) mass are different  from those of host halos.  However, since subhalos were  host halos  before  being  accreted, it  seems  likely that  the occupation statistics of subhalos as  a function of their mass {\\it at the time of accretion} are identical  to those of host halos at that time.  Indeed, it has been shown  that models based on this ansatz are extremely  successful  in  explaining  the correlation  functions  and luminosity functions of galaxies at different redshifts (e.g.  Vale \\& Ostriker  2004, 2006;  Conroy, Wechsler  \\& Kravtsov  2006).  However, since satellite galaxies  only make up a small  fraction ($\\lta 30\\%$) of the total galaxy population  (e.g.  van den Bosch \\etal 2007, 2008; Tinker  \\etal 2007;  Yang  \\etal 2008a),  this  consistency cannot  be considered  a sensitive test  of the  subhalo -  satellite connection. This is  also apparent from the  fact that models  that link satellite properties to  the {\\it  current} subhalo mass  can also fit  the data remarkably well (see e.g., Mandelbaum \\etal 2006; Kim \\etal 2008).  In  this  paper, we  use  the  conditional  stellar mass  function  of satellite galaxies and the relation between stellar mass and halo mass for  central  galaxies,  together  with  the  subhalo  mass  functions obtained from recent numerical  simulations, to determine the relation between satellite galaxies and dark matter subhalos. In particular, we investigate  what  fraction   of  satellite  galaxies  survives,  what fraction is  accreted into  the central galaxy,  and what  fraction is tidally  disrupted.  Our  approach is  empirical, because  the stellar mass functions and  the central stellar mass -  halo mass relation are obtained from a  large SDSS galaxy group catalogue.   The structure of this  paper is organized  as follows.  Section \\ref{sec_data}  gives a brief description  of the data and  the model used in  this paper.  In Section~\\ref{sec_disrupt}   we   present   our  predictions   of   the conditional  stellar  mass   functions  for  satellite  galaxies.   In Section~\\ref{sec_mtq} we discuss the  possible fates of the stars that are stripped  from the satellite galaxies.  Finally,  we summarize our results in Section~\\ref{sec_summary}.   Throughout this paper we adopt a $\\Lambda$CDM cosmology with  parameters that are consistent with the three-year  data   release  of  the  WMAP   mission  (hereafter  WMAP3 cosmology):  $\\Omega_{\\rm  m} =  0.238$,  $\\Omega_{\\Lambda} =  0.762$, $\\Omega_{\\rm  b} = 0.042$,  $n=0.951$, $h=H_0/(100  \\kmsmpc)=0.73$ and $\\sigma_8=0.75$ (Spergel \\etal 2007). ",
        "conclusions": "\\label{sec_summary}   Using the  conditional stellar  mass functions for  satellite galaxies obtained by Yang et al.  (2008b) and the sub-halo mass functions given by Giocoli  et al.  (2008),  we study the connection  between subhalos and  satellite  galaxies.  Assuming  that  at  the  time of  accretion satellite galaxies are associated  with subhalos according to the same stellar mass - halo mass relation as present-day central galaxies with halos, we predict the stellar  mass function of satellite galaxies and compare our model predictions  with observations. Our main results can be summarized as follows: \\begin{itemize} \\item Assuming  that the stellar  masses of satellite galaxies  do not evolve after they are accreted into the host halos, we find that the model over-predicts the population of satellite galaxies, especially in  low mass  halos.  One  solution,  albeit unlikely,  is that  the stellar mass fractions of haloes  at higher redshifts are lower than at  the present.   The other,  which  is considered  in the  present paper,  is that a  significant fraction  of satellite  galaxies have been stripped of their stars or even totally disrupted. \\item Assuming that the amount  of disruption of satellite galaxies is a function of the ratio between  the mass of the subhalo and that of the host halo, we find  that the surviving fraction can be described by $f(x = \\log(m_s/M_h))  = -0.041 -0.820x -0.422x^2 -0.078x^3$. The decrease  of $f$ with  increasing $m_s/M_h$  is consistent  with the idea that dynamical friction  brings subhalos and satellite galaxies to the  inner part of the  host halo where  stripping and disruption are more efficient. \\item The  majority of the  stars stripped from satellites  in massive halos  are predicted  to  end  up as  intra-cluster  stars, and  the predicted amounts  of the intra-cluster  component as a  function of the   velocity  dispersion   of   galaxy  system   match  well   the observational  results obtained  by  Gonzalez et  al.  (2007).   The fraction of the mass in the  stripped stars to that in the surviving central and satellite galaxies is predicted to be the highest ($\\sim 40\\%$ of the total) in  halos with masses $\\sim 10^{14} \\msunh$.  If all  these stripped  stars  end up  in  the intra-cluster  component (Max),  or maximum  of them  are  accreted into  the central  galaxy (Min), then we can predict that  a maximum $\\sim 19\\%$ and a minimum $\\sim 5\\%$ of the total stars  in the whole universe are in terms of the diffused intra-cluster component. \\item  Stars stripped from  satellite galaxies  are not  sufficient to build   up  the  central   galaxies  in   halos  with   masses  $\\la 10^{12.5}\\msunh$, and  so star formation  should not be  quenched in halos with masses up to at least $\\sim 10^{12.4}\\msunh$. \\end{itemize}  It should be pointed out once  again that our results are based on the assumption  that subhalos  have  the  same stellar  mass  - halo  mass relation as the  host halos at redshift $z=0$. If  the total amount of stars that can  form in a halo at higher redshift  is larger than that in a halo of the same  mass at lower redshift (e.g., Cooray 2005), the predicted disrupted  fraction of  satellite galaxies would  be larger, and so would be the predicted amount of ICS. On the other hand, if the total amount  of stars that can  form in high-$z$ halos  is lower than that in  its $z=0$ counterparts,  the predicted ICS fraction  would be lower.  It  is interesting that  the observed ICS fraction  is matched well with the assumption that the  total amount of stars that can form in a halo of a given mass is independent of redshift.  This assumption is also  consistent with the results  obtained by Wang  et al.  (2006) based  on a  semi-analytical model  of structure  formation.  However, current  semi-analytical   models  of  galaxy   formation  ignore  the existence  of the  ICS component,  so that  the stars  that  formed in subhalos  are assumed  either to  remain in  satellite galaxies  or to merge into central galaxies.  Consequently, such models either predict too  high  a  mass  for  the  central galaxies  in  massive  halos  or over-predict the number of satellite galaxies in group-sized halos.  As  described above, the  mass fraction  of stars  in the  diffuse ICS component is much  larger than that in the  central galaxies for halos with masses $\\ga 10^{14}\\msunh$, and  even in the whole universe, this stellar  mass  component is  a  significant  fraction  of $\\sim  19\\%$ (upper-limit)  or $\\sim  5\\%$  (lower-limit). The  implication of  the existence of such a stellar  component for galaxy formation has yet to be explored. Indeed, if the  total amount of stars in groups, clusters and  in the  universe  is larger  than  that implied  by the  observed stellar mass  function of galaxies, and if  semi-analytical models and numerical simulations do not resolve  the ICS component, then it would be incorrect  to use the observed stellar  mass/luminosity function of galaxies  to   constrain  star  formation   efficiency  and  feedback. Furthermore, the fraction  of the ICS component in  halos of different masses may  also convey important  information about the  evolution of galaxies in different environments.  Clearly, more theoretical work is required in order to make full  use of the information provided by the ICS component."
    },
    "0808/0808.0009_arXiv.txt": {
        "abstract": "We develop a formalism to study non-Gaussianity in both curvature and isocurvature perturbations. It is shown that non-Gaussianity in the isocurvature perturbation between dark matter and photons leaves distinct signatures in the CMB temperature fluctuations, which may be confirmed in future experiments, or possibly, even in the currently available observational data.  As an explicit example, we consider the QCD axion and show that it can actually induce sizable non-Gaussianity for the inflationary scale, $H_{\\rm inf} = O(10^9 - 10^{11})$\\,GeV. ",
        "introduction": "\\label{intro} The accumulating observational data, especially the WMAP observation of the cosmic microwave background (CMB)~\\cite{Komatsu:2008hk}, provided significant support for the inflationary paradigm.  The results of these measurements are consistent with nearly scale-invariant, adiabatic and Gaussian primordial density perturbations, known as the standard lore in the simple class of inflation models.  A possible detection of the deviation from the above properties will enable us to further constrain inflation models.  The scalar spectral index of the power spectrum is constrained as $n_s = 0.963^{+0.14}_{-0.15}$ at 68\\% C.L.~\\cite{Komatsu:2008hk}, which already excludes some inflation models.  No significant isocurvature component has been detected so far, and the current constraint on the ratio of the amplitudes of the isocurvature and curvature perturbations reads, $|S/\\zeta| \\lesssim 0.3$~\\cite{Komatsu:2008hk,Bean:2006qz}.  Recently, Yadav and Wandelt claimed an evidence of the significant non-Gaussianity in the CMB anisotropy data. Using the non-linearlity parameter $f_{\\rm NL}$ to be defined in the next section, their result is written as $47 < f_{\\rm NL} < 127$ at 95\\% C.L.~\\cite{Yadav:2007yy}. On the other hand, the latest WMAP five-year result is consistent with the vanishing non-Gaussianity: $-9 < f_{\\rm NL} < 111$ at 95\\% C.L., including $f_{\\rm NL}=0$~\\footnote{ Here we have quoted the value of $f_{\\rm NL}^{\\rm local}$ since we are interested in non-Gaussianity of the local type in this paper. }. Interestingly, however, the likelihood distribution of the WMAP result is biased toward positive values of $f_{\\rm NL}$.  Also there are some other studies searching for the non-Gaussianity~\\cite{Slosar:2008hx}, and it is not settled yet whether the non-Gaussianity exists.  At the present stage, therefore, it is fair to say that the observations are consistent with the nearly scale invariant and pure adiabatic perturbations with Gaussian statistics, while there is a hint of non-Gaussianity at the two sigma level.  The standard lore on inflation is based on a simple but crude assumption that it is only the inflaton that acquires sizable quantum fluctuations during inflation.  Its apparent success, however, does not necessarily mean that such a non-trivial condition is commonly met in the landscape of the inflation theory. In fact, there are many flat directions in a supersymmetric (SUSY) theory and the string theory. If some of them are light during inflation, they acquire quantum fluctuations, which may result in slight deviation from the standard lore.  One promising candidate is provided by the theory with a Peccei-Quinn (PQ) symmetry, which is introduced in order to solve the strong CP problem in the quantum chromodynamics (QCD)~\\cite{Peccei:1977hh,Kim:1986ax}.  There appears a pseudo-Nambu-Goldstone boson called axion associated with the spontaneous breakdown of the PQ symmetry.  The axion is a light scalar field and contributes to the cold dark matter (CDM) of the universe~\\cite{Preskill:1982cy}. In particular, the axion can have a large isocurvature perturbation~\\cite{Seckel:1985tj,Linde:1991km}.  As for non-Gaussianity, it is known that the slow-roll inflation generally predicts a negligible amount of non-Gaussianity, $f_{\\rm NL}= O(\\epsilon, \\eta)$~\\cite{Acquaviva:2002ud,Maldacena:2002vr,Seery:2005wm,Yokoyama:2007uu}.  Here $\\epsilon$ and $\\eta$ are the slow-roll parameters, which must be smaller than unity for the slow-roll inflation to last long enough. In the curvaton~\\cite{Mollerach:1989hu,Linde:1996gt,Lyth:2001nq} and/or ungaussiton~\\cite{Suyama:2008nt} scenarios, however, there are light scalars in addition to the inflaton, which can generate sizable non-Gaussianity~\\cite{Lyth:2002my,Bartolo:2003jx,Lyth:2005fi,Lyth:2006gd, Malik:2006pm,Ichikawa:2008iq,Beltran:2008ei, Suyama:2008nt}. As we will see, the axion can also induce sizable non-Gaussianity.   In this paper we point out that, if an isocurvature component possesses some amount of non-Gaussianity, it is transferred to the non-Gaussianity of the curvature perturbation, resulting in a possibly large value of $f_{\\rm NL}$.  If there are no other light scalar fields than the inflaton, the resultant density perturbations are necessarily adiabatic and almost Gaussian. As mentioned before, this may not be the case in the presence of many flat directions.  Suppose that there is a light scalar that acquires quantum fluctuations during inflation. Then its fluctuations produce isocurvature component.  If the scalar decays into radiation, the isocurvature perturbation is converted into the adiabatic one. This is exactly what occurs in the curvaton and/or ungaussiton scenarios.  However, as far as non-Gaussianity is concerned, the light scalar having large fluctuations needs not decay. Even if such a light field is not responsible for the total curvature perturbation, it can still provide a source of large non-Gaussianity. This interesting possibility was noted in Refs.~\\cite{Linde:1996gt,Bartolo:2001cw,Boubekeur:2005fj}. In this paper we have systematically studied the non-Gaussainity from the isocurvature perturbations and how it exhibits itself in the CMB anisotropy. Interestingly enough, we have found that the resultant non-Gaussianity induced by the isocurvature perturbations has distinctive signatures in the CMB, which should be distinguished from that in the curvaton and ungaussiton scenarios.    This paper is organized as follows.  In Sec.~\\ref{formalism}, a general formalism to study non-Gaussianity including isocurvature perturbations is presented.  In Sec.~\\ref{sec:temp}, we compute the bispectrum of the temperature fluctuations arising from non-Gaussianity in the isocurvature perturbations.  In Sec.~\\ref{axion} the formalism is applied to the case of the axion and it is shown that the axion can induce large non-Gaussianity while leaving a certain amount of the CDM isocurvature perturbation. Sec.~\\ref{conclusion} is devoted to discussion and conclusions. ",
        "conclusions": "\\label{conclusion} In this paper we have investigated a possibility that large non-Gaussianity is generated by isocurvature fluctuations. One interesting feature of this scenario is that the bispectrum and the power spectrum of the CMB temperature fluctuations exhibit characteristic scale dependence. In particular, the effective non-linearity parameter $f_{\\rm NL}^{\\Delta T}$ is significantly enhanced at large scales,  compared to the adiabatic case. Furthermore, our results indicate that large non-Gaussianity may be accompanied with an observable fraction of the isocurvature perturbation. If future observations confirm both large non-Gaussianity and a certain amount of isocurvature fluctuation component, our scenario will become very attractive. As a concrete example, we have shown that the axion can naturally induce such isocurvature perturbation in the CDM sector, leading to large non-Gaussianity. If the axion is indeed responsible for the large non-Gaussianity hinted by the current observations, the inflationary scale should be in the range of $O(10^9 - 10^{11})$\\,GeV. This opens up an interesting possibility that the axion can be probed through its non-Gaussianity contribution to the CMB temperature fluctuation, even if the energy density of the axion today is negligible compared to the total dark matter abundance~\\footnote{ There may be an anthropic reason that forces the axion abundance to take such value~\\cite{NT}. }.  Although we have restricted ourselves to the CDM isocurvature perturbation in this paper, the baryonic isocurvature perturbation can also generate large non-Gaussianity in a similar fashion \\cite{Kawasaki:2008jy}. There is indeed a scenario using the flat direction with the baryon number to generate baryonic isocurvature perturbations~\\cite{EM,KT,Kasuya:2008xp}.  We have assumed that the CDM isocurvature perturbation comes from the fluctuations in the CDM sector. However, this may not be the case, and it can arise from the radiation. That is to say, the CDM isocurvature perturbation is given by  $S=3(\\zeta_{\\rm CDM}-\\zeta_r)=-3\\zeta_r$ in some cases, including the curvaton/ungaussiton scenario.  In this case the sign of the non-linearity parameter $f_{S}$ becomes negative, which is expected to leave distinct features on the bispectrum of the temperature fluctuation. One may be able to use the features to probe into the origin of the isocurvature perturbation \\cite{Kawasaki:2008pa}.  In general, the isocurvature perturbation mainly affects the large scale temperature anisotropy, while its effect on the density perturbation is weaker than the adiabatic one. Thus if the non-Gaussianity is truly sourced by the isocurvature perturbation, it can be seen only in the CMB observations, and we should have null detection of non-Gaussianity from the analyses using the matter power spectra. Hopefully, future cosmological observations will enable us to establish or refute the existence of large non-Gaussianity. If it is there, we may be able to distinguish the origin, adiabatic or isocurvature. Undoubtedly, once detected, it will provide us with  useful information on the early universe and the high energy physics."
    },
    "0808/0808.0186_arXiv.txt": {
        "abstract": "Within an isospin- and momentum-dependent hadronic transport model it is shown that the recent FOPI data on the $\\pi^-/\\pi^+$ ratio in central heavy-ion collisions at SIS/GSI energies (Willy Reisdorf {\\it et al.}, NPA {\\bf 781}, 459 (2007)) provide circumstantial evidence suggesting a rather soft nuclear symmetry energy \\esym at $\\rho\\geq 2\\rho_0$ compared to the Akmal-Pandharipande-Ravenhall prediction. Some astrophysical implications and the need for further experimental confirmations are discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.1180_arXiv.txt": {
        "abstract": "We present the first results of a targeted survey carried out with the 305~m Arecibo telescope to detect \\hi-line emission from galaxies at redshift $z>0.16$. The targets, selected from the Sloan Digital Sky Survey database, are non-interacting disk galaxies in relatively isolated fields. We present here the \\hi\\ spectra and derived \\hi\\ parameters for ten objects detected in this pilot program. All are massive disk galaxies in the redshift interval 0.17$-$0.25 (i.e. 2$-$3 Gyr look-back time), with \\hi\\ masses \\Mhi~$=3-8$\\x 10$^{10}$ \\Msun\\ and high gas mass fractions (\\hi - to - stellar mass ratios \\about 10$-$30\\%). Our results demonstrate the efficacy of exploiting Arecibo's large collecting area to measure the \\hi\\ mass and rotational velocity of galaxies above redshift $z=0.2$. In particular, this sample includes the highest redshift detections of \\hi\\ emission from individual galaxies made to date. Extension of this pilot program will allow us to study the \\hi\\ properties of {\\em field} galaxies at cosmological distances, thus complementing ongoing radio synthesis observations of cluster samples at $z$\\about 0.2. ",
        "introduction": "\\label{s_intr} Because cold atomic hydrogen represents the reservoir for future star formation and the physics of the 21 cm transition is relatively simple, detection of the \\hi\\ line emission has the potential to contribute a key quantitative ingredient to models of how galaxies acquire and retain gas and convert it into stars. Of critical importance is the tracking of the gas content over the redshift interval between $z \\sim 1$ and 0, i.e. over the period when the observed cosmic star formation rate density has declined by an order of magnitude \\citep[e.g.,][and references therein]{lilly96,madau96,bell05}. However, our knowledge of the \\hi\\ 21~cm line emission properties of galaxies is currently limited to the very local universe, and only recently have technical improvements in radio frequency receivers and back-ends made attempts to explore the \\hi\\ emission from galaxies at $z> 0.1$ in any practical sense.  The first detection of \\hi\\ emission from a galaxy at cosmological distance ($z=0.1766$) was made by \\citet{zvv01} with the Westerbork Synthesis Radio Telescope (WSRT). More recently, \\cite{ver07} successfully carried out a pilot program with the WSRT to image $z\\sim 0.2$ galaxy clusters at 21 cm. They targeted the Abell 963 ($z=0.206$) and Abell 2192 ($z=0.188$) clusters for 20\\x 12 hours and 15\\x 12 hours respectively, and detected \\hi\\ emission from 42 galaxies. The Giant Metrewave Radio Telescope was used by \\citet{lah07} to constrain the \\hi\\ content of star forming galaxies at $z=0.24$. They obtained an average \\hi\\ spectrum by co-adding the non-detections of 121 galaxies with known optical redshifts. However, their detection is at best marginal.  Historically, Arecibo's huge collecting area has offered a critical advantage to \\hi\\ line studies of extragalactic objects, because sensitivity is of paramount importance to the detection of weak \\hi\\ signals. Recent improvements to the Arecibo instrumentation, both front-end and back-end, are allowing the practical exploration of the frequency range corresponding to \\hi\\ at $z>0.1$. Most notably, the installation of the new, single-pixel L-band receiver (``L-wide'') in February 2003 has given access to frequencies down to 1120 MHz (corresponding to redshifted \\hi\\ emission at $z=0.27$) with improved system noise and performance. Naturally, single-dish \\hi\\ observations of intermediate redshift galaxies are challenged by a lack of angular resolution. As the density of galaxies per unit solid angle increases with redshift, so also does the probability of finding multiple sources within the telescope beam. However, beam confusion can be broken kinematically, if high-quality optical redshifts are available for the sources within the beam. The availability of imaging and spectroscopic data from the Sloan Digital Sky Survey \\citep[SDSS;][]{yor00} has enabled the identification of sources suitable for long-integration \\hi\\ line observations with the Arecibo telescope. In this Letter, we report the first results of a pilot survey carried out at Arecibo to detect \\hi\\ emission from disk galaxies at $z\\sim 0.2$.  All the distance-dependent quantities in this work are calculated assuming $H_0 = 71$ \\kmsm, $\\Omega_{\\rm M}=0.27$ and $\\Omega_{\\rm vac}=0.73$.  \\vspace{24pt} ",
        "conclusions": "In this Letter, we have shown that carefully-planned Arecibo observations can deliver high-quality global \\hi\\ profiles for galaxies at redshifts up to $z=0.25$ with integration times of a few hours per object. The detected objects are massive disk galaxies with high gas mass fractions (\\about 10$-$30\\%), and are apparently rare in the local universe. At present, the small size of our sample, selection effects and the lack of adequate comparison samples at $z=0$ prevent us from drawing any firm conclusions about the evolution of the \\hi\\ content of galaxies between redshifts 0.25 and 0. Larger and well-defined samples of galaxies at $z \\sim 0.2$ and higher are needed in order to make progress in this area. Equally importantly, fair comparison samples at $z=0$ are {\\em also} needed. The latter might be provided by two large surveys that are ongoing at Arecibo. The first one is the Arecibo Legacy Fast ALFA (ALFALFA) survey \\citep{g05,g07}, a blind \\hi\\ survey of the extragalactic sky visible from Arecibo, which will make a complete census of \\hi-rich galaxies with $z \\leq 0.060$ and has already detected a significant population of objects with \\hi\\ masses $> 10^{10}$ \\Msun\\ \\citep{h08}. The second one is the GALEX Arecibo SDSS Survey \\citep[GASS;][]{gass}, a recently started program designed to measure the \\hi\\ content of \\about 1000 massive galaxies, randomly selected from the intersection of the SDSS spectroscopic survey, GALEX and ALFALFA footprints based on their redshift ($0.025< z < 0.05$) and stellar mass ($10.0 < {\\rm Log}~[M_{\\rm star}/{\\rm M}_{\\odot}] < 11.5$) only. ALFALFA and GASS will deliver {\\em unbiased} samples of local \\hi -rich and massive galaxies, respectively. In combination with the future extension of our pilot intermediate redshift field galaxy program and with complementary \\hi\\ synthesis mapping programs of similar galaxies in crowded cluster and group fields, it is clearly now feasible  to explore the evolution of the \\hi\\ content of disk galaxies over the last several Gyr for a statistically complete sample.\\\\"
    },
    "0808/0808.3185_arXiv.txt": {
        "abstract": "We extend a method for modeling synchrotron and synchrotron self-Compton radiations in blazar jets to include external Compton processes. The basic model assumption is that the blazar radio through soft X-ray flux is nonthermal synchrotron radiation emitted by isotropically-distributed electrons in the randomly directed magnetic field of outflowing relativistic blazar jet plasma. Thus the electron distribution is given by the synchrotron spectrum, depending only on the Doppler factor $\\delta_{\\rm D}$ and mean magnetic field $B$, given that the comoving emission region size scale $R_b^\\prime \\lesssim c \\dD t_v/(1+z)$, where $t_v$ is variability time and  $z$ is source redshift. Generalizing the approach of Georganopoulos, Kirk, and Mastichiadis (2001) to arbitrary anisotropic target radiation fields, we use the electron spectrum implied by the synchrotron component to derive accurate Compton-scattered $\\gamma$-ray spectra throughout the Thomson and Klein-Nishina regimes for external Compton scattering processes.  We derive and calculate accurate $\\gamma$-ray spectra produced by relativistic electrons that Compton-scatter (i) a point source of radiation located radially behind the jet, (ii) photons from a thermal Shakura-Sunyaev accretion disk and (iii) target photons from the central source scattered by a spherically-symmetric shell of broad line region (BLR) gas. Calculations of broadband spectral energy distributions from the radio through $\\gamma$-ray regimes are presented, which include self-consistent $\\gamma\\gamma$ absorption on the same radiation fields that provide target photons for Compton scattering. Application of this baseline flat spectrum radio/$\\gamma$-ray quasar model is considered in view of data from $\\gamma$-ray telescopes and contemporaneous multi-wavelength campaigns. ",
        "introduction": "A class of radio-loud blazar active galactic nuclei (AGNs) that emit luminous fluxes of $\\gtrsim 100$ MeV -- GeV $\\gamma$ rays was discovered with the Energetic Gamma Ray Experiment Telescope (EGRET) on the {\\em Compton Observatory} \\citep{har92,fic94,har99}. This result clarified the nature of 3C 273, which was first identified as a $\\gamma$-ray emitting AGN in COS-B satellite data \\citep{her77}.  The $\\gamma$ rays from blazars are certainly nonthermal in origin and associated with the radio jets formed by the supermassive black holes that power these sources. The largest subclass of EGRET AGNs are moderate redshift ($z\\approx 1$) flat-spectrum radio quasars (FSRQs) with blazar properties, including apparent superluminal motion, rapidly variable optical emission, high polarization, and intense broadened optical emission lines.  Another subclass of the EGRET AGNs consists of $\\gamma$-ray emitting BL Lacertae (BL Lac) objects, which are generally at lower redshifts ($z \\approx 0.1$ -- $0.3$) and, by definition, have weak or absent optical emission lines in their spectra. The X-ray selected BL Lac (XBL) subset were discovered to be a class of TeV $\\gamma$-ray sources by the Whipple Observatory \\citep{pun92,wee03}.  The common superluminal nature of the first identified $\\gamma$-ray blazars, namely 3C 273, 3C 279, and PKS 0528+134, led \\citet{dsm92} to propose a Compton-scattering origin for the $\\gamma$ rays. In this model, jet electrons Compton-scatter accretion-disk photons that intercept the jet plasma. The nonthermal jet electrons can also scatter internal synchrotron photons to produce a synchrotron self-Compton (SSC) component \\citep{bm96}. Given the broadened emission lines in the spectra of FSRQs, accretion-disk radiation scattered by surrounding gas of the broad line region (BLR) will provide a further source of target photons to be scattered to $\\gamma$-ray energies \\citep{sbr94}, as will radiation from a surrounding dusty torus \\citep{kat99,bsmm00}.  The accretion-disk and scattered radiation will attenuate jet $\\gamma$-rays through $\\gamma\\gamma$ pair-production attenuation \\citep{bk95,bl95}.  Expressions for the $\\gamma$-ray spectral energy distributions (SEDs) of blazars produced by Compton scattering processes have been derived and calculated for many specific models of the black-hole/blazar jet environment. In the case of external accretion-disk photons as the target photon source, where the accretion disk is described by an optically thick, geometrically thin thermal \\cite{ss73} accretion disk, Compton-scattered $\\gamma$-ray spectra were calculated in the Thomson regime by \\citet{ds93,ds02}. Calculations of the Thomson-scattered spectra for a quasi-isotropic target radiation field formed by BLR gas or hot dust were made by \\citet{sbr94}, \\cite{dss97} and \\citet{bsmm00}.  Detailed numerical calculations including both accretion disk and scattered radiation fields, have been made by, e.g., \\citet{kt05,bb00} and \\citet{br04}.  Compton scattering in the Klein-Nishina regime is not so simple to treat compared to analyses restricted to the Thomson regime, but is unavoidable for blazar analysis in the era of the {\\em Fermi Gamma-Ray Space Telescope} (formerly known as {\\em GLAST}) and the ground-based $\\gamma$ Cherenkov telescopes. For surrounding isotropic radiation fields in the stationary frame of the blazar AGN, \\citet{gkm01} suggested to transform the comoving electron distribution to the stationary frame and then scatter the target photons to $\\gamma$-ray energies, using the formula first derived by Jones \\citep{jon68,bg70}.  This approach is generalized in this paper to surrounding anisotropic radiation fields.  The usual spectral modeling approach proceeds by injecting power-law electrons and evolving these particles while they produce the output synchrotron and Compton-scattered radiation \\citep[e.g.,][]{ds93,bms97}. For example, \\citet{mod05} calculate electron energy evolution and spectral formation throughout the Thomson and Klein-Nishina regimes for different ratios of synchrotron and isotropic radiation field energy densities. They show that reduced Compton losses in the Klein-Nishina regime compared to synchrotron losses can lead to spectral hardening of the synchrotron component in the optical/X-ray regime \\citep[noted earlier by][]{da02}. A difficulty in this approach is that the electron energy-loss rate depends on the photon spectrum of the comoving radiation field, not just the total radiation energy density, and this field evolves with time. The modeler is faced with the prospect of simultaneously fitting the synchrotron and Compton components.  The acceleration scenario may well be over-simplified, and non-power-law particle injection distributions could be more realistic than power-law injection spectra, e.g., due to nonlinear effects in Fermi acceleration. Moreover, a separation between the acceleration and radiation zones may not be justified.  Here we extend a method of blazar analysis recently proposed for TeV blazars \\citep{fdb08} that avoids these difficulties.  For a standard $\\gamma$-ray blazar model, where isotropically distributed electrons spiral in a randomly oriented magnetic field with mean magnetic field strength $B$ in the fluid frame, the measured synchrotron flux directly reveals the electron spectrum responsible for the synchrotron radiation. The only uncertainties are the mean magnetic field $B$, the comoving size scale $R_b^\\prime$ of the emitting region, and the Doppler factor $\\dD = [\\Gamma (1-\\beta\\cos\\theta)]^{-1}$ ($\\G = 1/\\sqrt{1-\\beta^2}$ is the bulk Lorentz factor of the outflow).  With this electron spectrum, we then Compton-scatter target photons of the surrounding radiation fields using the head-on approximation to the total Compton cross section \\citep{ds93}, valid when the electron Lorentz factor $\\gamma \\gg 1$. This generalizes the approach of \\citet{gkm01} to surrounding anisotropic radiation fields.  The temporally-evolving electron spectrum in blazars can be derived in this approach from simultaneous multiwavelength blazar data. Values of $B$, $\\dD$, and jet power can then be deduced.  The related treatment for XBLs applied to PKS 2155-304, including more details about the derivation of the electron spectrum from the synchrotron component, the derivation and calculations of the SSC component and internal $\\gamma$-ray opacity by the synchrotron photons, is given by \\citet{fdb08}.  Analysis of blazar SEDs using this approach is presented in Section 2, where formulas to calculate Compton-scattered internal and external radiation and a $\\delta$-function approximation for $\\gamma\\gamma$ opacity from the internal radiation field are given. Derivations of the Compton-scattered spectrum for specific examples of external radiation fields consisting of a monochromatic point source of radiation radially behind the jet, a Shakura-Sunyaev disk model, and a model BLR radiation field are derived in Section 3.  Discussion of the results is found in Section 4. ",
        "conclusions": "We have presented accurate expressions for modeling synchrotron and Compton-scattered radiation from the jets of AGN that include target radiation fields from the accretion disk and BLR. This extends our technique for modeling synchrotron and SSC emission \\citep{fdb08} to include external Compton scattering in a relativistic jet of thermal radiation from the accretion disk, and accretion-disk radiation Thomson-scattered by electrons in the BLR.  These formulae use the full Compton cross section and are accurate throughout the Thomson and Klein-Nishina regimes at any angle with respect to the jet axis, so can also be used to model $\\gamma$-ray emission from radio galaxies, e.g., M87 \\citep{aha06}. We also derive expressions for the $\\g\\g$ opacity through the same scattered radiation field that serves as a target photon source for the jet electrons. The expressions, eqs.\\ (\\ref{tauggSStilder}) and (\\ref{taugge1r}), for opacity from the accretion-disk and scattered radiation field assume, however, that the photon travels along the jet axis, though it is straightforward to derive the more general case.  In the results presented here (Figs.\\ \\ref{f1} -- \\ref{f6} and \\ref{f10} -- \\ref{f12}) we have chosen parameters for demonstration purposes that give an exaggerated EC component, particularly a high $\\ell_{\\rm Edd}$.  Lowering the disk luminosity lowers the radiation considerably, as seen in Fig.\\ \\ref{f2}.  The disk radiation field in FSRQ blazars can be seen when the nonthermal blazar radiation is in a low state, as in the cases of 3C 279 \\citep{pia99}, 3C 454.3 \\citep{rai07} and, most clearly, 3C 273 \\citep[e.g.,][]{mon97}. These observations can be used to assign the accretion-disk luminosity when modeling a specific blazar, though the disk brightness could also vary during the flaring epoch.  The models presented here do, however, have limitations.  They do not yet include enhancements from secondary cascade radiation initiated by $e^\\pm$ pairs formed by $\\g$ rays interacting with lower energy radiation from the disk and the BLR. For the parameters considered here, this would not make a significant difference in the calculated SEDs because the energy flux of the attenuated radiation is a small fraction of the escaping flux. But even in this case, the reinjected pairs from the attenuated radiation will be isotropized if the reinjection occurs outside the relativistic flow, and will then make only a small contribution to the Doppler-boosted radiation.  Spectral features result from $\\gamma\\gamma$ absorption by BLR radiation, as seen in Fig.\\ \\ref{f12}. By assuming hard primary $\\gamma$-ray emission components, \\citet{aha08} argue that  hard intrinsic spectra from blazars such as 1ES 1101-232 \\citep{aha07-1101} could be formed through $\\gamma\\gamma$ attenuation. If the primary TeV radiation originates from an underlying jetted electron distribution, then a consistent model requires that a $\\gamma$-ray spectrum formed by Compton-scattering processes arises from the same radiation field responsible for $\\gamma\\gamma$ absorption. As our calculations show, soft Compton-scattered TeV spectra are formed due to Klein-Nishina effects in scattering, whether from the accretion disk or from photons scattered by the BLR. Thus either an extremely bright GeV component would be expected in the scenario of \\citet{aha08} (cf.\\ Figs.\\ \\ref{f11} and \\ref{f12}), which would easily be detected with {\\em Fermi}, or a hard primary $\\gamma$-ray spectrum must originate from other processes. Moreover, if this explanation was correct, then blazars with stronger broad emission lines should have harder VHE $\\gamma$-ray spectra than weak-lined BL Lacs.  Cascade radiation induced by ultra-relativistic hadrons could inject high-energy leptons to form a hard radiation component, though a sufficiently dense target field for efficient photohadronic losses will itself severely attenuate the TeV radiation \\citep{ad03}. Depending on the underlying acceleration model, a distinctive hadronic signature could appear at $\\gamma$-ray energies only, though one would expect that both electrons and protons would be accelerated synchronously.   Our calculations also do not as yet include absorption by the diffuse extragalactic background light (EBL), which would be important for EGRET $\\gamma$-ray FSRQs, which have a broad redshift distribution with a mean value $\\langle z\\rangle \\sim 1$, larger than the mean value $\\langle z\\rangle \\sim 0.3$ for BL Lacs \\citep{muk97}.  In cases where the Compton-scattered spectra decline steeply due to Klein-Nishina effects (e.g., Fig.\\ \\ref{f1}), the issue of EBL absorption is secondary.  In the one FSRQ, \\object{3C 279} at $z=0.536$, detected from 80 to $> 300$ GeV energies with MAGIC \\citep{mag08}, EBL effects cannot be neglected.  For the models studied here, the soft calculated $\\gamma$-ray spectra cannot in any case account for the measured, let alone intrinsic, spectrum of 3C 279.  Based on the simultaneous optical -- X-ray -- VHE $\\gamma$-ray SED of 3C279 during the MAGIC detection, B\\\"ottcher, Marscher \\& Reimer (2008, in preparation) show that one-zone leptonic models have severe problems explaining the flare, and require either extremely high Doppler factors or magnetic fields well below equipartition. Correlated X-ray, {\\em Fermi}, and TeV campaigns will offer opportunities to apply the results developed in this paper.  In conclusion, we have derived expressions to model FSRQ blazars that self-consistently include $\\g\\g$ attenuation on the same target photons that are Compton-scattered by the relativistic electrons.  By assuming that the lower energy, radio -- UV emission is nonthermal synchrotron radiation, then the underlying electron distribution can be determined given the Doppler factor, magnetic field, and size scale of the emission region. The equations derived here can be used to calculate the $\\gamma$-ray emission spectrum from SSC and external Compton scattering processes from this electron distribution. Separately, one can determine whether the inferred  electron distribution can be derived from specific acceleration scenarios. Complementary to the technique of injecting electron spectra and cooling, this method can be applied to multiwavelength data sets to analyze high-energy processes in the jets of AGN."
    },
    "0808/0808.4136_arXiv.txt": {
        "abstract": "{ We present two-dimensional (axisymmetric) neutrino-hydrodynamic simulations of the long-time accretion phase of a 15$\\,M_\\odot$ progenitor star after core bounce and before the launch of a supernova explosion, when non-radial hydrodynamic instabilities like convection occur in different regions of the collapsing stellar core and the standing accretion shock instability (SASI) leads to large-amplitude oscillations of the stalled shock with a period of tens of milliseconds. Our simulations were performed with the {\\sc Prometheus-Vertex} code, which includes a multi-flavor, energy-dependent neutrino transport scheme and employs an effective relativistic gravitational potential. Testing the influence of a stiff and a soft equation of state for hot neutron star matter, we find that the non-radial mass motions in the supernova core impose a time variability on the neutrino and gravitational-wave signals with larger amplitudes, as well as higher frequencies in the case of a more compact nascent neutron star. After the prompt shock-breakout burst of electron neutrinos, a more compact accreting remnant produces higher neutrino luminosities and higher mean neutrino energies. The observable neutrino emission in the SASI sloshing direction exhibits a modulation of several ten percent in the luminosities and around 1$\\,$MeV in the mean energies with most power at typical SASI frequencies between roughly 20 and 100$\\,$Hz. The modulation is caused by quasi-periodic variations in the mass accretion rate of the neutron star in each hemisphere. At times later than $\\sim$50--100$\\,$ms after bounce, the gravitational-wave amplitude is dominated by the growing low-frequency ($\\la$200$\\,$Hz) signal associated with anisotropic neutrino emission. A high-frequency wave signal results from nonradial gas flows in the outer layers of the anisotropically accreting neutron star. Right after bounce such nonradial mass motions occur due to prompt post-shock convection in both considered cases and contribute mostly to the early wave production around 100$\\,$Hz. Later they are instigated by the SASI and by convective overturn that vigorously stir the neutrino-heating and cooling layers, and also by convective activity developing below the neutrinosphere. The gravitational-wave power then peaks at about 300--800$\\,$Hz, connected to changes in the mass quadrupole moment on a timescale of milliseconds. Distinctively higher spectral frequencies originate from the more compact and more rapidly contracting neutron star. Both the neutrino and gravitational-wave emission therefore carry information that is characteristic of the properties of the nuclear equation of state in the hot remnant. The detectability of the SASI effects in the neutrino and gravitational-wave signals is briefly discussed. ",
        "introduction": "\\label{sec:introduction}  Neutrinos and gravitational waves are the most direct potential probes of the processes that occur deep inside of a dying star, accompanying or causing the initiation of the stellar explosion. Neutrinos were already detected in connection with supernova SN~1987A (Bionta et al.\\ 1987, Hirata et al.\\ 1987, Alexeyev et al.\\ 1988), although with poor statistics so that the extraction of information for constraining the explosion mechanism was not possible. The simultaneous measurement of signals of both types remains a very realistic hope for the next Galactic supernova.  Capturing neutrinos and gravitational waves from the same source has the advantage of providing complementary insight into the conditions of the stellar core. While neutrino signals reflect the density structure and thermodynamic conditions in the high-density plasma of the collapsing core and forming neutron star, gravitational waves carry crucial information about the dynamics and nonradial motions of the stellar matter, e.g.\\ of its rotational state or of hydrodynamic instabilities such as convection that deform and stir the condensing central compact remnant of the explosion.  The prompt burst of electron neutrinos, for example, signals the breakout of the supernova shock from the neutrinosphere and the shock heating of matter in this region. The neutrino emission after core bounce is a sensitive probe of the mass accretion rate on the forming neutron star, hence of the density structure of the infalling layers of the dying star (e.g., Liebend\\\"orfer et al.\\ 2003, Buras et al.\\ 2006b). The onset of the explosion is expected to show up as a more or less sudden drop in the mass accretion rate of the forming neutron star. A soft supernuclear equation of state will lead to a more compact and hotter remnant, radiating higher neutrino luminosities and more energetic neutrinos (e.g., Janka et al.\\ 2005, Marek 2007). And a possible phase transition to non-nucleonic matter in the neutron star core may impose characteristic features like a second neutrino burst (Sagert et al.\\ 2008) or, if triggering the collapse to a black hole, may cause an abrupt termination of the neutrino emission (e.g., Burrows 1988; Keil \\& Janka 1995; Sumiyoshi et al.\\ 2006, 2007; Fischer et al.\\ 2008).  Theoretical work on the gravitational-wave signals from stellar core collapse and explosion has a long history of successively refined numerical models. In particular the infall and bounce phases, which are theoretically relatively well understood parts of the evolution, have received a lot of interest, because a strong and characteristic signature could make them a promising source of a detectable gravitational-wave burst (for a review-like introduction to the subject and a nearly complete list of publications, see Dimmelmeier et al.\\ 2008). For this to be the case, the core of the progenitor star must develop a sufficiently large deformation during its infall, and for that it must rotate enough rapidly. This, however, does not seem to be compatible with predictions from the latest generation of stellar evolution models for the vast majority of massive stars, which are expected to have lost most of their angular momentum before collapse (Heger, Woosley, \\& Spruit 2005). Only in rare, very special cases, possibly accounting for the few tenths of a percent of all stellar core collapses that produce gamma-ray bursts, such stars seem to be able to retain a high angular momentum in their core and to thus produce relativistic jets and highly asymmetric and extremely energetic, probably magnetohydrodynamically driven explosions (for a recent review, see e.g. Woosley \\& Bloom 2006). If this hypothetical connection was true, the core bounce phase would not really offer grand perspectives for the measurement of gravitational waves.  In contrast, only relatively little work has so far been done on determining the wave signals from the post-bounce evolution of a supernova, although these signals are likely to carry important information about the still incompletely understood explosion mechanism and the associated core dynamics (see the review by Ott 2008). M\\\"uller \\& Janka (1997) and later M\\\"uller et al.\\ (2004) on the basis of significantly improved numerical models, showed that in the case of delayed, neutrino-driven explosions convective overturn behind the stalled shock as well as convection inside the nascent neutron star can account for sizable gravitational-wave signals, which should be detectable with a high probability from a Galactic supernova when the Advanced Laser Interferometer Gravitational-Wave Observatory (LIGO~II) is running (see also Fryer et al.\\ 2002, 2004).  While the delayed neutrino-heating mechanism relies on the support by strong nonradial hydrodynamic instabilities in the region of neutrino-energy deposition, magnetohydrodynamic explosions, if linked to rapid rotation, are expected to develop relatively soon after core bounce and thus to occur faster than the growth of the mentioned nonradial instabilities. They tap the reservoir of differential rotation in a collapsing stellar environment and therefore require rapidly spinning cores (see Burrows et al.\\ 2007b, Thompson et al.\\ 2005). Thus they are good candidates for sizable gravitational-wave pulses from the moment of core bounce (e.g., Kotake et al.\\ 2006, Ott et al.\\ 2004). The newly proposed acoustic explosion mechanism, in which large-amplitude core gravity-mode oscillations of the neutron star convert accretion power to pressure and shock waves that feed the supernova shock with acoustic energy (Burrows et al.\\ 2006, 2007a), is predicted to be associated with enormous and very characteristic gravitational-wave activity due to the fast periodic movement of roughly a solar mass of dense matter at late times ($\\sim$1~second) after core bounce (Ott et al.\\ 2006).  In the present paper our focus is on an analysis of the features in the neutrino and gravitational-wave signals that might provide evidence for the action of the so-called standing accretion shock instability (SASI), which has been shown to lead to low-$\\ell$ mode (in terms of an expansion in spherical harmonics with order $\\ell$), large-amplitude nonradial shock deformation and violent sloshing motions of the stalled supernova shock (Blondin et al.\\ 2003; Blondin \\& Mezzacappa 2007; Bruenn et al.\\ 2006; Scheck et al.\\ 2004, 2006, 2008; Ohnishi et al.\\ 2006, Foglizzo et al.\\ 2007; Yamasaki \\& Foglizzo 2008). This SASI activity is found to reach the nonlinear regime at roughly 100$\\,$ms after core bounce and to grow in amplitude over possibly hundreds of milliseconds (Scheck et al.\\ 2008; Marek \\& Janka 2007; Burrows et al.\\ 2006, 2007a). It does not only act as seed of powerful secondary convection but can also provide crucial aid for the neutrino-heating mechanism by pushing the accretion shock to larger radii and by thus stretching the time accreted matter is exposed to neutrino heating in the gain layer (Buras et al.\\ 2006b, Scheck et al.\\ 2008, Marek \\& Janka 2007, Murphy \\& Burrows 2008). Moreover, the SASI is found to play the driving force of the g-modes pulsations of the neutron star core that are the essential ingredient of the acoustic mechanism.  Observational signatures of the presence of the SASI would therefore be extremely important for our understanding of how massive stars begin their explosion. Supernova asymmetries and pulsar kicks are one, yet not unambiguous observational hint. Neutrinos and gravitational waves, which originate directly from the region where the blast is initiated, however, may remain the only way to obtain direct information. It is therefore a highly relevant question to ask whether any characteristic structures are imprinted on the neutrino and gravitational-wave emission by the SASI activity in the supernova core. In contrast to convection, which exhibits the fastest growth for the higher-$\\ell$ modes (see Foglizzo et al.\\ 2006), the SASI is expected to possess the largest growth rates for the dipolar and quadrupolar deformations (corresponding to $\\ell = 1,\\,2$; Blondin \\& Mezzacappa 2006; Foglizzo et al.\\ 2007; Yamasaki \\& Foglizzo 2008; Foglizzo 2001, 2002). Even in the fully nonlinear situation the geometry and motion of the shock and post-shock layer are found to be governed by these lowest modes.  We present here an analysis of the neutrino and gravitational-wave signals that are calculated on the basis of the two-dimensional post-bounce and pre-explosion simulations of a 15$\\,M_\\odot$ star recently published by Marek \\& Janka (2007). We constrain ourselves to two nonrotating models, which allow us to discuss the differences that can be expected from a stiff and a soft nuclear equation of state (results of the corresponding 1D simulations can also be found in Marek \\& Janka 2007). We find that the SASI sloshing of the shock and the associated quasi-periodic mass motions lead to a time-modulation of the neutrino emission and to gravitational-wave amplitudes whose size and characteristic frequency depend on the compactness of the proto-neutron star during the first half of a second after core bounce. Though the SASI contributions to the power spectra are superimposed by a significant high-frequency ``noise'' due to convective fluctuations, at least the combination of measurements should make it possible to identify the SASI activity in the supernova core.  Our paper is organized as follows. In Sect.~\\ref{sec:numerics} we will briefly outline the main numerical and physics ingredients of our simulations, and the basic features of the two simulations we compare. In Sect.~\\ref{sec:results} we discuss our results with respect to the SASI effects on the shock motion, neutrino emission, and gravitational-wave signal, and in Sect.~\\ref{sec:conclusions} we will summarize our findings and draw conclusions, including a discussion whether the SASI modulations are detectable by neutrino and gravitational-wave experiments. ",
        "conclusions": "\\label{sec:conclusions}   We have presented results of two stellar core-collapse simulations performed with the {\\sc Prometheus-Vertex} code for a nonrotating 15$\\,M_\\odot$ progenitor star in axial symmetry. One of these models was computed with a soft version of the Lattimer \\& Swesty (1991) nuclear equation of state, and the other one with the stiffer Hillebrandt \\& Wolff (1985) EoS, which produces a significantly less compact proto-neutron star. For these simulations we have compared the neutrino and gravitational-wave signals between core bounce and 400$\\,$ms later. In particular we were interested in signal features that may yield information about the action of the standing accretion shock instability (SASI) in the supernova core during the long phase of post-bounce accretion that precedes the onset of the supernova explosion in all recent multi-dimensional calculations of stellar core collapse.  Our main results can be summarized as follows: \\begin{itemize} \\item[(1)] In both simulations we found that the weakened bounce shock produces a negative entropy gradient that partly overlaps with the negative lepton number gradient left behind by the escaping shock breakout burst of electron neutrinos. This region is Ledoux unstable and develops prompt post-shock convection, which leads to an early burst of gravitational-wave emission lasting several ten milliseconds after core bounce. Its frequency spectrum shows a prominent peak around 70--100$\\,$Hz. \\item[(2)] A later, long-lasting phase of strong gravitational-wave emission with a long-time trend in growing wave amplitudes sets in at $t \\ga 100\\,$ms after bounce. It is produced by nonradial mass motions and anisotropic neutrino emission that are caused by convective overturn inside the proto-neutron star and in the neutrino-heating layer behind the stalled shock, and in particular by vigorous sloshing of the accretion shock due to the development of low-$\\ell$ mode SASI activity. The wave train for the matter signal consists of a superposition of quasi-periodic variations on timescales from a few milliseconds to several ten milliseconds. The spectra of the dominant $\\ell = 1,\\,$2 SASI modes peak between about 20 and roughly 100$\\,$Hz, and the gravitational-wave emission also shows significant power at frequencies up to $\\sim$200$\\,$Hz. The most prominent maxima of the gravitational-wave spectra, however, are located at higher frequencies of 300--800$\\,$Hz and originate mainly from the outer layers of the neutron star where rapid gas flows are triggered by the supersonic impact and deceleration of accretion funnels. \\item[(3)] The high-frequency peak of the gravitational-wave spectrum and less clearly also the secondary maximum near the SASI frequency exhibit a dependence on the compactness of the nascent neutron star and thus on the properties of the high-density EoS: the more compact neutron star leads to more powerful shock oscillations and significantly larger gravitational-wave amplitudes earlier after core bounce. Also the characteristic frequencies of the low-$\\ell$ SASI modes and of the gravitational-wave signal are higher. While the main peak of the wave spectrum is located at 300--600$\\,$Hz for the stiffer EoS, it can be found at 600--800$\\,$Hz for the softer EoS with the more compact remnant. \\item[(4)] The more compact neutron star also radiates neutrinos with higher luminosities and greater mean energies. Anisotropic neutrino emission produces a low-frequency component of the gravitational-wave signal at $f \\la 100$--200$\\,$Hz with an amplitude that dominates the matter signal by up to a factor of two. The emission anisotropy is mostly established in the cooling and heating regions of the accretion layer outside of the neutrinospheres and is determined by the anisotropic transport of muon and tau neutrinos through this layer, which leads to a slightly higher flux of muon and tau neutrinos near the equatorial plane of the polar grid, i.e., perpendicular to the direction of the SASI oscillations of the supernova shock. Despite the stronger emission of $\\nu_e$ and $\\bar\\nu_e$ from the accretion layer in the polar regions (i.e.\\ in the direction of the SASI shock expansion and contraction), the gravitational-wave amplitude associated with the anisotropic energy loss in neutrinos exhibits a nearly monotonic trend to negative values. \\item[(5)] The neutrino luminosities and mean energies, in particular the ones that can be received by an observer in the polar (and thus SASI) direction, show a quasi-periodic time variability with an amplitude of several ten percent (up to about 50\\%) of the minimum values for the luminosities and of roughly 1$\\,$MeV (up to 10\\%) for the mean energies. The luminosity fluctuations are correlated and in phase with the energy variations and are somewhat bigger for $\\nu_e$ and $\\bar\\nu_e$ than for heavy-lepton neutrinos. They originate from alternating phases of gas accumulation and compression in the accretion layer in particular around the poles, caused by the expansion and contraction of the shock in the course of the SASI oscillations. The frequency spectra of the neutrino luminosities show most power between about 20 and 200$\\,$Hz and significant power in a decaying tail up to about 400$\\,$Hz as a consequence of emission variability associated with faster, convective modulations of the accretion flow between supernova shock and nascent neutron star. Since SASI and convective activity occur inseparably in the gain layer and stir the cooling layer below, power is found to be distributed over a wide range of frequencies in the Fourier spectra of the luminosities. \\end{itemize}  The SASI induced temporal modulations of the luminosities and mean energies are features by which neutrinos can provide direct evidence of the dynamical processes that occur in the supernova core. Only few other cases are known where neutrinos carry such information, for example the prompt flash of electron neutrinos as a signature of the shock breakout from the neutrinosphere, or a possible second neutrino burst or abrupt termination of the neutrino emission, which are expected if the neutron star collapses to a more compact object or a black hole. A detection of the predicted characteristic, quasi-periodic variation in the neutrino event rate from a Galactic supernova would confirm the action of the SASI in the supernova core and would thus provide extremely valuable support of our present understanding of the core-collapse physics: the SASI is discussed as possible origin of global supernova asymmetries and pulsar kicks and as a potentially crucial agent on the route to the explosion.  Moreover, nonradial mass motions such as convection and the SASI in the supernova core after bounce produce gravitational-wave emission with significant intensity. We found that the wave spectrum exhibits a broad hump between about 20 and 200$\\,$Hz and a main maximum with clearly more power at considerably higher frequencies. The latter is associated with SASI and accretion induced gas flows in the neutron star surface layers, where much more mass is involved than in the fairly dilute post-shock region. Both the primary and the secondary maximum of the spectrum are sensitive to the structure of the neutron star: larger gravitational-wave amplitudes and correspondingly higher spectral maxima as well as higher frequencies are obtained with a softer equation of state and more compact remnant. The relative SASI variations of the neutrino luminosities do not appear to be very sensitive to the nuclear equation of state, and a clear EoS-dependent difference in the modulation frequencies is also not visible, but the absolute values of the luminosities and mean energies of the neutrinos radiated by the more compact and hotter, accreting proto-neutron star in our simulations are 10--20\\% higher.  Anisotropic neutrino emission is also an important source of gravitational waves with amplitudes that can dominate the matter signal. The corresponding spectrum possesses most power at frequencies below 100--200$\\,$Hz. The emission anisotropy is a function with rapid short-time variability (with periods of milliseconds to $\\ga$10$\\,$ms) and with an overlying, nearly monotonic long-time trend toward slightly stronger emission from equatorial regions, i.e., perpendicular to the direction of the SASI shock oscillations. This trend decides about the sign of the neutrino gravitational-wave amplitude. A detailed analysis of our models reveals that it is the result of complex transport physics in which the neutrinospheric flux, which has a low pole-to-equator anisotropy, is modified strongly in the inhomogeneous environment of the accretion layer around the forming neutron star. The pole-to-equator asymmetry changes with radius several times in both directions and the contributions from $\\nu_e$, $\\bar\\nu_e$, and heavy-lepton neutrinos partly cancel each other. A reasonable determination of these transport asymmetries, which is not only important for gravitational-wave calculations but also for pulsar kick estimates based on neutrino recoil, requires detailed transport calculations. Even our sophisticated ray-by-ray approach may not be sufficient for accurate predictions of measurable effects, but a full multi-angle treatment of the transport may be necessary.  From this discussion it is clear that any reliable assessment of the detectability of SASI effects in the neutrino and gravitational-wave emission will ultimately require results from three-dimensional models. The growth of SASI modes with nonvanishing azimuthal order $m$ was shown to be of potential importance, and the presence of such nonaxisymmetric asymmetries can have a big influence on the structure and the dynamics of the accretion layer (Blondin \\& Mezzacappa 2007, Yamasaki \\& Foglizzo 2008, Iwakami et al.\\ 2008). It is probable that also the emission asymmetries and time modulations of the emitted signals exhibit significant differences compared to 2D models, in particular because in 3D the SASI axis is not forced to coincide with an axis of symmetry and with the axis of the numerical grid.  The SASI and convective activity, its duration as well as strength, must be expected to vary with the progenitor star and to depend on the length of the postbounce accretion phase until the onset of the explosion (for hints to that, see the results for 11.2$\\,M_\\odot$ and 15$\\,M_\\odot$ explosions in Marek \\& Janka 2007). In the nonlinear stage, the SASI and convective motions are of chaotic nature and therefore stochastic differences also in the large-scale pattern and global asymmetry are likely to develop from small initial differences even in progenitor stars (models) of the same kind (see, e.g., the different explosion asymmetries obtained in a large set of 2D simulations performed by Scheck et al.\\ 2006). Nevertheless, basic features of the instabilites, in particular their mere presence and long-time growth behavior and the sizable modulation they impose on the emitted signals over a period of possibly hundreds of milliseconds, should be common to different realizations. Judged from our present understanding, however, a template-like uniformity of the signal structure with a well defined frequency (or frequency range) and amplitude that exhibit a regular dependence on the explosion and progenitor properties, appears not to be very likely. More model calculations are therefore needed to find out which signal properties such as the modulation frequency, are sufficiently generic to yield information about the neutron star equation of state.   \\subsection{Detectability of SASI modulations}  Should the SASI signatures and their magnitude discussed in this work be confirmed by 3D results, the question arises whether they will be measurable with existing or emerging neutrino and gravitational-wave experiments in the case of the next Galactic supernova. A detailed prediction of supernova neutrino signals in terrestrial detectors is a highly complex problem, because the neutrinos radiated by the nascent neutron star undergo flavor conversions so that the flavor arriving at the Earth is not necessarily the same as the one leaving the source. On the one hand, the flavor evolution is affected by neutrino-matter oscillations during the neutrino propagation through the density gradients in the supernova mantle and envelope (where the H and L resonances are encountered at densities around 10$^3\\,$g$\\,$cm$^{-3}$ and 10$\\,$g$\\,$cm$^{-3}$, respectively) and through the Earth. On the other hand, neutrino conversion is also induced by the nonlinear self-interaction associated with neutrino-neutrino forward scattering in the dense neutrino background close to the neutron star (for an application of these phenomena in the context of the neutrino signals from oxygen-neon-magnesium core supernovae, see Lunardini et al.\\ 2008, and from black hole formation in stellar collapse events, see Nakazato et al. 2008; for a discussion of the flavor conversion physics, see the references to original publications in the latter two papers). The flavor composition of the neutrino burst in the detector and thus the detector response therefore do not only depend on the type and structure of the progenitor star but also on a number of degrees of freedom associated with unknown neutrino properties and the location of the experiment, e.g.\\ the neutrino-mass hierarchy (normal or inverted), the still undetermined mixing angle between the mass eigenstates 1 and 3, and the nadir angle of the neutrino path through the Earth.  Because of the involved complexity we defer an investigation of the flavor conversion effects and their consequences for the neutrino measurements to a future paper. However, since the SASI modulations affect the luminosities and mean energies of neutrinos and antineutrinos of all flavors with similar strength, we can ignore the flavor oscillation effects for a simple estimate to convince us that the SASI modulations of the radiated signal should be detectable in principle. Super-Kamiokande, for example, is expected to capture roughly 8000 electron antineutrinos through their absorption by protons, if a Galactic supernova occurs at a distance of 10$\\,$kpc and releases at total energy of $5\\times 10^{52}\\,$erg in $\\bar\\nu_e$ with an average energy of 15$\\,$MeV. A luminosity of $4\\times 10^{52}\\,$erg$\\,$s$^{-1}$, which is a typical value for the emission maxima during the SASI phase, should account for an event rate of approximately 7000 $\\bar\\nu_e$ captures per second, corresponding to about 70 events in the $\\sim$10$\\,$ms intervals of peak emission. If $L_\\nu\\langle\\epsilon_\\nu^2\\rangle \\langle\\epsilon_\\nu\\rangle^{-1}$ is about 50\\% lower in luminosity minima, about one half of this event rate and event number can be expected during the low-emission periods. Therefore the SASI modulations will lead to a variability of the measured signal that can be larger, though not much, than the statistical $\\sqrt{N_{\\nu,{\\mathrm{det}}}}$ fluctuations, $N_{\\nu,{\\mathrm{det}}}$ being the number of detected neutrinos in a time interval of 10$\\,$ms. For Hyper-Kamiokande, whose fiducial volume would be roughly 30 times bigger than that of Super-Kamiokande, the Poisson noise would be reduced to a few percent of the event number within 10$\\,$ms, so that the SASI modulations should create a clear signature. Even more powerful for this purpose is the IceCube experiment. It will be able to identify the supernova neutrino burst by the Cherenkov glow caused by the neutrino energy deposition in the ice, which will show up as time-correlated noise in all phototubes. For a supernova at a distance of 10$\\,$kpc, the $\\bar\\nu_e$ luminosity maxima during the SASI phase are expected to cause a photon count rate of more than $10^6$ per second in the 4800 optical modules (see Dighe et al.\\ 2003). The $\\sim$$1.2\\times 10^4$ events within 10$\\,$ms intervals and the variations of several 1000 events between periods of luminosity peaks and minima (in the case of IceCube these variations scale with $L_\\nu\\langle\\epsilon_\\nu^3\\rangle \\langle\\epsilon_\\nu\\rangle^{-1}$, because the crucial quantity is the energy deposition in the detector) dominate by far the Poisson fluctuations associated with the background counting rate of 300$\\,$Hz per optical module; the Poisson variability of the background corresponds to $\\sqrt{N_\\gamma}\\sim 120$ photon events in all phototubes, which is $\\sim$1\\% of the supernova signal in a 10$\\,$ms interval (an impression of the detection capability even for a supernova model with a lower fluctuation amplitude of the neutrino emission than in the models discussed here, can be obtained from the upper lefthand panel of Fig.~2 in Dighe et al.\\ 2003).  The detectability of gravitational-wave signals very similar to the ones presented in this work was investigated in much detail by M\\\"uller et al.\\ (2004). Due to the restriction to a 90-degree grid (one hemisphere) in the case of the 15$\\,M_\\odot$ core-collapse model s15r considered in the latter paper, the contributions of dipolar SASI modes were not included in there. Moreover, since the collapse evolution was followed to later times in our most recent simulations (about 400$\\,$ms instead of $\\sim$270$\\,$ms after bounce) and the quadrupole amplitudes exhibit a trend in rise during the later post-bounce phases, the signal-to-noise (S/N) ratios and maximum detection distances given for LIGO~I and Advanced LIGO by M\\\"uller et al.\\ (2004) are likely to be only lower limits. Since the S/N ratio scales roughly linearly with the Fourier transform of the wave amplitude, the detection prospects grow with longer integration periods of the signal and thus higher Fourier power. The measurement of dominant gravitational-wave power at frequencies of several hundred Hz with growing strength over a period of hundreds of milliseconds would therefore be a strong indication of SASI and convective activity around the nascent neutron star. M\\\"uller et al.\\ (2004) estimated a S/N ratio of roughly 70 in Advanced LIGO for a 15$\\,M_\\odot$ supernova at 10$\\,$kpc, corresponding to the possibility of capturing gravitational waves from such sources up to a distance of roughly 100$\\,$kpc. A third-generation underground interferometer facility (Einstein Telescope) would lead to a gain in sensitivity of approximately a factor of 20 in the most relevant frequency range between 100$\\,$Hz and 1000$\\,$Hz (Spiering 2007). This would increase the expected detection rate of supernovae by up to a factor of 3 (Ando et al.\\ 2005). A detailed analysis of the gravitational-wave signals and their observational implications from a greater sample of core-collapse models will be presented in a forthcoming paper."
    },
    "0808/0808.2595_arXiv.txt": {
        "abstract": "Recent astrophysical data indicates that dark matter shows a controversial behaviour in galaxy cluster collisions. In case of the notorious Bullet cluster, dark matter component of the cluster behaves like a collisionless system. However, its behaviour in the Abell 520 cluster indicates a significant self-interaction cross-section. It is hard for the WIMP based dark matter models to reconcile such a diverse behaviour. Mirror dark matter models, on the contrary, are more flexible and for them diverse behaviour of the dark matter is a natural expectation. ",
        "introduction": "The evidence of dark matter is compelling at all observed astrophysical scales \\cite{1-1,1-2}. At that the dark matter reveals itself only through its gravitational interactions. There are no observational facts that indicate existence of any significant interactions between the ordinary and dark matter particles. However, some recent astrophysical observations show that dark matter particles may have significant self-interactions. This is surprising because other observations support the WIMP paradigm that the dark matter is essentially collisionless. The mirror dark matter \\cite{1-3, 1-4,1-5}, first introduced in \\cite{1-KOP}, has a potential to reconcile these seemingly contradictory observational facts. For other attempts to do this, on the base of BEC dark matter, see \\cite{1-6}. Below, after briefly commenting the mirror matter idea, we consider the above mentioned observational facts and argue why, in our opinion, they indicate the existence of the mirror matter as the main dark matter component. ",
        "conclusions": "There is a growing astrophysical evidence of diverse behaviour of dark matter in galaxy cluster collisions. These observations indicate that the mirror dark matter models deserve careful examination and observational verification. If the mirror dark matter were as popular as the SUSY dark matter, we would say that it is already discovered. However, it would be more fair to conclude that we need more observational evidence to firmly prove this fascinating conjecture."
    },
    "0808/0808.2240_arXiv.txt": {
        "abstract": "\\vspace{1cm} \\centerline{\\bf ABSTRACT}\\vspace{2mm} Recently, Gamma-Ray Bursts (GRBs) were proposed to be a complementary cosmological probe to type Ia supernovae (SNIa). GRBs have been advocated to be standard candles since several empirical GRB luminosity relations were proposed as distance indicators. However, there is a so-called circularity problem in the direct use of GRBs. Recently, a new idea to calibrate GRBs in a completely cosmology independent manner has been proposed, and the circularity problem can be solved. In the present work, following the method proposed by Liang {\\it et al.}, we calibrate 70 GRBs with the Amati relation using 307 SNIa. Then, following the method proposed by Shafieloo {\\it et al.}, we smoothly reconstruct the cosmic expansion history up to redshift $z=6.29$ with the calibrated GRBs. We find some new features in the reconstructed results. ",
        "introduction": "\\label{sec1} The current accelerated expansion of our universe~\\cite{r1} has been one of the most active fields in modern cosmology since its discovery in 1998 from the observations of type Ia supernovae (SNIa)~\\cite{r2}. Later, the observations of cosmic microwave background (CMB) anisotropy~\\cite{r3} and large-scale structure~\\cite{r4} confirmed this discovery. Although today there are already many observational methods, SNIa have been proved to be one of the most powerful tools to probe this mysterious phenomenon. In 1992, the famous Phillips relation was found~\\cite{r5}, which claimed that for nearby SNIa there exists a clear correlation between their intrinsic brightness at maximum light and the duration of their light curve. Since then, some empirical technics~\\cite{r6,r7,r8} have been developed to use the Phillips relation to make SNIa into standard candles. See e.g.~\\cite{r9,r10} for brief historical reviews.  However, SNIa are plagued with extinction from the interstellar medium. Hence, the currently maximum redshift of SNIa is only about $z\\simeq 1.755$. As argued in~\\cite{r10}, the observations at $z>1.7$ are fairly important to distinguish cosmological models. On the other hand, the redshift of the last scattering surface of CMB is $z\\simeq 1090$. So, the observations at intermediate redshift are important. Recently, Gamma-Ray Bursts (GRBs) were proposed to be a complementary probe to SNIa. So far, GRBs are the most intense explosions observed in our universe. Their high energy photons in the gamma-ray band are almost immune to dust extinction. Up to now, there are many GRBs observed at $0.1<z\\leq 8.1$, whereas the maximum redshift of GRBs is expected to be $10$ or even larger~\\cite{r11}. Strictly speaking, GRBs are not standard candles, with radiated energies spanning several orders of magnitude. But the use of empirical luminosity correlations to standardize them as distance indicators has been proposed by several authors. In some sense, this is reminiscent of the case of the Phillips relation for SNIa. These empirical correlations include the Amati relation~\\cite{r12}, those derived from it (for examples, the Ghirlanda relation~\\cite{r13}, the Yonetoku relation~\\cite{r14}, the Liang-Zhang relation~\\cite{r15}, the Firmani relation~\\cite{r16}), and others~\\cite{r17,r18,r19,r20}. Therefore, the so-called GRB cosmology emerges recently. We refer to e.g.~\\cite{r21,r22,r23,r10} for comprehensive reviews.  To our knowledge, in~\\cite{r24} GRBs were used to constrain $\\Omega_m$ and dark energy for the first time. Similar works also include~\\cite{r25} for examples. However, there is a so-called circularity problem in the direct use of GRBs~\\cite{r21}. In this case, to calibrate the empirical GRB luminosity relations, one need to assume a particular cosmological model with some model parameters {\\it a priori}, mainly due to the lack of a set of low redshift GRBs at $z<0.1$ which are cosmology independent. When one uses these ``calibrated'' GRBs (which are actually cosmological model dependent) to constrain cosmological models, the circularity problem occurs. To alleviate the circularity problem, some statistical methods have been proposed, such as the scatter method~\\cite{r26}, the luminosity distance method~\\cite{r26}, and the Bayesian method~\\cite{r27}. These methods were used extensively in the literature~\\cite{r15,r28,r29}. However, they still cannot solve the circularity problem {\\em completely}. Another method trying to avoid the circularity problem was proposed in~\\cite{r30}, in which the parameters of the empirical GRB luminosity relation and the cosmological model parameters were fitted to the observational data simultaneously. However, for {\\em any} given cosmological model, this method can {\\em always} obtain some parameters for the cosmological model and the empirical GRB luminosity relation. In this sense, {\\em any} cosmological model is ``viable'' (except for a few obviously absurd models). So, it cannot be used to rule out any cosmological model. Therefore, it is also an unsatisfactory method to solve the circularity problem completely.  To overcome the circularity problem completely, one should calibrate GRBs in a cosmology independent manner. Due to the lack of a set of low redshift GRBs at $z<0.1$ which are cosmology independent, it was proposed to calibrate the empirical GRB luminosity relation using a sufficient number of GRBs within a small redshift bin centered around any redshift~\\cite{r31} (see also e.g.~\\cite{r21}). However, this method might be unrealistic, since the current sample of observed GRBs is not large enough. Recently, a new idea to calibrate GRBs in a completely cosmology independent manner has been proposed in~\\cite{r32,r33} independently. Similar to the case of calibrating SNIa as secondary standard candles by using Cepheid variables which are primary standard candles, we can also calibrate GRBs as standard candles with a large amount of SNIa. This idea of the distance ladder was also briefly mentioned in~\\cite{r34}. However, in fact, \\cite{r34} used the $\\Lambda$CDM model rather than SNIa data themselves to calibrate GRBs and hence the calibration is still cosmology dependent actually. In~\\cite{r32}, the GRB luminosity relation was calibrated with the empirical formula of the luminosity distance of SNIa, while the physical meaning of this empirical formula is still unclear and its reliability should be tested carefully. Instead, in~\\cite{r33} the GRB luminosity relation was calibrated with the interpolated distance moduli from the Hubble diagram of SNIa themselves. We refer to the original papers~\\cite{r32,r33} for details. Using these cosmology independently calibrated GRBs, we can constrain the parameters of cosmological models without circularity problem. In fact, the constraints on $\\Lambda$CDM model and $w_0-w_a$ parameterized dark energy model have been considered in~\\cite{r32,r33,r35}.  In the present work, we will calibrate GRBs to reconstruct the cosmic expansion history up to redshift $z=6.29$. The calibration method of GRB relation adopted here is the one proposed in~\\cite{r33}. However, we will use the datasets of SNIa and GRBs which are larger than the ones used in~\\cite{r33}. So, more calibrated high redshift GRBs can be available with smaller uncertainties. In Sec.~\\ref{sec2}, we use 38 GRBs at $z<1.4$ to calibrate the Amati relation while the distance moduli of these 38 low redshift GRBs are interpolated from 307 SNIa data~\\cite{r36,r37}. Then, the distance moduli of 32 GRBs whose $z>1.4$ can be derived from the calibrated Amati relation. In Sec.~\\ref{sec3}, following the method proposed in~\\cite{r38,r39}, we smooth 307 SNIa data and 32 calibrated GRBs data whose $z>1.4$ to reconstruct the cosmic expansion history up to $z=6.29$. Conclusion and discussions are briefly given in Sec.~\\ref{sec4}.   \\begin{center} \\begin{figure}[htbp] \\centering \\includegraphics[width=0.99\\textwidth]{calnew.eps} \\caption{\\label{fig1} Left panel: The Hubble diagram of 307 SNIa (black diamonds) and 38 low redshift GRBs (red stars) whose distance moduli are derived by using cubic interpolation. Right panel: 38 GRBs data (red stars) in the $\\log E_{\\rm p,i}\\,/({\\rm 300\\,keV})-\\log E_{\\rm iso}\\,/{\\rm erg}$ plane. The best-fit calibration line is also plotted. See text for details.} \\end{figure} \\end{center}   \\vspace{-10mm} % ",
        "conclusions": "\\label{sec4} In this work, we used the cosmology independent method proposed in~\\cite{r33} to calibrate the GRBs and derived the distance moduli $\\mu$ for 32 high redshift GRBs at $z>1.4$, whereas the numerical results are given in Table~\\ref{tab1}. We have used the 307 SCP Union SNIa compilation~\\cite{r36} and the 70 GRBs compiled in~\\cite{r41}. Since these GRBs are calibrated in a completely cosmology independent manner, one can use the calibrated 32 GRBs at $z>1.4$ to constrain cosmological models without circularity problem completely. On the other hand, following the method proposed by Shafieloo {\\it et al.}~\\cite{r38}, in this work we smoothly reconstructed the cosmic expansion history up to $z=6.29$ using 307 SNIa compiled in~\\cite{r36} and 32 calibrated GRBs at $z>1.4$. It is worth noting that this method is also model independent.  Here are some comments on the technical details~\\cite{r74}. Firstly, SNIa as standard candles are affected by several potential problems, e.g., dust extinction, color evolution, reliability of the Phillips relation up to high $z$, uncertainty on the progenitors, possible existence of different sub-classes, etc. Thus, calibrating GRB with SNIa propagates these uncertainties and sources of systematics also into the calibrated spectrum-energy correlations. In addition, even when using a cosmology independent method, in this way GRBs are no more an ``independent'' cosmological probe. In other words, by calibrating with SNIa one can avoid the circularity problem intrinsic in the use of GRB spectrum-energy correlations and increase the accuracy in the determination of spectral parameters. However, with respect to using spectrum-energy correlations alone, one might introduce more systematics and introduce a ``circularity'' with SNIa themselves. Secondly, in the present work, we used the bisector of the two ordinary least squares~\\cite{r47} to derive the best-fit parameters of the calibrated $E_{\\rm p,i}-E_{\\rm iso}$ correlation. In fact, we closely followed the method used by Schaefer~\\cite{r10}. In this method, one changed the value of $\\sigma^2_{E_{\\rm iso,sys}}$ {\\em by hand} to force the $\\chi^2/dof$ of the 38 points at $z<1.4$ to be $1$. One might consider that this is not a robust method. As a good alternative, we refer to the maximum likelihood method used, e.g., by Amati {\\it et al.}~\\cite{r41}.  As mentioned above, the reconstructed $H(z)$, $q(z)$ and $w_{tot}(z)$ are interesting in some sense. To make these results robust, we have tried various values of $\\Delta$ and various initial guess models. The resulted $H(z)$, $q(z)$ and $w_{tot}(z)$ are still similar to the previous ones. The features in the reconstructed $H(z)$, $q(z)$ and $w_{tot}(z)$ at high redshift $z>2$ remain. Some remarks are in order. Firstly, one might doubt the validity of using GRBs as standard candles. If GRBs are not real standard candles, the reconstructed results at $z>1.4$ are unreliable of course. In fact, the debate is not settled in the GRB community today. Some authors argued that GRBs cannot be used as standard candles, see~\\cite{r49,r50,r51,r52,r53,r70,r71,r72,r73} for examples. On the other side, some authors advocated that GRBs can be used as standard candles to probe cosmology, see~\\cite{r10,r11,r12,r13,r14,r15,r16,r17,r18,r19,r20,r21,r22, r23,r24,r25,r26,r27,r28,r29,r30,r31,r32,r33,r34,r35,r54,r55, r56,r57,r58,r59,r60,r61,r62,r63} for instance. Today, the situation is still unclear and many authors choose to wait and see. Secondly, if the GRBs can be used as standard candles, since the sample of available GRBs at high redshifts is still not large enough, some bias might exist in the reconstructed $H(z)$, $q(z)$ and $w_{tot}(z)$ and lead to the interesting features. So, before more and better high redshift GRBs are available, we cannot say anything conclusively. Finally, on the other side, some hints were found for oscillating $H(z)$, $q(z)$ and $w_{tot}(z)$ in the literature, see e.g.~\\cite{r64,r65,r66,r67,r68} and references therein. So, the reconstructed results obtained here might be meaningful in some sense and deserve further investigations.  Obviously, due to the large scatter and the lack of a large amount of well observed GRBs, there is a long way to use GRBs extensively and reliably to probe cosmology. The Fermi Gamma-ray Space Telescope~\\cite{r69} which was launched recently might change this situation. The GRB sample will be appreciably larger, and the new GRBs will be better studied in some sense. In the coming Fermi era, we hope that the GRB cosmology could have a bright future."
    },
    "0808/0808.0245_arXiv.txt": {
        "abstract": "We present self-similar solutions for advection -dominated accretion flows with thermal conduction in the presence of outflows. Possible effects of outflows on the accretion flow are parametrized and a saturated form of thermal conduction, as is appropriate for the weakly-collisional regime of interest, is included in our model. While the cooling effect of outflows is noticeable, thermal conduction provides an extra heating source. In comparison to accretion flows without winds, we show that the disc rotates faster and becomes cooler  because of the angular momentum and energy flux which are taking away by the winds. But thermal conduction opposes the effects of winds and not only decreases the rotational velocity, but increases the temperature. However, reduction of the surface density and the enhanced  accretion velocity are amplified by both of the winds and the thermal conduction. We find  that for stronger outflows, a higher level of saturated thermal conduction is needed to significantly modify the physical profiles of the accretion flow. ",
        "introduction": "The thin accretion disk model describes flows in which the viscous heating of the gas radiates out of the system immediately after generation (Shakura \\& Sunyaev 1973). However, another kind of accretion has been studied during recent years where radiative energy losses are small so that most of the energy is advected with the gas. These Adection-Dominated Accretion Flows (ADAF) occur in two regimes depending on the mass accretion rate and the optical depth. At very high mass accretion rate, the optical depth becomes very high and the radiation can be trapped in the gas. This type of accretion which is known under the name 'slim accretion disk' has been studied in detail by Abramowicz et al. (1988). But when the accretion rate is very small and the optical depth is very low, we may have another type of accretion (Narayan \\& Yi 1994; Abramowitz et al. 1995; Chen 1995). However, numerical simulations of radiatively inefficient accretion flows revealed that low viscosity flows are convectively unstable and convection strongly influences the global properties of the flow (e.g., Igumenshchev, Abramowicz, \\& Narayan 2000). Thus, another type of accretion flows has been proposed, in which convection plays as a dominant mechanism of transporting angular momentum and the local released viscous energy (e.g., Narayan, Igumenshchev, \\& Abramowicz 2000).  This diversity of models tells us that modeling the hot accretion flows is a challenging and controversial problem. We think, one of the largely neglected physical ingredient in this field, is {\\it thermal conduction}. But a few authors tried to study the role of \"turbulent\" heat transport in ADAF-like flows (Honma 1996; Manmoto et al. 2000). Since thermal conduction acts to oppose the formation of the temperature gradient that causes it, one might expect that the temperature and density profiles for accretion flows in which thermal conduction plays a significant role to appear different compared to those flows in which thermal conduction is less effective. Just recently, Johnson \\& Quataert (2007) studied the effects of electron thermal conduction on the properties of hot accretion flows, under the assumption of spherical symmetry. In another interesting analysis, Tanaka \\& Menou (2006) showed that thermal conduction affects  the global properties of hot accretion flows substantially. They generalized standard ADAF solutions to include a saturated form of thermal conduction. In the second part of their paper, a set of two dimensional self-similar solutions of ADAFs in the presence of thermal conduction has been presented. The role of conduction is in providing the extra degree of freedom necessary to launch thermal outflows according to their 2D solutions.   On the other hand, ADAFs with winds or outflows have been studied extensively during recent years, irrespective of possible driven mechanisms of winds. But thermal conduction has been neglected in all these ADAFs solutions with winds. In advection-dominated inflow-outflow solutions (ADIOS), it is generally assumed that the mass flow rate has a power-law dependence on radius, with the power law index, $s$, treated as a parameter (e.g., Blandford \\& Begelman 1999; Quataert \\& Narayan 1999; Beckert 2000; Misra \\& Taam 2001; Fukue 2004; Xie \\& Yuan 2008). Beckert (2000) presented self-similar solutions for ADAFs with radial viscous force in the presence of outflows from the accretion flow or infall. Turolla \\& Dullemond (2000) investigated how, and to what extent, the inclusion of the source of ADAF material affects the Bernoulli number and the onset of a wind. In their model, the accretion rate decreases with radius ($s<0$). Misra \\& Taam (2001) studied the effect of a possible hydrodynamical wind on the nature of hot accretion disc solutions. They showed that their solutions are locally unstable to a new type of instability called wind-driven instability, in which the presence of a wind causes the disc to be unstable to long-wavelength perturbations of the surface density. Kitabatake, Fukue \\& Matsumoto (2002) studied supercritical accretion disc with winds, though angular momentum loss of the disc, because of the winds, has been neglected. Comparisons with observations reveal that the X-ray spectra of such a wind-driven self-similar flow,  can explain the observed spectra of black hole candidates in quiescence (Quataert \\& Narayan 1999; Yuan, Markoff \\&  Falcke 2002). Lin, Misra \\& Taam (2001) find that the spectral characteristics of high-luminosity black hole systems suggest that winds may be important for these systems as well.   Considering extensive works on hot accretion flows with winds and significant role of thermal conduction in deriving outflows (Tanaka \\& Menou 2006), we study ADAFs with outflows {\\it and} thermal conduction using height-integrated set of equations. A phenomenological way is adopted in which we parameterize the rate at which mass, angular momentum and energy are extracted by outflow or wind. In the next section, we present basic equations of the model. Self-similar solutions are investigated in section 3. The paper concludes with a summary of the results in section 4. ",
        "conclusions": "Theoretical arguments and observations suggest that mass loss via winds may be important in sub -Eddington, radiatively inefficient, accretion flows. On the other hand, thermal conduction may play a significant role in such systems (Tanaka \\& Menou 2006). Using a simplified model, we included both winds and thermal conduction in unified model in order to understand their possible combined effects on the dynamics of the system. Accounting for a variable mass accretion rate in the flow in proportion to $R^{s}$ and a saturated form of thermal conduction with coefficient $\\phi_{\\rm s}$, a set of self-similar solutions are presented in this study.  We have varied $s$ and $\\phi_{\\rm s}$ in our models to judge their sensitivity to these parameters. The most important finding of our analysis is that as more mass, angular momentum and energy flux are taking away from the disc (i.e. stronger winds), possible modifications of the physical profiles due to the thermal conduction will occur at higher level of conduction.  Although our self-similar solutions are too simplified to be used to calculate the emitted spectrum, their typical behavior show importance of thermal conduction in the presence of winds. For future the global solution rather than the self-similar solutions is required. This is because most of the radiation of an ADAF comes from its innermost region where the self-similar solution breaks down."
    },
    "0808/0808.1769_arXiv.txt": {
        "abstract": "We present the results of deep spectroscopy for the central region of the cluster lens SDSS~J1004+4112 with the Subaru telescope. A secure detection of an emission line of the faint blue stellar object (component E) near the center of the brightest cluster galaxy (G1) confirms that it is the central fifth image of the lensed quasar system. In addition, we measure the stellar velocity dispersion of G1 to be $\\sigma_*=352\\pm13$ km~s$^{-1}$. We combine these results to obtain constraints on the mass $M_{\\rm BH}$ of the putative black hole (BH) at the center of the inactive galaxy G1, and hence on the $M_{\\rm BH}$-$\\sigma_*$ relation at the lens redshift $z_l=0.68$. From detailed mass modeling, we place an upper limit on the black hole mass,  $M_{\\rm BH}<2.1\\times10^{10}M_\\odot$ at 1$\\sigma$ level ($<3.1\\times10^{10}M_\\odot$ at 3$\\sigma$),  which is consistent with black hole masses expected from the local and redshift-evolved  $M_{\\rm BH}$-$\\sigma_*$ relations, $M_{\\rm BH}\\sim 10^9-10^{10}M_\\odot$. ",
        "introduction": "Quasars lensed by foreground clusters of galaxies serve as a powerful probe of the mass distributions of clusters. SDSS J1004+4112 is the first example of such quasar-cluster lens systems \\citep{inada03,oguri04,ota06}. It consists of four bright quasar images ($z=1.734$) with a maximum image separation of $14\\farcs6$ produced by a massive cluster at $z=0.68$.  In addition to the quasar images, multiply-imaged background galaxies are also observed \\citep{sharon05}. An advantage for using lensed quasars is that they can provide time delays, which have been actually detected for this system \\citep{fohlmeister07,fohlmeister08}, containing unique information on the lens potential.  What makes the quasar lens SDSS~J1004+4112 particularly unique is the probable central fifth image (component E) of the lensed quasar system. High-resolution {\\it Hubble Space Telescope (HST)} Advanced Camera for Survey (ACS; \\cite{clampin00}) images show a blue point source located $0\\farcs2$ from the center of the brightest cluster galaxy G1 \\citep{inada05}. The central lensed image, if confirmed, has several important implications for the central structure of lensing objects. For instance, central images can constrain the inner mass profiles of lensing galaxies, particularly the masses of the central supermassive black holes (e.g., \\cite{mao01}; \\cite{rusin01}; \\cite{keeton03}; \\cite{rusin05}).  While the correlation between the black hole mass ($M_{\\rm BH}$) and the stellar velocity dispersions ($\\sigma_*$) of the host galaxies has been established for local galaxies \\citep{ferrarese00,gebhardt00,tremaine02}, measurements of the redshift evolution for the correlation, derived for galaxies hosting active galactic nuclei, are still controversial (e.g., \\cite{peng06}; \\cite{shen08}; \\cite{woo08}). Thus independent constraints from central images are thought to provide insights into this relation.  In this {\\it Letter}, we present results of two deep spectroscopic follow-up observations for SDSS J1004+4112, conducted at the Subaru 8.2-meter telescope. One is spectroscopy of component E to confirm its lensing nature, and the other is spectroscopy of galaxy G1 to measure its velocity dispersion. The latter observation is particularly important in separating the mass distribution of dark matter from that of baryons (e.g., \\cite{sand08}). We then attempt mass modeling of the lens system. Together with the measurement of $\\sigma_*$ for galaxy G1, it provides a direct constraint on the $M_{\\rm BH}$-$\\sigma_*$ relation at the lens redshift of $z=0.68$. In what follows, we assume a standard flat universe with $\\Omega_M=0.26$, $\\Omega_\\Lambda=0.74$, and $H_0=72{\\rm km\\,s^{-1}Mpc^{-1}}$ (e.g., \\cite{tegmark06}). ",
        "conclusions": "We have presented the results of two spectroscopic observations at the Subaru telescope. With the first observation, we confirmed that the central point source found by \\citet{inada05} is indeed the central fifth image of the lensed quasar by detecting the C\\emissiontype{IV} emission line of the quasar. This represents the first spectroscopic confirmation of the central odd image. In the second spectroscopic observation, we determined the stellar velocity dispersion of galaxy G1 to be $\\sigma_*=352\\pm13$~${\\rm km\\,s^{-1}}$.  We used these results to derive constraints on the $M_{\\rm BH}$-$\\sigma_*$ relation at the lens redshift $z=0.68$, assuming the presence of a BH at the center of the brightest cluster galaxy G1. With a parametric model which successfully reproduces the model constraints, we obtained limits of $M_{\\rm  BH}<2.1\\times10^{10}M_\\odot$ at 1$\\sigma$ and $M_{\\rm BH}<3.1\\times10^{10}M_\\odot$ at 3$\\sigma$ that are consistent with $M_{\\rm BH}$ expected from the extrapolation of the known $M_{\\rm BH}$-$\\sigma_*$ relation. It is worth noting that these constraints are derived for an ``inactive'' galaxy with no nuclear activity. Current studies of the redshift evolution of the $M_{\\rm BH}$-$\\sigma_*$ relation make use of the Balmer line widths, and therefore are restricted to BHs in active galaxies. Given possible differences in the $M_{\\rm BH}$-$\\sigma_*$ between active and inactive local galaxies \\citep{greene06}, constraints on $M_{\\rm BH}$ for distant inactive galaxies from gravitational lensing will be essential to fully understand the origin of the relation.  \\bigskip  N.~I. acknowledges support from the Special Postdoctoral Researcher Program of RIKEN. This work was in part supported by Department of Energy contract DE-AC02-76SF00515. C.~S.~K. acknowledges support from NSF grant AST 07-08082. G.~P.~S acknowledges support from a Royal Society University Research Fellowship. E.~F. acknowledges support from the Smithsonian Institution."
    },
    "0808/0808.1844_arXiv.txt": {
        "abstract": "We test four commonly used astrophysical simulation codes; Enzo, Flash, Gadget and Hydra, using a suite of numerical problems with analytic initial and final states. Situations similar to the conditions of these tests, a Sod shock, a Sedov blast and both a static and translating King sphere occur commonly in astrophysics, where the accurate treatment of shocks, sound waves, supernovae explosions and collapsed haloes is a key condition for obtaining reliable validated simulations. We demonstrate that comparable results can be obtained for Lagrangian and Eulerian codes by requiring that approximately one particle exists per grid cell in the region of interest. We conclude that adaptive Eulerian codes, with their ability to place refinements in regions of rapidly changing density, are well suited to problems where physical processes are related to such changes. Lagrangian methods, on the other hand, are well suited to problems where large density contrasts occur and the physics is related to the local density itself rather than the local density gradient. ",
        "introduction": "Over the last decade, numerical simulations have developed into a cornerstone of astrophysical research. From interactions of vast clusters of galaxies to the formation of proto-planetary discs, simulations allow us to evolve systems through time and view them at every angle. They are an irreplaceable test-bed of our physical understanding of the Universe.   A number of hydrodynamics codes have been developed that are widely used in this field and still more are being developed. Fundamentally, they all do the same job; they solve the equations of motion to calculate the evolution of matter through time. Whether the matter represents a nebula for the birth of a star or a network of galaxy clusters, the basic technique remains the same.   However, the algorithms used to solve these equations vary from code to code and this results in differences in the resulting data. Understanding the origin of these variations is vital to the understanding of the results themselves; is an observed anomaly an interesting piece of new physics or a numerical effect? As observational data takes us deeper into the Universe, it becomes more important to pin down the origin of these numerical artifacts.   Additionally, it is difficult to compare results from simulations run with different codes. With observations, papers clearly state the properties of the instrument such as the diameter of the mirror and the wavelengths it is most sensitive to. While a brief description of the code is always included in theoretical papers, there exists no obvious conversion to other numerical techniques and therefore the results are more difficult for the reader to interpret.   The problem of code comparison is not new and it is a topic that has recently created a great deal of interest. The reason for its current importance is a positive one; improved numerical techniques and increased computer power have resulted in simulations reaching greater resolutions than could have been imagined even a few years ago. However, this high refinement comes at a price; as we start to pick out the detail of these complex fluid flows, the physics we need to consider gets dramatically more complicated.  This brings us to the main question code comparison projects are trying to answer; can we use the same tools for this new regime of problems?   A number of papers have come out that tackle this. One of the most famous is the Santa Barbara Comparison Project \\citep{Frenk1999} which compared many of the progenitors of today's codes by running a model of a galaxy cluster forming. Taking a different tack \\citet{Thacker} compared the performance of a dozen different implementations of a single approach (in this case smoothed particle hydrodynamics, SPH) on standard astrophysical problems that included the Sod shock, examining the range of outcomes that were available from a single technique. They concluded that one of the weaknesses of SPH was the weak theoretical grounding which allows several equally viable formulations to be derived. More recent work includes \\citet{Agertz2007}, who focus specifically on the formation of fluid instabilities, comparing the formation of Kelvin-Helmholtz and Rayleigh-Taylor instabilities in six of the most utilised codes. Additionally, \\citet{OShea2005} has completed a direct comparison between two particular codes ({\\it Enzo} and {\\it Gadget2}, see below) looking at the formation of galaxies in a cosmological context.   All these projects give detailed insights into the differences between the codes, but are unable to provide a quantitative measure of how well a code performs in a particular aspect. This is especially true of the cosmological-based tests of \\citet{Frenk1999, OShea2005} where the problem is not sufficiently well-posed for convergence onto a single answer. \\citet{Agertz2007} sets out a simpler problem and compares it to analytical predictions, but the system is still sufficiently complex not to have an exact solution. Additionally, no previous comparison has attempted to quantitatively compare different codes to one another; asking whether it is possible to obtain identical results and with what conditions. Without this crucial piece of information, it is impossible to fully assess a piece of work performed by an unfamiliar code or to judge which code might be the most suited to a given problem type. This has resulted in somewhat general comments being made about the differences between numerical techniques which has led to many myths about a code's ability becoming accepted dogma.   The set of tests we present in this paper are designed to tackle these difficulties with the intention that they might become part of an established test programme all hydrodynamical codes should attempt. We present four problems that specifically address different aspects of the numerical code all of which have expected `correct' answers to compare to. The first two tests, the Sod shock tube and Sedov blast, are both strong shock tests with analytical solutions. The third and forth tests concern the stability of a galaxy cluster and are primarily tests of the code's gravitational solver. For all four tests we directly compare the codes against the analytic solution and present an estimate of the main sources of any systematic error.   The remainder of this paper is organised as follows: in section 2 we give a short summary of the main features of each of the four codes we have employed. In section 3 we deal with the Sod shock and Sedov blast tests, setting out the initial and final states and comparing each code against them. We repeat this exercise for a static and translating King sphere in section 4. Finally we discuss our results and summarise our conclusions in section 5. ",
        "conclusions": "We have presented four well posed tests with known solutions which we have used to compare four well used astrophysical codes; {\\it Enzo, Flash, Gadget2} and {\\it Hydra}. We have examined each code's ability to resolve strong shocks and incorporate an accurate gravity solver capable of resolving forces over many orders of magnitude in density.  The tests presented here were specifically designed to be difficult and push the codes to their limit; appropriately so since many situations in astrophysics demand extreme physical conditions. Despite these requirements, all the codes did pass our tests satisfactorily and what we highlight here is the strengths and weaknesses of the different techniques.  Significant issues were {\\it Enzo (Zeus)}'s failure to produce a spherical shock, a problem that leads to a lack of energy conservation and to poor recovery of the position of a Sedov blast. For extreme shocks the artificial viscosity implementation of {\\it Gadget2} leads to spurious entropy driven bubbles and for very static configurations the centre-of-mass can slowly drift due to the difficulty of recovering the zeroth order term in the tree force. For a translating cluster, both implementations of {\\it Enzo} appear to lose momentum over time.  All our tests lead to the conclusion that in order to achieve roughly the same resolution in grid and particle codes a good rule of thumb is to require one particle per cell in the region of interest. This is difficult to achieve for the SPH codes over shock fronts or voids and likewise harder for AMR codes to match in the centre of collapsed objects. This is reflected strongly in our results where the AMR codes largely excel at the shock tests in section~\\ref{sec:shock} and the SPH codes at the gravitationally bound cluster tests in section~\\ref{sec:cluster}. For a cosmological simulation these criteria would equate to having the minimum mesh cell size about the same size as the gravitational softening or half the SPH search length.  As implied by the relation above, for strong shocks SPH codes require significantly more effort than AMR codes. This is because AMR codes have the ability to add extra resolution in regions of rapidly changing density, whereas SPH codes can effectively only refine with the density itself, so they end up undertaking lots of unnecessary work far from the shock where the forces are small in regions of high but uniform density. However, for models such as the King sphere, where AMR codes traditionally refine using a density criterion anyway, SPH codes can achieving high resolution far more efficiently in these dense structures because they do this very naturally by following individual particles. For AMR codes to achieve similar results, high resolution and a significant increase in computational time is required. In these cases AMR codes would still have an advantage if it was necessary to resolve the low density region, where the SPH codes simply run out of particles and have to smooth over very large volumes.  Additional care must be taken for the parallel implementations of {\\it Hydra} and {\\it Enzo}. {\\it Hydra}'s parallel version was not used in this paper and further testing would be needed to ensure that it performs correctly. The gravity solver in {\\it Enzo} was demonstrated to produce different results in serial and parallel in the King model tests (although not in other situations) and this should be explicitly tested when running a simulation with this code. An updated gravity solver for {\\it Enzo} which rectifies this problem is currently being developed. While serial runs allow much parameter space exploration, rapid code development and inter-comparison, true state-of-the-art large simulations are confined to codes that can utilise high performance computing facilities and successfully run on a large number of processors.  To a large extent the choice of simulation code largely comes down to the problem being attacked. For problems with large dynamic range or when the object is rapidly translating across the volume particle based methods are much less computationally expensive in order to achieve the same result. Conversely, for problems where the interesting physics is in regions of rapidly changing density rather than in high density regions themselves, AMR codes excel thanks to their more adaptive refinement criteria.  The computational time required for each code to perform these tests varied greatly. Ideally, elapsed time is largely irrelevant since the required resolution and problem size should be set by the appropriate physics being tackled rather than available computer resources. However, we include Table~{\\ref{table:times}} for completeness and to give the reader an indication of times involved. In Table~\\ref{table:times}, the computational systems referred to are; `UF HPC' is the University of Florida's high performance computing center which uses (typically) InfiniBand-connected 2.2 GHz AMD Opteron duel cores with 6 Gb RAM. `UF Astro' is the University of Florida Astronomy's minicluster with 2.66 GHz Intel quad-cores and 4 Gb RAM. `Columbia' is the Columbia University Astronomy's cluster with gigabit-connected 2.2 GHz Opteron cores and 3 Gb RAM. `Miranda' is at the University of Durham and consists of Myrinet-connected 2.6 GHz Opteron cores with 4 Gb RAM. `Nottingham' is the Nottingham University HPC with gigabit-connected 2.2 GHz Opteron cores with 2 Gb of RAM.  \\begin{table*} \\caption{Computational time to perform the presented tests, including the number of processors used and the type of system. Readers should note that comparing CPU times between simulations performed on disparate computers is fraught with danger. For true runtime performance, the codes would need to be bench marked on identical systems which was beyond the scope of this paper. We therefore include this table purely for guidance and warn against over interpretation. $\\dagger$ The times quoted for the static and translating cluster runs in {\\it Flash} are estimated total runtimes based on the speed up attained with the latest version of the FFTW gravity solver. The original runs were performed using an early version of the hybrid FFTW gravity solver which was still under development and as such was not optimised. The original runs took place on 24 processors and runtimes for the pre-development level FFTW solver were 217 hrs 42 mins for the static cluster and 564 hrs 5mins for the translating cluster. This drastic improvements arises from a decrease in the memory requirements and improved communication patterns. A reduction in computational time in the Sedov blast test was additionally achieved by using a simplified Riemann solver that was not required to follow multiple fluids with different adiabatic indices (a feature the default configuration of {\\it Flash} has implemented to model thermonuclear flashes). This reduced the memory overheads.} \\begin{tabular}{lccl} \\hline {\\bf Riemann Shock Tube: Planar} & hh:mm & no. processors & system  \\\\ \\hline {\\it Enzo (PPM)} & 00:36 & 16 & UF HPC (Opteron) \\\\ {\\it Enzo (Zeus)} & 04:00 & 1 & UF Astro (Intel Xeon) \\\\ {\\it Flash} & 00:24 & 32 & Miranda (Opteron) \\\\ {\\it Hydra} ($10^6$ particles) & 00:36  & 1 & Nottingham (Opteron)\\\\ {\\it Gadget2} ($10^6$ particles) & 01:31 & 2 & Nottingham (Opteron) \\\\ \\hline {\\bf Riemann Shock Tube: Oblique} & & &\\\\ \\hline {\\it Enzo (PPM)} & 01:22 & 16 &  UF HPC  (Opteron) \\\\ {\\it Enzo (Zeus)} & 05:06 & 8 & UF HPC (Opteron) \\\\ {\\it Flash} & 00:26 & 45 & Miranda (Opteron)  \\\\ {\\it Hydra} ($10^6$ particles) & 00:36 & 1 & Nottingham (Opteron) \\\\ {\\it Gadget2} ($10^6$ particles)  & 01:30 & 2 & Nottingham (Opteron) \\\\ \\hline {\\bf Sedov Blast Test} &  &  &  \\\\ \\hline {\\it Enzo (PPM)} & 00:24 & 16 & UF HPC (Opteron) \\\\ {\\it Enzo (Zeus)} & 04:41 & 6 & UF HPC (Operton) \\\\ {\\it Flash} & 15:33 & 16 & Miranda (Opteron) \\\\ {\\it Hydra} & 06:49 & 1 & Nottingham (Opteron) \\\\ {\\it Gadget2} & 19:41 & 4 & Nottingham (Opteron) \\\\ \\hline {\\bf Static Cluster Test} & &  &  \\\\ \\hline {\\it Enzo (PPM)} & 61:50 & 1 & UF HPC (Opteron) \\\\ {\\it Enzo (Zeus)} & 23:49 & 1 & UF HPC (Opteron) \\\\ {\\it Flash} & 111:01 $\\dagger$ & 12 & Miranda (Opteron) \\\\ {\\it Hydra} & 37:51 & 1 & Nottingham (Opteron) \\\\ {\\it Gadget2} & 09:48 & 4 & Nottingham (Opteron) \\\\ \\hline {\\bf Translating Cluster Test} &  &  &  \\\\ \\hline {\\it Enzo (PPM)}&  78:50  & 1 & Columbia (Opteron) \\\\ {\\it Enzo (Zeus)} & 70:45 & 1 & UF Astro (Intel Xeon)\\\\ {\\it Flash} & 374:12 $\\dagger$ & 12 & Miranda (Opteron) \\\\ {\\it Hydra} & 39:23 & 1 & Nottingham (Opteron) \\\\ {\\it Gadget2} & 09:18 & 4 & Nottingham (Opteron) \\\\ \\hline \\end{tabular} \\label{table:times} \\end{table*}   While this paper attempts to cover the major features of astrophysical simulations, it does comprise only four tests. Further examples would extend this paper beyond its scope (and readability), but the differences between codes cannot be fully cataloged without further testing. This paper then, is designed as a starting point for a suite of tests to be developed from which codes can be quantitatively compared for the jobs they are intended for. To assist groups wishing to run these tests on their own code and to compare new updates, we are making these results and initial conditions available on the web at {\\tt http://www.astro.ufl.edu/codecomparison}."
    },
    "0808/0808.1305_arXiv.txt": {
        "abstract": "We report the search for low-mass \\xray\\ binaries in quiescence (qLMXBs) in the globular cluster (GC) NGC~6304 using \\xmm\\ observations. We present the spectral analysis leading to the identification of three candidate qLMXBs within the field of this GC, each consistent with the \\xray\\ spectral properties of previously identified qLMXBs in the field and in other GCs -- specifically, with a hydrogen atmosphere neutron star with radius between 5--20\\km.  One (\\sourcefour, with $\\rinfty =11.6\\ud{6.3}{4.6}$\\,(D/5.97\\kpc)\\km\\ and $\\kteff=122\\ud{31}{45} \\eV$ is located within one core radius ($r_c$) of the centre of NGC~6304. This candidate also presents a spectral power-law component contributing 49 per cent of the 0.5--10\\keV\\ flux.  A second one (\\sourcenine\\ with $\\rinfty=10.7\\ud{6.3}{3.1}$\\,(D/5.97\\,kpc)\\,km and $\\kteff=115\\ud{21}{16}$\\,eV) is found well outside the optical core (at $\\sim32\\,r_c$) but still within the tidal radius.  From spatial coincidence, we identify a bright 2MASS infrared counterpart which, at the distance of NGC~6304, seems to be a post-asymptotic giant branch star.  The third qLMXB (\\sourcefive\\ with $\\rinfty=23 \\ud{69}{10}$\\,(D/5.97\\kpc)\\km\\ and $\\kteff=70\\ud{28}{20} \\eV$) is a low signal-to-noise candidate for which we also identify from spatial coincidence a bright 2MASS infrared counterpart, with 99.916 per cent confidence.  Three qLMXBs from this GC is marginally consistent with that expected from the encounter rate of NGC~6304.  We also report a low signal-to-noise source with an unusually hard photon index ($\\alpha=-2.0\\ud{1.2}{2.2}$).  Finally, we present an updated catalogue of the \\xray\\ sources lying in the field of NGC~6304, and compare this with the previous catalogue compiled from \\rosat\\ observations. ",
        "introduction": "Transiently accreting low-mass \\xray\\ binaries in quiescence (qLMXBs) produce useful constraints on physical models of the interiors of neutron stars.  They were first discovered during the post-outburst periods of transient LMXBs Cen X-4 and Aql X-1, in which their faint emission ($L_{\\rm X}\\sim10^{32}-10^{33}\\cgslum$), interpreted as thermal black-body emission, had emission areas a factor of 100 smaller than implied by the projected area of a 10\\km\\ neutron star (NS).  Further observations of Cen X-4 and new observations of 4U 2129+47 \\citep{garcia94} and 4U 1608$-$522 \\citep{asai96b} confirmed these observational properties of the class of qLMXBs.  At that time, it was suggested that some low-level mass accretion on to the compact object was required to power the observed luminosity, although what set the low-mass accretion rate ($\\dot{M}$) was unclear \\citep{jvp87,verbunt94}.  An alternate explanation to low-$\\dot{M}$ accretion has become the dominant explanation for the quiescent emission of qLMXBs \\cite[BBR98 hereafter]{brown98}.  An up to date summary of the theoretical developments of deep crustal heating and the related observations is provided here.  In this different interpretation of the quiescent thermal emission, the luminosity is provided not by accretion on to the compact object, but by heat deposited in the neutron star crust by pressure-sensitive nuclear reactions.  These reactions \\citep[HZ08 hereafter]{haensel90,haensel03,gupta06,haensel08}, referred to as deep crustal heating (DCH), occur as matter is piled at the top of the neutron star surface, forcing the column below to greater density and electron Fermi energies.  For example, beginning with pure \\nuc{56}{Fe} at the top of the crust, the increasing electron Fermi energy induces electron capture at a density of $1.494 \\tee{9}\\cgsdensity$; because of the odd-even staggering of nuclear masses, a second electron capture follows to produce \\nuc{56}{Cr} and releases 0.041\\MeV\\ per accreted nucleon (HZ08).  At a density $1.114\\tee{10}\\cgsdensity$ the electron Fermi energy is again above threshold for capture on to \\nuc{56}{Cr}, and \\nuc{56}{Ti} is produced with an energy deposition of 0.036\\MeV\\ per accreted nucleon.  This process continues, resulting in a series of electron captures, neutron emission, and pycnonuclear reactions, through the depth of the crust, until the matter reaches the density of undifferentiated equilibrium nuclear matter, at which point it is no longer constituted of differentiable nuclei.  In doing so, this reaction chain deposits 1.9\\MeV\\ per nucleon, distributed throughout the crust (HZ08). Although the amount of heat deposited into the outer crust depends on the composition of nuclei produced by the fusing of accreted H and He \\citep{gupta06}, the total amount of heat is dominated by neutron emissions and pycnonuclear reactions in the inner crust, giving a total heat deposition that is relatively insensitive to the nuclear history of the accreted matter (HZ08).  In a steady-state, these reactions give rise to a time-average luminosity, which is directly proportional to the time-average mass accretion rate (BBR98): \\begin{equation} L = 9\\tee{32}\\,\\frac{\\langle \\dot{M} \\rangle}{10^{-11}\\unit{\\msun\\, \\peryear}}\\,\\frac{Q}{1.5\\unit{MeV/amu}}\\cgslum \\label{eq:dch} \\end{equation} \\noindent where $Q$ is the average heat deposited in the NS crust per accreted nucleon.  Of course, nuclear burning in the NS atmosphere will produce not pure iron, but a wide range of nuclei \\citep{schatz99, schatz01}; this has been shown \\citep{haensel03} to possibly change $Q$ to range between 1.1-1.5\\MeV\\ per accreted nucleon.  Moreover, the uncertain time-scale for the pycnonuclear burning of \\nuc{34}{Ne} \\citep{yakovlev06}, which arises in the Fe-chain, leaves uncertain whether that reaction is followed, or some other reaction takes its place on a shorter time-scale, at slightly greater density.  However, in the final analysis, the amount of heat $Q$ deposited in the crust is roughly equal to the average mass defect of nuclei at the top of the crust (near densities $\\approx\\ee{9}\\cgsdensity$) with undifferentiated equilibrium nuclear matter at the bottom of the crust, near nuclear saturation density $2.6\\tee{14}\\cgsdensity$ (HZ08).  Some qLMXBs exhibit significantly less photon luminosity than predicted by Eq.~(\\ref{eq:dch}). In several cases, such as KS~1731$-$260 \\citep{wijnands01b,rutledge02c,cackett06} and 1H~1905+000 \\citep{jonker07b} the long-term mass-transfer rate is not well constrained.  For the accreting millisecond pulsar SAX J1808.4$-$3657, however, predictions of $\\langle \\dot{M}\\rangle$ from the integrated fluence agree with predictions based on the orbital period and gravitational radiation losses \\citep{bildsten01}. For this system, the low quiescent luminosity is consistent with strong neutrino emission from the core, such as from nucleon or hyperon direct Urca processes \\citep{heinke07}.  The energy source of the quiescent thermal luminosity aside, the thermal spectrum can be explained (BBR98) as due to a realistic neutron star hydrogen atmosphere model \\citep{zavlin96,rajagopal96,heinke06}, rather than a black-body approximation imposed in previous works.  In the case of qLMXBs, the atmosphere has been accreted from the low-mass post-main sequence stellar companion (BBR98).  When accretion shuts off, metals will settle gravitationally on a time-scale of $\\sim$\\,seconds \\citep{bildsten92}, leaving a pure H-atmosphere.  The emergent photon spectra of H atmosphere NSs is physically different -- although parametrically similar -- to that of blackbody thermal spectra \\citep{zavlin96}, which dramatically changes the derived emission area radii from the $\\approxlt 1\\km$ derived from the black-body assumption, to radii which are consistent with the entire area of the neutron star \\citep{rutledge99}.  The DCH luminosity/H atmosphere spectral interpretation has been applied in the detection (and non-detections) of a large number of historically transient LMXBs, including: Cen X-4 \\citep{campana00,rutledge01}; Aql X-1 \\citep{rutledge01b} 4U~1608-522 \\citep{rutledge99}; 4U~2129+47 \\citep{rutledge00}; the transient in NGC~6440, SAX~J1748$-$2021 \\citep{intzand01}; X1732-304 in Terzan 1 \\citep{wijnands02}; XTE~J2123$-$058 \\citep{tomsick04}; EXO~1747$-$214 \\citep{tomsick05}; MXB~1659$-29$ \\citep{cackett06}; 1M~1716$-$315 \\citep{jonker07a}; 2S~1803$-$245 \\citep{cornelisse07}; 4U~1730$-$22 \\citep{tomsick07}; 1H~1905+000 \\citep{jonker07b}.  The spectral analysis of qLMXBs with the H atmosphere interpretation leads to the determination of \\rinfty, the emission radius (also called, the projected radius) of the neutron star.  It differs from the physical radius by a factor $g_{r}$, the gravitational redshift: \\begin{equation} \\rinfty = \\frac{\\rns}{g_{r}} = \\rns  \\left(1 - \\frac{2 G M_{\\rm NS}}{\\rns\\ c^{2}} \\right)^{-1/2} \\label{eq:rinf} \\end{equation} \\noindent The effective temperature is also affected by the gravitational redshift through the relation $T_{\\rm eff}^{\\infty}=g_{r}T_{\\rm eff}$.  A major uncertainty in accurate measurements of \\rinfty\\ remains the distance $D$ to the system. While high signal-toi-noise (S/N) \\xray\\ spectra are capable of determining $\\rinfty/D$ to statistical certainty, the 10--50 per cent systematic uncertainties in the distances to field sources directly contribute an additional 10--50 per cent uncertainty in \\rinfty.  The observational motivation for measuring \\rinfty\\ -- placing observational constraints on the dense matter equation of state -- requires measurements of \\rinfty\\ to $\\approx 5$ per cent accuracy, or better \\citep{lattimer04}.   The observational solution is to discover new qLMXBs, the distances to which are known with greater certainty than the present group of field sources.  Globular clusters (GCs) are the obvious place to search for these, as they host an over-abundance of LMXBs and their distances are precisely known or measurable.  Historically, it was suggested that \\xray\\ binary systems in GCs are formed by capture from the remnants of massive stars, i.e. neutron stars or black holes \\citep{clark75}. It was also noted that the ratio of \\xray\\ sources to the surrounding stellar density is two orders of magnitude larger for GCs than for the Milky Way.  For example, 12 out of $\\sim$ 100 known bright \\xray\\ sources were lying within GCs while the population of GC represents only $10^{-4}$ of the total mass of the Milky Way \\citep{hut92}.  The encounter rate is related to the physical properties of GCs, so that the high stellar density in their core creates a propitious environment for large encounter rates \\citep{verbunt02}.  Recently, a correlation was shown between the GC encounter rate and the number of hosted \\xray\\ sources \\citep{pooley03} or between the GC encounter rate and the number of hosted accreting NS binary systems \\citep{heinke03b}.  In addition, cluster metallicity can play a significant role in the number of LMXBs in GCs \\citep{kundu02,maccarone03,ivanova06,sivakoff07}.   To date, only a handful of qLMXBs in GCs capable of producing high S/N \\xray\\ spectra are known: in Omega Cen \\citep{rutledge02b,gendre03a}; X5 and X7 in 47~Tuc \\citep{heinke03,heinke06}; in NGC~6397 \\citep{grindlay01b}; in M13 \\citep{gendre03b}; in M28 \\citep{becker03}; A-1 in M30 \\citep{lugger07}; in M80 \\citep{heinke03c}; in M55 and NGC~3201 \\citep{webb06}. Others have been reported in the literature, but with lower S/N, or high values of \\nh, which makes them significantly less useful for the primary motivation for searching for qLMXBs: providing measurements of \\rinfty.  In each case, the qLMXB was identified by its \\xray\\ spectrum alone, consistent with a H-atmosphere neutron star with radius of $\\approx 10\\km$, without observation of an optical/IR counterpart, or evidence of an historical \\xray\\ outburst.  In two cases -- X5 in 47 Tuc \\citep{edmonds02} and Omega Cen \\citep{haggard04} -- the optical counterpart of the binary was identified following the spectral classification, lending support to the assertion that \\xray\\ spectral classification can identify qLMXBs.  Some qLMXBs also exhibit a power-law (PL) spectral component which dominates at photon energies above 2\\keV.  The first detection of such a component was made on the qLMXB Cen X-4 \\citep{asai96b,asai98}, where the power-law contributes about 40 per cent of the 0.5--10\\keV\\ flux.  Observations of 47 Tuc (qLMXBs X5 and X7) also showed weak power-law components (with photon index $\\alpha=2.6{\\rm -}3$) \\citep{grindlay01}.  However, analysis of more sensitive observations concluded that PL components are not required \\citep{heinke03}.  A more recently identified candidate qLMXB in 47~Tuc, has shown a power-law component that represents $65\\ud{28}{22}$ per cent of the 0.5--10\\keV\\ unabsorbed flux \\citep{heinke05}.  NGC~6440 have possible qLMXBs with significant PL components, but the large absorption in the direction of this GC prevents a justified classification of the \\xray\\ sources due to the low photon flux where the thermal spectrum is dominant \\citep{intzand01}.  In general, qLMXBs in GCs display no or little power-law components while qLMXBs in the field more frequently tend to have a significant power-law component.  A proposed interpretation relates the power-law emission as the residual of a recent accretion episode \\citep{grindlay01}. An alternate explanation evokes it to shock emissions via the emergence of a magnetic field \\citep{campana00b}.  However, neither of these two hypothesis are yet supported by observational evidence.  A program of short-integration \\xmm\\ observations, using the pn camera, have been undertaken to survey GCs for qLMXB candidates, to increase the number of known qLMXBs in GCs.  This paper reports the discovery of three candidate qLMXBs in NGC~6304 from spectral classification, \\sourcefour, \\sourcenine\\ and the low signal-to noise source, \\sourcefive.  These were found in short \\xmm/pn exposures targeted specifically for the discovery of examples of this class of objects.  The organisation of this paper is as follows.  In \\S\\ref{sec:obs} and \\S\\ref{sec:astrometry}, the data reduction, the observations and the astrometric analysis are described.  In \\S\\ref{sec:spectra}, the spectral analysis of the detected sources are presented with detailed explanations for interesting sources, including a hard source with the photon index $\\alpha= -2.0\\ud{1.2}{2.2}$.  In \\S\\ref{sec:candidates}, the observational properties of the three candidate qLMXBs are described. \\S\\ref{sec:discuss} covers a comparison with a previous observation of NGC~6304 with \\rosat\\ (\\S\\ref{sec:rosat}) and a comparison with the expected number of qNS in GCs (\\S\\ref{sec:number}).  Finally, \\S\\ref{sec:conclude} summarizes the results of this observation and concludes.  \\subsection{NGC~6304} \\label{ngc6304} The globular cluster NGC~6304 is located at the position R.A.=17h14m32.1s and ${\\rm Dec.}=-29\\deg 27\\arcmin 44.0\\arcsec$, (J2000) approximately 2.2\\kpc\\ from the galactic centre.  Its integrated visual magnitude is $V=8.22$ and it has a color excess $E\\left(B-V\\right)=0.53$.  NGC~6304 is moderately compact, with the core radius $r_c=0.21\\arcmin$, the half mass radius $r_{\\rm HM}= 1.41\\arcmin$ and the tidal radius $r_t=13.25\\arcmin$.  It is a non core-collapse GC with a concentration $c=\\log\\left(r_{t}/ r_{c}\\right)=1.80$ (assuming a King's profile) and with a central luminosity density $\\log\\left(\\rho_{0}\\right)= 4.41\\unit{\\lsun\\,pc^{-3}}$.  Finally, its metallicity $\\left[{\\rm m/H}\\right]$ is -0.56, implying Z = 0.00488, where $Z_{\\odot} = 0.0177$ was used \\citep{montalban04}.  All properties come from the updated catalog of GCs \\citep{harris96} or from a more recent work \\citep{valenti07}.  A more recent measurement of the cluster distance modulus, $\\left(m-M\\right)_{0}= 13.88\\pm0.03$, provides an heliocentric distance of $5.97 \\ppm{0.08} \\kpc$ \\citep{valenti05}. ",
        "conclusions": "\\label{sec:conclude}  The globular cluster NGC~6304 hosts three candidate qLMXBs.  The neutron star radii and temperatures obtained for the three candidates are consistent with typical values for neutron stars and in accordance with radii of other known field and GCs qLMXBs.  No variability is measured from any of the candidates over the time scale of the observation, with very weak limits; the large X-ray variability in the \\xmm/pn background precludes a more detailed variability analysis.  A comparison with \\rosat\\ observations 14 yr prior (R94), finds the fluxes of the candidates are consistent with those of the earlier observation (R94).  The absence of a detailed error analysis in the fluxes from the published \\rosat\\ analysis precludes placing statistical limit on the amount of variability; however, as the number of detected counts in the \\rosat\\ observations was small ($<$20), the magnitude of intensity variability is limited to be less than a factor of $\\sim$few on the 14-yr time-scale.  \\sourcefour\\ is located in the core (0.79\\unit{r_c}) of the GC.  Its spectrum is acceptably fitted with a NS Hydrogen atmosphere model with $\\kteff=122 \\ud{31}{45}\\eV$ and $\\rinfty=11.6\\ud{6.3}{4.6}\\km$, combined with a power-law component of photon index $\\alpha= 1.2\\ud{0.7}{0.8}$.  The power-law component dominates the spectrum at photon energies above 2\\keV, and represents 49 per cent of the observed flux (0.5--10 keV).  This contribution, although its source is not perfectly understood, is similar to that found in some other qLMXBs, and is comparable in the fraction of the flux for which it accounts to the field qLMXB Cen~X-4 \\citep{asai96b,rutledge01,menou01,campana04}. While the high stellar density in the core of NGC~6304 precludes identification of a 2MASS counterpart for \\sourcefour, a visual overlap with a 2MASS image of the GC shows a possible association with an uncatalogued -- and likely unresolved -- star.  The association will require further investigation with higher resolution optical/IR imaging.  \\sourcefive, a second candidate, has a low signal-to-noise and correspondingly larger error bars for the X-ray spectral parameters used to identify it as a qLMXB candidate: the NS H-atmosphere temperature is $70 \\ud{28}{20}\\eV$, and $\\rinfty=23\\ud{69}{10}\\km$.  A faint 2MASS counterpart ($m_{J}=13.063$) is identified (with 99.916 per cent confidence), and the \\xray\\ emission is unlikely to emerge from the giant star itself.  \\sourcenine, a third candidate, is a bright \\xray\\ source consistent with a neutron star H atmosphere spectral model at the distance of the GC.  Its measured properties are $\\kteff=115 \\ud{21}{16}\\eV$ and $\\rinfty = 10.7\\ud{6.3}{3.1}\\km$.  The 2MASS IR counterpart (with 99.988 per cent confidence) appears to belong to the post-asymptotic giant branch; this is the first qLMXB with this type of companion. Other explanations for the source are investigated, concluding that a coronally active field star is not a possible explanation, due to the unusual optical/IR colors and unusual X-ray spectrum.  Optical/IR spectroscopy of \\TwoMassnine\\ can definitively determine whether the star is a foreground main sequence star, or an evolved post-AGB star at the distance of NGC~6304.  A faint unidentified \\xray\\ source showing an unusual excess of high energy photons resulting in a photon index of $\\alpha = -2.0\\ud{1.2}{2.2}$.  The nature of this source is unclear. There is no evidence the X-ray source was variable during the observation.  No off-band counterparts in the 2MASS or DSS catalogues are found within $3\\sigma$ of the X-ray position.  Deep observations in the off-bands, as well as repeated detection and study in the X-ray, can inform interpretation of this source.  The qLMXB candidates in NGC~6304 are astronomically interesting. First, the number of candidates is consistent with predictions from the calculated encounter rate, supporting previous observational analyses which characterised the qLMXB rate vs. encounter rate in GCs \\citep{pooley03,heinke03b,gendre03b}.  This indicates that the qLMXB rate vs. encounter rate is not merely a parameterisation resulting from observations, but can predict the number of qLMXBs in unobserved GCs.  NGC~6304 contains one of the most luminous known qLMXBs.  \\sourcefour, with a intrinsic luminosity $L_{X}=\\ee{33}\\cgslum$ (0.5--10\\keV), is the fourth most luminous GC qLMXB after 47 Tuc X7 ($L_{X}=1.5\\tee{33}\\cgslum$ (0.5--10\\keV)), 47 Tuc X5 ($L_{X}=1.4\\tee{33}\\cgslum$ (0.5--2.5\\keV)) and the qLMXB in M28 ($L_{X}=1.2\\tee{33}\\cgslum$ (0.5--10\\keV)).  These represent the upper limit of the expected luminosity for LMXB in quiescence.  In addition, \\sourcefour, with its high flux ($F_{X}=2.3\\tee{-13}\\cgsflux$, after 47 Tuc X7, 47 Tuc X5 and the qLMXB in M28 with fluxes 5.3, 4.3 and $3.3\\ud{1.9}{1.1}$ $\\tee{-13}\\cgsflux$, respectively - see Table~\\ref{tab:NS}), is a good candidate to constrain the equation of state of dense matter.  NGC~6304 hosts two candidate qLMXBs lying at large distance from the hosting core, but still within the tidal radius; \\sourcenine\\ and \\sourcefive\\ are located at $\\sim32r_{c}$ and $\\sim15r_{c}$, respectively.  Previously, all GC qLMXBs were located within $<7 r_c$ of the optical centre of their GC (see Table~\\ref{tab:Rcore}), the most distant belonging to U24 in NGC~6397, located $6.8r_{c}$ from the centre; noting, however, that NGC~6397 is a core collapse GC with $r_{c}=0.05\\arcmin$.  To sum up, the qLMXB candidates in NGC~6304 are of much greater distance from the GC centre than has been observed previously, and may require revisiting theory of qLMXB formation and binary evolution in GC.  Deeper X-ray exposures are required to confirm the classification of the candidates qLMXB found in NGC~6304.  Specifically, it is necessary to constrain \\rinfty\\ to an uncertainty of few percent (as already performed for confirmed qLMXBs in $\\omega$ Cen and 47 Tuc, see Table \\ref{tab:NS}) rather than few tens of percent here.  For \\sourcefour, better angular resolution will confirm the isolated nature of the candidate qLMXB, or may resolve multiple \\xray\\ sources in the core. New deep observations of \\sourcefive\\ and \\sourcenine\\ will produce a localisation unbiased by bad detector columns.  At optical wavelengths, observations of the aforementioned candidates can determine unequivocally the evolutionary state of the counterparts (a giant for \\TwoMassfive, and a post-AGB star for \\TwoMassnine), as well as measure their binary orbital periods through ellipsoidal variations.  The spatial isolation and brightness of the IR counterparts implies this work can be done easily using ground based, moderate sized telescopes.  Optical spectroscopic observations can confirm that the counterparts have radial velocities consistent with that of NGC~6304.  Finally, high spatial resolution observations of the GC core in the optical/IR bands can pinpoint a possible counterpart for \\sourcefour\\ and confirm the qLMXB nature of this \\xray\\ source."
    },
    "0808/0808.1902_arXiv.txt": {
        "abstract": "Extrasolar super-Earths (1-10 M$_{\\earth}$) are likely to exist with a wide range of atmospheres. Some super-Earths may be able to retain massive hydrogen-rich atmospheres.  Others might never accumulate hydrogen or experience significant escape of lightweight elements, resulting in atmospheres more like those of the terrestrial planets in our Solar System.  We examine how an observer could differentiate between hydrogen-rich and hydrogen-poor atmospheres by modeling super-Earth emission and transmission spectra, and we find that discrimination is possible by observing the transmission spectrum alone.  An Earth-like atmosphere, composed of mostly heavy elements and molecules, will have a very weak transmission signal due to its small atmospheric scale height (since the scale height is inversely proportional to molecular weight). On the other hand, a large hydrogen-rich atmosphere reveals a relatively large transmission signal.  The super Earth emission spectrum can additionally contrain the atmospheric composition and temperature structure. Super-Earths with massive hydrogen atmospheres will reveal strong spectral features due to water, whereas those that have lost most of their hydrogen (and have no liquid ocean) will be marked by CO$_2$ features and a lack of H$_2$O. We apply our study specifically to the low-mass planet orbiting an M star, Gl 581c ($M sin i$ =  5 M$_{\\earth}$), although our conclusions are relevant for super-Earths in general.  The ability to distinguish hydrogen-rich atmospheres might be essential for interpreting mass and radius observations of planets in the transition between rocky super-Earths and Neptune-like planets. ",
        "introduction": "The search for extrasolar planets has recently resulted in the discovery of a new class of planets, super-Earths, with masses between about 1 and 10 Earth masses (M$_{\\earth}$) \\citep{riv05, bea06, udr07}. Since no such planets exist in our Solar System, the bulk composition of these planets and of their atmospheres remains largely unknown, although some theoretical work on the subject has been presented. In contrast to rocky planets in our Solar System, super-Earths are predicted to have large surface gravities, - upwards of 25 m/s$^2$ for a 5 M$_{earth}$ planet - \\citep{val06, val07, for07, sea07, sot07} and some super-Earths may therefore be able to retain massive H-rich atmospheres.  Others will bear a closer resemblence to Earth, with atmospheres depleted in hydrogen and composed of predominantly heavier molecules. A third interesting case is that some super-Earths may lie in between these two extremes, with moderate levels of atmospheric hydrogen due to either incomplete escape of hydrogen over the planet's lifetime and/or from outgassing of a significant secondary H$_2$ atmosphere. The question arises as to whether these cases will differ in their gross observable properties, and how. In the interest of discovering habitable planets and learning more about atmospheric evolution on super-Earths, it will be necessary to find a method for discriminating between the possible types of atmospheres.  This issue has led us to model the observational signature for the various possible atmospheres on the low-mass planet, Gl 581c \\citep{udr07}.  The bulk composition of planets in the mass range between Earth and Neptune carries important clues to understanding planet formation.  Even the simple statistical distribution of very water-rich versus dry super-Earths in an observed sample would be sufficient \\citep{val07}. Unfortunately, with planet mass and radius alone the internal structure models have degeneracies, which become almost untractable when a hidrogen-rich envelope covers the planet \\citep{ada07}. Our study in this paper points to a potential solution to breaking this degeneracy, by distinguishing a hydrogen-rich atmosphere spectroscopically.  A number of observations of transiting hot Jupiters have already allowed for constraints to be placed on the nature of their atmospheres.  IR observations of the secondary eclipse (when the planet passes behind its host star) have resulted in measurements of brightness temperatures \\citep{dem05, cha05, dem06, har07, dem07, demo07, knu08} as well as emission spectra \\citep{gri07, ric07, swa07, cha08} for several hot Jupiters.  Optical transmission spectra for HD 209458b have led to the discovery of several species in the planet's atmosphere including sodium \\citep{cha02}, hydrogen \\citep{vid03, ehr08}, and tentative detections of oxygen and carbon \\citep{vid04}. Recent observations of the IR transmission spectrum of HD 189733b have led to the first discoveries of water \\citep{tin07, bea08} and methane \\citep{swa08} in the atmosphere of an extrasolar planet.  Additionally, attempted observations of the optical secondary eclipse of HD 209458b using the {\\it MOST} space telescope have allowed for strong limits to be placed on the planet's optical albedo \\citep{row06, row07}.  Since super-Earths are smaller than hot Jupiters, these types of measurements will prove more challenging for this new class of planets \\citep{sea08}.  However, it will most likely be these same types of observations that will eventually aid in characterizing the  atmospheres of low-mass planets.   The discovery of the low-mass planet Gl 581c ($M sin i =  5 M_{\\earth}$) has generated a large amount of interest in the literature, and a number of authors have already explored the question of habitability for this planet after an initial claim by \\citet{udr07} that Gl 581c may reside inside its star's habitable zone.  It now appears that the planet is probably too hot to host liquid water unless its surface is mostly obscured by highly reflective clouds \\citep{sel07}, or the planet is tidally locked, with the night side receiving sufficient heat circulation from the day side of the planet \\citep{chy07}.  To further confirm this conclusion, \\citet{von07, von07b} have examined the possibility of sustained photosynthetic life on Gl 581c and have ruled this out on the grounds that the planet cannot keep sufficient CO$_2$ in its atmosphere for photosynthesis to occur, while at the same time retaining a surface temperature cool enough for liquid water.  However, in terms of climate stability,  \\citet{beu07} have modeled the dynamic evolution of the Gl 581 system.  They find that the orbit (and therefore the climate) of Gl 581c should be secularly stable over 10$^8$ yrs, leaving open a possibility that life of a more exotic variety could still develop on this planet.  In this paper, we address a different question relating to the nature of Gl 581c's atmosphere.  What observations are needed to constrain the atmospheric composition and overall hydrogen content for Gl 581c or similar super-Earths? To this end, we investigate several scenarios for the composition of the super-Earth atmosphere ranging from hydrogen-rich to hydrogen-poor. The cases that we explore are (i) a massive hydrogen-rich (reducing) atmosphere obtained through accretion of nebular gas, (ii) a mildly reducing atmosphere that has lost most but not all of its hydrogen due to atmospheric escape processes, and (iii) an oxidizing CO$_2$-rich atmosphere similar in composition to that of Venus that contains essentially no hydrogen.  From the modeled spectra, we determine how to observationally differentiate between the various scenarios.  We describe our model atmosphere in \\S~\\ref{model} and our three scenarios for atmospheric composition in \\S~\\ref{comp}.  In \\S~\\ref{results} we reveal the modeled emission and transmission spectra for Gl 581c, and we determine how hydrogen content can be established through observations.  Finally, we summarize our results and discuss their implications in \\S~\\ref{discussion}. ",
        "conclusions": "}  We have determined a method for differentiating between Super-Earths with large hydrogen-rich atmospheres and those with atmospheres composed mostly of heavier molecules.  Since the relative depth of spectral features in a planet's transmission spectrum results in a direct determination of the atmospheric scale height, the mean molecular weight of that atmosphere can then be calculated by a simple relation (see Eq.~\\ref{H_eq}).  For the case of Gl 581c, a hydrogen-rich atmosphere with its large scale height reveals spectral lines in transmission at a level of 10$^{-4}$.  On the other hand, the transmission features for a hydrogen-poor atmosphere will be more than an order of magnitude smaller at a level of less than 10$^{-5}$.  The presence of clouds will not alter this conclusion unless they are unrealistically high in the atmosphere - at a height approaching 10 scale heights above the planet's surface, where the atmospheric density has dropped off to a point that is is difficult to imagine the presence of an optically thick cloud deck. We therefore emphasize that transmission spectroscopy should be a robust method for determining the atmospheric scale height.   The emission spectrum should also be generally helpful in constraining atmospheric hydrogen content.  However interpretation can be more complicated than for transmission spectra due to the dependence of the emission spectrum on the atmospheric temperature gradient.  For a hydrogen-dominated super-Earth, the emission spectrum should be marked mostly by water features. As we then transition in the cases that we have presented, from H-rich to H-poor, we find that the planetary emission starts showing weaker H$_2$O features, while absorption due to CO$_2$ increases.  Additionally, the presence of methane or ammonia features for hydrogen-containing atmospheres could be an interesting indication of the photochemistry at work, since these molecules are thought to be destroyed by photodissociation over the lifetime of the planet. Even a weak detection of these features could therefore be helpful in constraining models for atmospheric evolution.   The exact mass of Gl 581c's atmosphere, which has been chosen somewhat arbitrarily in this paper, has little effect on our spectra presented above.  This is due to the simple fact that regions beyond an optical depth of unity are essentially hidden from observers.  Therefore, as long as the atmosphere becomes optically thick at some height above the planetary suface, a more massive atmosphere would  reveal the same emission and transmission spectra as a less massive one barring any secondary effects such as differing cloud structures. Pressure and temperature conditions near the planetary surface, which would be indicative of the atmospheric mass cannot be observed via spectroscopy for an optically thick atmosphere.  For this reason it will be difficult to constrain atmospheric mass with spectral observations.  Another negative consequence is that the question of habitability will remain unanswered as long as the planet's surface conditions are unknown and unobservable.  Mass-radius measurements may provide the first constraints on atmospheric mass for super-Earths, however there are large degeneracies in the models that will be difficult to overcome \\citep{ada07}.   Our simple model for the super-Earth atmosphere has allowed us to understand some of the important effects of hydrogen content on transmission and emission spectra.  However, future more detailed models will be needed to address some of the issues that we have simplified here.  A T-P profile that is coupled with a photochemical model is needed to correctly predict some of the feedback mechanisms that occur as photochemistry alters the atmospheric composition.  Surface processes also much be considered as they can further affect atmospheric composition through cycles (such as the carbonate-silicate cycle), outgassing, or interactions with a liquid ocean. Additional questions to be addressed by a general super-Earth model are those that plague all planetary atmosphere models such as the presence of clouds and a proper treatment of condensation.   Currently, no transiting super-Earths have been detected.  However, CoRoT \\citep{bag03, bar05} and Kepler \\citep{bor04, bas05} missions are both designed with the capability of detecting transiting super-Earths orbiting Sun-like stars. Depending on occurrence rates, a large number of such planets could be discovered in the next 5 years by both of these space-based missions. Additionally, ground-based transit surveys such as MEarth \\citep{nut07}, which will specifically target M-dwarfs, as well as radial velocity surveys both have the capabilities to detect planets in the super-Earth range.  Most interestingly, the combination of these searches should allow for super-Earths to be found in a wide range of orbits and around a variety of types of stars.   Spitzer is currently the only telescope that has demonstrated the capability of observing exoplanet emission and transmission spectra in the mid-IR. Unfortunately, it is unlikely that Spitzer will be able to make the same types of observations for super-Earths, with the possible exception of a warm planet with a hydrogen-rich atmosphere -- see Figure ~\\ref{transmission}. Otherwise, measuements would need to be taken at a level of several to tens of ppm.  The James Webb Space Telescope (JWST) \\citep{gar06} should fare better, since it has been designed to make these types of measurements at the precision needed.  With its expected instrument complement that will allow for narrow band photometry and low to medium resolution IR spectroscopy, JWST should allow for follow-up observations to characterize the atmospheres of transiting Super-Earths.  Valenti et al.~(in prep.)~have simulated JWST NIRSpec spectra for transiting Earth-like planets orbiting M-stars, using atmospheric models from \\citet{ehr06}. They determine that 10 transits (for a 1.4 hr transit duration) and 28 hours of observing time are needed to detect water absorption features in transmission at a 4-$\\sigma$ detection threshold. Scaling these values to the hydrogen-rich case for Gl 581c results in a 20-hour integration time and 6 transits needed to detect its transmission features with their expected depth of 10$^{-4}$.  Others have suggested that JWST can do better by up to a factor of two (Swain 2008, submitted).  If this is the case, then even less observing time will be needed.  These types of observations should be feasible with proper telescope scheduling and should allow for the full richness of super-Earth atmospheres of different sizes, compositions, and orbits to be explored."
    },
    "0808/0808.3136_arXiv.txt": {
        "abstract": "In March 2009, NASA will launch the Kepler satellite---a mission designed to discover habitable Earth-like planets around distant Sun-like stars. The method that Kepler will use to detect distant worlds will only reveal the size of the planet relative to the size of the host star, so part of the mission is devoted to characterizing other suns using asteroseismology. In this proceedings, I give a broad overview of the Kepler mission and the data that it will produce, with a special emphasis on how it could improve our understanding of solar and stellar dynamos. I conclude with an update on the development of a stellar modeling pipeline for interpreting asteroseismic observations. ",
        "introduction": "In the region of Canada that is now Quebec and Ontario, a Native American tribe known as the Algonquin developed a detailed myth about the annual path of the Big Dipper around the north celestial pole, which they chanted in their ``Song of the Stars'':  \\begin{quotation} \\noindent We are the stars which sing, \\\\ We sing with our light;  \\\\ We are the birds of fire, \\\\ We fly over the sky. \\end{quotation}  \\noindent If we think of asteroseismology as the study of ``stars which sing with their light'', then NASA's Kepler mission is poised to record the largest stellar choir ever assembled. In this proceedings, I will give a brief overview of the Kepler mission, followed by a review of the specific types of data the satellite is likely to produce which have the potential to improve our understanding of solar and stellar dynamos. I will conclude with an update on the work I have been doing to prepare for the massive wave of data soon expected from Kepler. ",
        "conclusions": "The Kepler mission promises to revolutionize the quality and quantity of asteroseismic data available for solar-type stars, and it will significantly contribute to research topics related to the solar-stellar connection. The mission needs asteroseismology to determine the absolute sizes of any potentially habitable Earth-like planets that might be discovered. It will also yield a wide variety of data that will be useful for calibrating dynamo models, sampling many different sets of physical conditions and evolutionary phases to complement the well-studied solar example. A uniform analysis of the asteroseismic data will help minimize systematic errors between the parameter values obtained for different stars, and such an analysis will be facilitated by our TeraGrid-based community modeling tool. The ``stars which sing with their light'' have many secrets to tell us, and the Kepler satellite will allow us to listen as never before."
    },
    "0808/0808.0819_arXiv.txt": {
        "abstract": "In this letter, we propose a model of inflationary cosmology with a bounce preceded and study its primordial curvature perturbations. Our model gives rise to a primordial power spectrum with a feature of oscillation on large scales compared with the nearly scale-invariant spectrum generated by the traditional slow rolling inflation model. We will show this effect changes the Cosmic Microwave Background (CMB) temperature power spectrum and the Large Scale Structure (LSS) matter power spectrum. And further with a detailed simulation we will point out this signal is detectable to the forthcoming observations, such as PLANCK and LAMOST. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.2441_arXiv.txt": {
        "abstract": "We present photometric analysis and follow-up spectroscopy for a population of extremely red stellar objects extracted from the point-source catalogue of the INT Photometric H$\\alpha$ Survey (IPHAS) of the northern galactic plane. The vast majority of these objects have no previous identification. Analysis of optical, near- and mid-infrared photometry reveals that they are mostly highly-reddened asymptotic giant branch stars, with significant levels of circumstellar material. We show that the distribution of these objects traces galactic extinction, their highly reddened colours being a product of both interstellar and circumstellar reddening. This is the first time that such a large sample of evolved low-mass stars has been detected in the visual and allows optical counterparts to be associated with sources from recent infrared surveys.  Follow-up spectroscopy on some of the most interesting objects in the sample has found significant numbers of S-type stars which can be clearly separated from oxygen-rich objects in the IPHAS colour-colour diagram. We show that this is due to the positions of different molecular bands relative to the narrow-band H$\\alpha$ filter used for IPHAS observations. The IPHAS $(r' - $H$\\alpha$) colour offers a valuable diagnostic for identifying S-type stars. A selection method for identifying S-type stars in the galactic plane is briefly discussed and we estimate that over a thousand new objects of this type may be discovered, potentially doubling the number of known objects in this short but important evolutionary phase. ",
        "introduction": "Evolved low and intermediate mass stars ($0.8 < M_{\\odot} < 8.0$) in the asymptotic giant branch (AGB) stage of evolution are some of the most luminous stars in the Galaxy and represent one of the final evolutionary stages of all low and intermediate mass stars. Observations of these evolved stars show various degrees of photospheric abundance enhancements due to carbon and s-process elements being brought to the surface by the third dredge-up \\citep{iben83}. In particular this has led to two specific subtypes of AGB stars known as S-type stars and carbon stars. S-type stars are thought to be the transition objects between the oxygen-rich M giants (C/O$ < 1$) and the carbon stars (C/O$ > 1$) as carbon is dredged up to the surface \\citep{keen54}. The relative rarity of S-type stars has been attributed to the short transition time from oxygen-rich to carbon-rich surface chemistry, though there is evidence that this may not be the complete picture \\citep[e.g.][]{chan91}.  Mass-loss during the AGB phase can replenish the interstellar medium (ISM) with this processed material and rates of mass-loss are observed to reach up to 10$^{-5}$~M$_{\\odot}$yr$^{-1}$ at the tip of the AGB \\citep[e.g.][]{knap85,zijl92}. The cool extended atmospheres of AGB stars are also one of the principle sites of dust formation in the Galaxy. The circumstellar material around such objects contributes to their reddening and, by the time they reach the tip of the AGB, they may become completely obscured at optical wavelengths. These stars are responsible for much of the processed material returned to the ISM and in doing so drive the evolution of the Galaxy \\citep{buss99}.  The INT Photometric H$\\alpha$ Survey of the northern galactic plane \\citep[IPHAS,][]{drew05} is an imaging survey using the Wide Field Camera (WFC) on the Isaac Newton Telescope (INT). The imaging in broad-band Sloan $r'$ and $i'$ filters, accompanying narrow-band H$\\alpha$ filter ($\\lambda_c = 6568$~\\AA, FWHM~=~95~\\AA) observations, reaches to at least $r' = 20$. One of the survey's main aims is to identify large numbers of objects in important short-lived phases of stellar evolution \\citep[see e.g.][]{with08,corr08}. The survey covers the entire northern galactic plane in the Galactic latitude range $-5^{\\circ} \\le b \\le +5^{\\circ}$, a total sky area of 1800 square degrees. More than 90\\% of the survey area has already been observed and an initial release of data has been made \\citep[see][]{gonz08}. The filters employed by IPHAS naturally allow it to highlight red and reddened populations in the galactic plane and it has uncovered large numbers of extremely red stellar objects (ERSOs), with ($r' - i'$) colour extending to almost $\\sim$6, the majority of which have no previous identification. This work comprises a study of these objects to determine their typical properties and assess the cause of their highly reddened colours.  In Section~2 we present the IPHAS observations and select a population of extremely reddened stellar objects from the point source catalogue (PSC). In Section~3 we analyse the photometric properties of these sources to determine their typical properties and in Section~4 we discuss the small fraction of objects with previous identifications. In Section~5 we investigate the galactic distribution of the ERSO sample and compare with galactic reddening maps. In Section~6 we cross-correlate the IPHAS photometry of these objects with photometry from several recent infrared surveys, and assess the fraction of objects with infrared excesses. In Section 7 we present near-IR spectroscopy of some of the most extreme sources in the colour-colour diagram, including the discovery of several S-type stars in a distinctive region of the IPHAS colour-colour plane. Finally, in Section 8 we outline a method for separating these interesting objects from the complete IPHAS catalogue. ",
        "conclusions": "Over 25,000 sources were selected from the IPHAS point source catalogue based on their extremely red colours, the majority of which have no previous identifications. These extremely red stellar objects (ERSOs) are predominantly late-type giant stars. Initial follow-up spectroscopy has revealed the molecular bands indicative of cool photospheres. The distribution of these objects in the galactic plane shows strong evidence for tracing galactic extinction, indicating that their highly reddened colours are mainly due to this cause. As would be expected for late-type giant stars, a magnitude-limited study of the infrared colour excesses of these sources reveals that the majority have circumstellar material which must also be contributing to their extreme reddening.  Follow-up spectroscopy reveals that chemically diverse objects such as S-type stars will fall below O-rich evolved stars in the IPHAS colour-colour plane because of the positions of their molecular bands relative to the narrow-band H$\\alpha$ filter. A method to identify these objects from the IPHAS point source catalogue is proposed and a simple estimate shows that over a thousand new S-type stars in the northern galactic plane could be discovered this way, doubling the currently known number of objects. The extension to the richer southern galactic plane using data from the forthcoming VPHAS+ survey would also be prosperous in this search.  In a further paper, we will present WHT LIRIS near-IR spectra of a large sample of ERSOs, in addition to optical spectra of ERSOs obtained using the MMT HectoSpec multi-object spectrograph. These spectra will be used to investigate the relationship between evolved star surface chemistry and photometric colours such as $(r' - $H$\\alpha$)."
    },
    "0808/0808.1959_arXiv.txt": {
        "abstract": "We have developed a macroscopic description of coherent electro-magnetic radiation from air showers initiated by ultra-high energy cosmic rays in the presence of the geo-magnetic field. This description offers a simple and direct insight in the relation between the properties of the air shower and the time-structure of the radio pulse. As we show the structure of the pulse is a direct reflection of the important length scales in the shower. ",
        "introduction": "In recent years the interest in the use of radio detection for cosmic ray air showers is increasing with the promising results obtained from recent LOPES~\\cite{Fal05,Ape06} and CODALEMA~\\cite{Ard06} experiments, which triggered plans to install an extensive array of radio detectors at the Pierre Auger Observatory~\\cite{Ber07}. For these reasons there is a strong interest to understand the link between the properties of the extensive air shower (EAS) and the time structure of the emitted pulse.  Already in the earliest works on radio emission from air showers~\\cite{Jel65,Por65,Kah66,All71}, the importance of coherent emission was stressed. Two mechanisms, Cherenkov radiation and geo-magnetic radiation were investigated. In more recent work~\\cite{Fal03,Sup03,Hue05}, the picture of coherent synchrotron radiation from secondary electrons and positrons in the Earth's magnetic field was proposed and extensive results on geo-synchrotron emission are given in~\\cite{Hue07}.  The emission of an electromagnetic pulse by the charges in an EAS is described by Classical Electrodynamics~\\cite{Jac-CE}. In this sense it is an easy matter. The complication arises is the designation of all possible charge and current distributions that drive the electromagnetic emission process. For this models are needed such as the collective model described in \\secref{MGMR}. This macroscopic description of geomagnetic radiation~\\cite{Sch08,Wer08} (MGMR) is in the basic idea very similar to that used in Ref.~\\cite{Kah66}. However, independent of any particular model, some of the general features of the emitted pulse can be calculated based the length scales involved in an EAS as is discussed in \\secref{L-scales}. For simplicity of the discussion we will limit ourselves here to an ideal case of a vertical airshower (along the $z$-axis) while the Earth-magnetic field is taken parallel to the Earth (along the $y$-axis). The case for more realistic geometry can be found in the literature~\\cite{Wer08}. ",
        "conclusions": "We have shown that the macroscopic model for radio emission offers a very clear and simple picture for the emitted radio pulse. Different length scales determine the structure of the radio pulse, depending on the distance from the shower core. Close to the core the pulse length is determined by the pancake thickness while at larger distances the shower profile is reflected in the pulse structure. From the radio pulse at various distances from the shower core thus both the pancake thickness as well as the height of the shower maximum can be determined. This offers interesting prospects to determine the cosmic-ray composition.  The present discussion has been limited to the main contribution to the emitted signal coming from the electrical current induced by the Earth magnetic field. In Ref.~\\cite{Wer08} also the contribution are investigated which are due to induced electric dipole moments and charge excess for realistic EAS simulations. These are shown to give correction of the order of 10\\%, leaving the basic conclusions unchanged. Also the effect of more realistic geometries and a finite index of refraction of air have been studied."
    },
    "0808/0808.1168_arXiv.txt": {
        "abstract": " ",
        "introduction": "Inflation model \\cite{Guth:1980zm} provides an elegant mechanism to solve many puzzles in the hot big bang model due to the quasi-exponential expansion of the universe. The quantum fluctuation generated during inflation naturally seeds the small anisotropies in the cosmic microwave background radiation and the formation of the large-scale structure. In general many scalar fields are present in the early universe. The dynamics of inflation is governed by inflaton(s), but the primordial curvature perturbation can be generated by inflaton(s) \\cite{Guth:1982ec} or curvaton(s) \\cite{Linde:1996gt}, or both of them. However the feature of the primordial power spectrum caused by inflaton is similar to curvaton.    The CMB non-Gaussianity \\cite{Komatsu:2001rj} opens a new window to probe the physics of the early universe. A well-discussed ansatz of non-Gaussianity has the local shape. This kind of non-Gaussianity can be characterized by some non-linearity parameters: \\begin{equation} \\zeta({\\bf x})=\\zeta_g({\\bf x})+\\frac{3}{5}f_{NL}\\zeta_g^2({\\bf x})+\\frac{9}{25}g_{NL}\\zeta_g^3({\\bf x})+\\cdots~, \\end{equation} where $\\zeta$ is the curvature perturbation in the uniform density slice, $\\zeta_g$ denotes the Gaussian part of $\\zeta$. The non-linearity parameters $f_{NL}$ and $g_{NL}$ parameterizes the non-Gaussianity from the irreducible 3-point and 4-point correlation functions respectively.    In the case of single field inflation model $f_{NL}^{local}\\sim {\\cal O}(n_s-1)$ \\cite{Maldacena:2002vr}, which is constrained by WMAP ($n_s=0.960_{-0.013}^{+0.014}$) \\cite{Komatsu:2008hk} to be much less than unity. In some general inflation models \\cite{Chen:2006nt} a large non-Gaussianity is possibly obtained, but the shape of non-Gaussianity is different from the local shape. As we know, curvaton model can easily generate a large local-type non-Gaussianity \\cite{Linde:1996gt,Lyth:2002my}. Different group reported different result of data analysis to $f_{NL}^{local}$ based on WMAP 3yr data: \\e 27<f_{NL}^{local}<147 \\q in \\cite{Yadav:2007yy} and \\e -70<f_{NL}^{local}<91 \\q in \\cite{Hikage:2008gy} at $95\\%$ CL. WMAP group reported the 5yr result \\cite{Komatsu:2008hk}: \\begin{equation}\\label{fnlexp} -9<f_{NL}^{local}<111 \\quad (95\\% \\ \\hbox{CL}). \\end{equation} Up to now a Gaussian distribution of the primordial curvature perturbation is still consistent with experiments. However, the lower bound on $f_{NL}^{local}$ in Eq.\\eqref{fnlexp} is greatly reduced from WMAP3 to WMAP5. If a large positive $f_{NL}^{local}$ is confirmed by more years integration of WMAP, or the Planck mission, it can help us to distinguish curvaton model from inflation model. Recently many issues related to curvaton were widely discussed in \\cite{Huang:2008ze,Huang:2008rj,Enqvist:2008gk}.  Curvaton is a light scalar field with weak self-interaction. Even though its energy density is subdominant during inflation, it can make the main contribution to the primordial curvature perturbation. In curvaton model, $f_{NL}$ is inversely proportional to the fraction of curvaton energy density in the energy budget at the epoch of curvaton decay, and the value of $g_{NL}$ is mainly contributed by the curvaton self-interaction. In the literatures, the interaction term of curvaton is assumed to be neglected and then $g_{NL}$ is small. Now the cosmological observations become more and more accurate, and even $g_{NL}$ can be detected in the near future if it is not too small. So it is worth working out how $g_{NL}$ is related to the curvaton self-interaction.   In \\cite{Enqvist:2008gk,Enqvist:2005pg} the authors only considered the case with small positive interaction term and they concluded that $g_{NL}\\sim {\\cal O}(-10^4)-{\\cal O}(-10^5)$ for the reasonable cases. In this paper we will see that it is also reasonable to consider the case with small negative interaction term for curvaton model and discuss the dynamics of curvaton more carefully.   This paper is organized as follows. In Sec. 2, we discuss some possible forms of the curvaton potential. In Sec. 3, the dynamics of curvaton with potential slightly deviating from quadratic form is studied in detail. In Sec. 4, we treat the non-quadratic potential as a perturbation, and analytically calculate $f_{NL}$ and $g_{NL}$ to the leading order. How to determine the non-quadratic term from experiments will be investigated in Sec. 5. At the end, we give a summary in Sec. 6.    \\setcounter{equation}{0} ",
        "conclusions": "The fluctuations of curvaton field evolve non-linearly on the superhorizon scales if the curvaton potential deviates from the exactly quadratic form. In this paper, we suggest that the leading order of the non-quadratic term for curvaton field can be negative, for example in the axion-type curvaton model. We also investigate the curvaton dynamics and the non-linearity parameters in curvaton model with non-quadratic curvaton potential in detail. For a positive non-quadratic term, $f_{NL}$ can be very small even when $f_D\\ll 1$. But it does not happen in the case with a negative non-quadratic term. We also see that the second order non-linearity parameter $g_{NL}$ measures the size of the self-interaction of curvaton field. If $g_{NL}$ is positive, the leading order of the self-interaction term should be negative, and the axion-type curvaton model is preferred. The next generation of experiments such as Planck will improve the accuracy to $\\Delta f_{NL}\\sim 6$ and $\\Delta\\tau_{NL}\\sim 560$ \\cite{Kogo:2006kh}. We hope the non-linearity parameters will be detected soon.    In this paper, the self-interaction term is treated as a perturbation. The case in which the self-interaction term dominates the curvaton potential during inflation will be discussed in \\cite{Huang:2008zj}. On the other hand, multiplicity of curvaton fields is expected in the fundamental theories going beyond the standard model. In particular, axions are typically present in large numbers in string compactifications. In \\cite{Huang:2008rj}, these axion fields are suggested to be taken as curvatons. It is worth investigating the non-linearity parameters in N-vaton model \\cite{Huang:2008rj} carefully in the future.     \\vspace{1cm}  \\noindent {\\bf Acknowledgments}  We would like to thank E.~J.~Chun and M. Li for useful discussions.     \\newpage"
    },
    "0808/0808.1397_arXiv.txt": {
        "abstract": "Two of the three instruments proposed to ESO for the second generation instrumentation of the VLTI would use integrated optics for beam combination. Several design are studied, including co-axial and multi-axial recombination. An extensive quantity of combiners are therefore under test in our laboratories. We will present the various components, and the method used to validate and compare the different combiners. Finally, we will discuss the performances and their implication for both VSI and Gravity VLTI instruments. ",
        "introduction": "To pursue the development and increase the capabilities of the VLTI, ESO selected three instrumental projects in phase A. The three new instruments are : MATISSE, VSI, and GRAVITY. MATISSE is a mid-infrared spectro-interferometer. VSI and GRAVITY are two near-infrared instruments, both with the specificity of using integrated optics beam combiners \\cite{1999A&AS..138..135M,1999A&AS..139..173B,2000ApOpt..39.2130H, 2001A&A...376L..31B,2002A&A...390.1171L,2006A&A...450.1259L}. To that end, new integrated optics (IO) beam combiners were developed. The requirements for VSI were: \\begin{itemize} \\item Being able to work in the J, H and K band. \\item Being able to combine up to 6 telescopes. \\end{itemize} on the other hand, GRAVITY had several criterium to meet: \\begin{itemize} \\item High sensitivity in the K band (more than 50\\% throughput for the beam combiner as a whole). \\item Non-temporal modulation to allow long integration time exposures. \\end{itemize}  The basic concept for the combiner was an ABCD type-recombination \\cite{1977JOSA...67...81S}. Such a combiner was already presented by M. Benisty et al. \\cite{2006SPIE.6268E..73B}. Multi-axial 6 beam combiners component are also investigated as a good alternative.  The goal of this new development round is to: (i) prove the validity of the technology for the K and J band, and (ii) improve the throughput and therefore the sensitivity by testing different arrangement of the waveguide paths. Several new combiners were therefore designed, and were recently delivered to our laboratory. ",
        "conclusions": "OI circuits presented in Figure~\\ref{fig:waf} have recently been delivered to our laboratory. We have shown that the transmission of the components is adapted for astronomical observations in the K band. Optimization of the throughput is nevertheless still under investigation, with possibilities offered by using new couplers, optical paths, and/or fiber to combiner injection methods.  We also developed the tools to fully characterize the complex transfer function of the new components, to have rigorous comparison between the different combiner engineered. The knowledge of the $V2PM$ matrix should also allow signal to noise estimation, to provide sound and practical way to choose the best combiner for the VLTI."
    },
    "0808/0808.4117_arXiv.txt": {
        "abstract": "Dust plays an essential role in the unification theory of active galactic nuclei (AGNs). This review summarizes our current understanding of the extinction and infrared emission properties of the circumnuclear dust in AGNs as well as the inferred dust composition and size distribution. ",
        "introduction": "\\vspace{-3mm} Dust is the cornerstone of the unification theory of active galactic nuclei (AGNs). This theory proposes that all AGNs are essentially ``born equal'': all types of AGNs are surrounded by an optically thick dust torus and are basically the same object but viewed from different lines of sight (see e.g. Antonucci 1993; Urry \\& Padovani 1995). The large diversity in the observational properties of AGNs (e.g. optical emission-line widths and X-ray spectral slopes) is simply caused by the viewing-angle-dependent obscuration of the nucleus: those viewed face-on are unobscured (allowing for a direct view of their nuclei) and recognized as ``type 1'' AGNs, while those viewed edge-on are ``type 2'' AGNs with most of their central engine and broad line regions being hidden by the obscuring dust.  Apparently, key factors in understanding the structure and nature of AGNs are determining the geometry of the nuclear obscuring torus around the central engine and the obscuration (i.e. extinction, a combination of absorption and scattering) properties of the circumnuclear dust. An accurate knowledge of the dust extinction properties is also required to correct for the dust obscuration in order to recover the intrinsic optical/ultraviolet (UV) spectrum of the nucleus from the observed spectrum and to probe the physical conditions of the dust-enshrouded gas close to the nucleus.  The presence of an obscuring dust torus around the central engine was first indirectly indicated by the spectropolarimetric detection of broad permitted emission lines (characteristic of type 1 AGNs) scattered into our line of sight by free electrons located above or below the dust torus in a number of type 2 AGNs (e.g. see Heisler et al.\\ 1997, Tran 2003). Direct evidence for the presence of a dust torus is provided by infrared (IR) observations. The circumnuclear dust absorbs the AGN illumination and reradiates the absorbed energy in the IR. The IR emission at wavelengths longward of $\\lambda$\\,$>$\\,1$\\mum$ accounts for at least 50\\% of the  bolometric luminosity of type 2 AGNs. For type 1 AGNs, $\\simali$10\\% of the bolometric luminosity is emitted in the IR (e.g. see Fig.\\,13.7 of Osterbrock \\& Ferland 2006). A near-IR ``bump'' (excess emission above the $\\simali$2--10$\\mum$ continuum), generally attributed to hot dust with temperatures around $\\simali$1200--1500$\\K$ (near the sublimation temperatures of silicate and graphite grains), is seen in a few type 1 AGNs (Barvainis 1987; Rodr\\'{i}guez-Ardila \\& Mazzalay 2006). Direct imaging at near- and mid-IR wavelengths has been performed for several AGNs and provides constraints on the size and structure of the circumnuclear dust torus (e.g. see Jaffe et al.\\ 2004, Elitzur 2006). Spectroscopically, the 10$\\mum$ silicate {\\it absorption} feature (see \\S3.3) and the 3.4$\\mum$ aliphatic hydrocarbon {\\it absorption} feature (see \\S3.2) are widely seen in heavily obscured type 2 AGNs; in contrast, the 10$\\mum$ silicate {\\it emission} feature has recently been detected in a number of type 1 AGNs (see \\S3.3).  To properly interpret the observed IR continuum emission and spectroscopy as well as the IR images of AGNs, it requires a good understanding of the absorption and emission properties of the circumnuclear dust. To this end, one needs to know the composition, size, and morphology of the dust -- with this knowledge, one can use Mie theory (for spherical dust) to calculate the absorption and scattering cross sections of the dust from X-ray to far-IR wavelengths, and then calculate its UV/optical/near-IR obscuration as a function of wavelength, and derive the dust thermal equilibrium temperature (based on the energy balance between absorption and emission) as well as its IR emission spectrum. This will allow us to correct for dust obscuration and constrain the circumnuclear structure through modeling the observed IR emission and images. The former is essential for interpreting the obscured UV/optical emission lines and probing the physical conditions of the central regions; the latter is critical to our understanding of the growth of the central supermassive black hole.  However, little is known about the dust in the circumnuclear torus of AGNs. Even our knowledge of the best-studied dust -- the Milky Way interstellar dust -- is very limited. In this review, I will take a comparative study of the extinction and IR emission as well as the UV/IR spectroscopic properties and the inferred composition, size and morphology of the dust in AGNs and the dust in the interstellar medium (ISM) of the Milky Way and other galaxies.  \\vspace{-4.5mm} ",
        "conclusions": ""
    },
    "0808/0808.3501_arXiv.txt": {
        "abstract": "This paper reports on the Suzaku results of thermal and non-thermal features of 30~Dor~C, a supernova remnant (SNR) in a superbubble of the Large Magellanic Cloud (LMC). The west rim exhibits a non-thermal X-ray spectrum with no thermal component. A single power-law model is rejected but a power-law model with spectral cutoff is accepted. The cutoff frequency of $(3-7)\\times 10^{17}$~Hz is the highest among the shell type SNRs like SN 1006 ($\\sim 6\\times 10^{16}$~Hz), and hence 30~Dor~C would be the site of the highest energy accelerator of the SNR shock. The southeast (SE) and northeast (NE) rims have both the thermal and non-thermal components. The thin-thermal plasmas in the both rims are in collisional ionization equilibrium state. The electron temperature of the plasma in the SE rim ($kT_e \\sim 0.7$~keV) is found to be higher than the previously reported value. The power-law index from SE is nearly the same as, while that from the NE is larger than that of the West rim. The SNR age would be in the range of $(4-20)\\times 10^3$~yr. Thus, 30~Dor~C is likely to be the oldest shell-like SNR with non-thermal emission. ",
        "introduction": "\\label{sec:introduction}  OB associations typically contain a few 10 massive stars. The combined actions of fast stellar winds and core-collapse supernova (SN) explosions of the member stars create large ($\\gtrsim 10$~pc) shell-like structures, called superbubbles (SBs), by sweeping up the ambient medium (e.g., Mac Low \\& McCray 1988). Therefore, SBs are very energetic objects which store a large amount of energy injected from massive stars; the swept up interstellar medium (ISM) can be successively heated to high temperature by  stellar winds and/or SNe. Furthermore, large tenuous cavities created inside the SB walls allow that the blast shocks of the interior supernova remnants (SNRs) expand rapidly without decelerating for a long time. Therefore, the timescale of efficient cosmic-ray acceleration can be much longer than that of most isolated SNRs.   For the study of SNRs and SBs, the Large Magellanic Cloud (LMC) is an ideal site, because of little foreground extinction and well-known distance (50~kpc: Alves 2004). Dunne, Points, Chu~(2001) systematically studied several SBs in the LMC using ROSAT, and showed their X-ray emissions are brighter than that theoretically expected for a wind-blown bubble. This fact suggests that the X-ray emitting walls of the SBs are re-heated by the shock of interior SNRs, and hence the emissions are enhanced. They also found that the plasma temperatures of the shocked SB walls are typically of the order of 0.1~keV.   30~Dor~C, located in the southwestern region of the 30~Doradus complex, is one of the SBs in the LMC identified by the morphology of the H$\\alpha$ emission (Mathewson et al.~1985; Kennicutt \\& Hodge 1986). This SB hosts an OB association LH~90, with the age of a few Myr (Lucke \\& Hodge 1970). A shell-like structure with a diameter of $\\sim \\timeform{6'}$ ($\\sim 87$~pc at 50~kpc distance) was discovered by the radio observation (Mills et al.~1984).   X-rays from 30~Dor~C were first detected with the Einstein satellite (Long et al.~1981). Then, Dunne, Points, Chu~(2001) found a clear shell-like structure confined within the H$\\alpha$ shell by the ROSAT observation. They thus concluded that the X-ray emission arises from the interaction between the SB and an interior SNR. Bamba et al.~(2004: hereafter B04) discovered synchrotron X-rays from the shell region of 30~Dor~C with Chandra and XMM-Newton. They found that the total luminosity of the synchrotron emission was $\\sim 10$ times larger than that of SN~1006 (Koyama et al.~1995), and hence argued that the large energy supply was probably by multiple SN explosions.   This paper reports the X-ray observation of 30~Dor~C with Suzaku (Mitsuda et al.~2007).  Utilizing the good sensitivity of Suzaku, we study the detailed properties of the thermal and non-thermal emissions from 30~Dor~C. The distance to the LMC is assumed to be 50~kpc, following Alves~(2004). ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Thermal Emission} \\label{ssec:thermal}  The thermal component of the SE spectrum was well represented by the CIE plasma model with the electron temperature of $\\sim 0.7$~keV. Since the bright H$\\alpha$ emission coincides with the SE rim (Mathewson et al.~1985), the origin of the thermal component is considered to be the SB shell re-shocked by the blast-wave of the interior SNR. In addition, we detected the significant thin-thermal component from the NE rim, for the first time. This region is also correlated with the H$\\alpha$ emission, albeit the brightness is lower than that of the SE region. Therefore, the same origin with the SE can be considered. The low elemental abundance of Fe relatively to the other lighter elements is probably due to that the SB shell is dominantly contributed by stellar winds and core-collapse SNRs occurred in the past.   The size of the SE elliptic region is $\\timeform{1.2'} \\times \\timeform{1.1'}$, which corresponds to 18~pc $\\times$ 16~pc at a distance of 50~kpc. Assuming the plasma depth to be 22~pc, a half radius of the SNR, the emission volume is estimated to be $V = 5.7 \\times 10^{59}$~cm$^3$. Thus, the EM of the VAPEC component in the SE region corresponds to the proton and electron densities of $n_p \\simeq n_e = 0.37~f_{\\rm SE}^{-0.5}$~cm$^{-3}$, where $f_{\\rm SE}$ is the filling factor. The plasma mass is, therefore, estimated to be $M = 175~f_{\\rm SE}^{0.5}$~\\Mo. Similarly, from the size of the NE ellipse of $\\timeform{1.9'} \\times \\timeform{0.9'}$, the plasma densities and the mass are respectively estimated to be $n_p \\simeq n_e = 0.12~f_{\\rm NE}^{-0.5}$~cm$^{-3}$ and $M = 57~f_{\\rm NE}^{0.5}$~\\Mo, where $f_{\\rm NE}$ is the filling factor for the VAPEC component in the NE region.   Both the SE and NE plasmas were found to be in CIE condition. For full ionization equilibrium, the ionization timescale defined as $n_et$ is required to be $\\geq~10^{12}$~cm$^{-3}$~s, where $t$ is the plasma age or the time since the gas was shock-heated (Masai 1984). Therefore, the plasma ages of the SE and NE plasmas are estimated to be higher than $\\sim 9\\times 10^4$~yr and $\\sim 3\\times 10^5$~yr, respectively. As mentioned in subsection~\\ref{ssec:non-thermal}, the real age of the SNR is likely to be lower than $\\sim 2\\times 10^4$~yr. Thus, the plasma ages are significantly older than the SNR age. This fact can be interpreted that the SB shell had been ionized by preceding stellar winds and/or SNe, before the re-heating by the interior SNR. In CIE plasma, energy equipartition between electron and protons can be assumed. Therefore, the thermal energies ($E = 3n_e V kT_e$) of the SE and NE plasmas are estimated to be $6.7\\times 10^{50}$~ergs and $1.2\\times 10^{50}$~ergs, respectively.   The electron temperature derived for the SE rim (0.58--0.76~keV) is significantly higher than the value of the previous report of B04 (0.19--0.23~keV). If we fitted the SE spectrum fixing the temperature to 0.21~keV, the best-fit value of B04, the model failed to reproduce the fluxes of the Ne\\emissiontype{X} and Mg\\emissiontype{XII} K$\\alpha$ lines, which are clearly detected in our spectrum (see figure~\\ref{fig:se}). Indeed, the K$\\alpha$-line flux ratios of Ne\\emissiontype{X}/Ne\\emissiontype{IX} and Mg\\emissiontype{XII}/Mg\\emissiontype{XI} predicted by the 0.21~keV CIE plasma are only less than 10\\%. Since the emission lines from these hydrogen-like ions were not detected in the Chandra spectrum probably due to the poor photon statistics and relatively poor energy resolution, the systematically lower temperature would be obtained by B04.   \\subsection{Non-Thermal Emission} \\label{ssec:non-thermal}  The high quality spectrum of the West rim rejected the simple power-law model. We found that either the broken power-law or the SRCUT was required to fit the West spectrum, for the first time. Similar spectral steepening had been found in the synchrotron X-rays from the other SNRs, SN~1006 (Bamba et al.~2008) and  RX~J1713.7--3946 (Takahashi et al.~2008). Therefore, the previous argument (B04) becomes more robust that the non-thermal X-rays from 30~Dor~C are synchrotron emissions from high energy electrons accelerated by the shock front of the SNR.  In the SRCUT model for the West rim, the roll-off frequency was obtained to be (3--7)$\\times 10^{17}$~Hz for the typical range of the spectral index at the radio band of $\\alpha =$ 0.5--0.6. We also tried to fit the SE and NE spectra using the SRCUT models with $\\alpha = 0.5$ instead of the power-law components. Then, the roll-off frequencies were obtained to be $> 1.1\\times 10^{17}$~Hz and (0.7--1.4)$\\times 10^{17}$~Hz, respectively. All these values are significantly higher than that of SN~1006 ($\\sim 6\\times 10^{16}$~Hz: Bamba et al.~2008), and similar to RX~J1713.7--3946 ($\\sim 2\\times 10^{17}$~Hz: Takahashi et al.~2008), the highest value among the shell-like SNRs. Since $\\nu _{\\rm rolloff} \\propto BE_e^2$, where $B$ and $E_e$ are the magnetic field and the maximum energy of accelerated electrons (Reynolds 1998), the electrons would be accelerated to higher energy than that in the case of SN~1006, if the comparable strengths of the magnetic field are assumed.   30~Dor~C has a radius of $\\sim 44$~pc, which is much larger than any other shell-like SNRs with non-thermal emission (e.g., RX~J1713.7--3946, RX~J0852.0--4622, RCW~86: $r < 20$~pc). If we simply assume that the blast-wave have been expanded almost freely in the cavity with the initial speed (possibly $\\sim 10^4$~km~s$^{-1}$), then the SNR age would be $\\sim 4\\times 10^3$~yr. These may be the lower limit of the age. On the other hand, in order to emit intense synchrotron X-rays, the shock speed higher than $\\sim 2000$~km~s$^{-1}$ is required by the standard diffusive shock acceleration theory (e.g., Aharonian \\& Atoyan 1999). Therefore, the upper limit of the age is estimated to be $\\sim 2\\times 10^4$~yr. We thus conclude that the age of 30~Dor~C would be between $\\sim 4\\times 10^3$~yr and $\\sim 2\\times 10^4$~yr; 30~Dor~C is the oldest among the known non-thermal SNRs.   Our observation demonstrates even older-system can be a site of higher energy acceleration than typical young SNRs, such as SN~1006. RCW~86 is another shell-like SNR with non-thermal shell, which has a relatively old age of $\\sim 1800$~yr. Similarly to 30~Dor~C, RCW~86 also coincides with the OB association, and is considered to be fast expanding inside the large cavity (Vink et al.~2006; Yamaguchi et al.~2008). The long timescales of the efficient acceleration, achieved in 30~Dor~C and RCW~86, would be due to the extremely low density of the cavity created inside the SB shell; the SNRs in SBs can be better cosmic-ray accelerator than the isolated SNRs in homogeneous ISM.   Non-thermal X-rays had been detected from other star-forming regions, RCW~38 (Wolk et al.~2002) and Westerlund~1 (Muno et al.~2006) in our Galaxy, and DEM~L192 (N51D: Cooper et al.~2004) in the LMC. However, the emissions from these regions show center-filled morphologies, in contrast to those from 30~Dor~C and RCW~86. Moreover, their photon indexes (e.g., $\\Gamma \\sim 1.3$ for DEM~L192) are lower (harder) than those of typical shell-like SNRs. Therefore, the origins of these emissions are still unclear. Future detailed studies of the thermal and non-thermal properties of these star-forming regions using sensitive observatory, such as Suzaku, would help to reveal their origins. Observations of more SBs are also necessary to determine if non-thermal emissions are commonly present or not in SB regions.        We have analyzed Suzaku/XIS data of 30~Dor~C. The results are summarized as follows:  \\begin{enumerate}  \\item  The West rim is dominated by the non-thermal emission, whereas the SE and NE rims show the thermal and non-thermal emissions.  \\item  The spectrum from the West rim is well-reproduced by a power-law with spectral cutoff. The cutoff frequency of $(3-7)\\times 10^{17}$~Hz is the highest, although the age is the oldest, among the known shell-like SNRs with non-thermal emission.  \\item  The thermal components of the SE and NE spectra are well-represented by the plasma at a collisional ionization equilibrium.  \\item  The electron temperature of the thermal plasma in the SE rim ($kT_e \\sim 0.7$~keV) is found to be higher than the value of the previous report.  \\end{enumerate}   \\bigskip The authors would like to thank the Suzaku Science Working Group and the members of Suzaku SN~1987A team. We also thank Yasunobu Uchiyama for his useful comments. H.Y.~is supported by the Special Postdoctoral Researchers Program in RIKEN. A.B.~is supported by JSPS Research Fellowship for Young Scientists."
    },
    "0808/0808.1218_arXiv.txt": {
        "abstract": "In previous studies, we have examined a resonant excitation of disk oscillations in deformed disks. In these studies, however, mathematical treatment around the resonant points was not rigorous. In this paper the inadequate point is corrected, with no essential changes in the final results. For this excitation process to work, disks must be general relativistic. That is, the non-monotonic radial distribution  of epicyclic frequency in relativistic disks is essential for the presence of the resonance and for trapping of oscillations. In this paper, the growth rate of resonant oscillations is expressed in a form more suitable for numerical calculations. ",
        "introduction": "In previous papers (Kato 2004, 2008a), we proposed a resonant excitation mechanism of disk oscillations in deformed disks. The purpose is to propose an excitation mechanism of quasi-periodic oscillations (QPOs) observed in neutron-star and black-hole X-ray binaries (e.g., Kato and Fukue 2007, Kato 2008b). In this resonant excitation model of disk oscillations, a deformation of disks from an axially-symmetric equilibrium state is essential. The deformation to be considered is a warp or an eccentric deformation of disks in the equatorial plane.  An outline of the model is as follows. A non-linear coupling between a disk oscillation (hereafter we call it the original oscillation) and a deformed part of the disks (warp or eccentric deformation) brings about some forced disk oscillations (we call them intermediate oscillations). The intermediate oscillations make a resonant coupling with the unperturbed disk at particular radii of the disk. After this resonant coupling, the intermediate oscillations feedback to the original oscillation by a nonlinear coupling with the deformed part of the disk. Since this nonlinear feedback process involves a resonance, the original oscillation is amplified or dampened. Kato (2004, 2008a) examined this nonlinear feedback processes, and derived an excitation criterion and growth (damping) rate of the resonant oscillations. It is of importance to note that in Keplerian disks this resonant excitation process works only when the disks are general relativistic. That is, a non-monotonic radial distribution of epicyclic frequency is necessary for appearance of resonance and for trapping of oscillations.  In these studies the intermediate disk oscillations were assumed to be local near to the resonant radius in the sense that their radial wavelength is shorter than the characteristic radial length of disks. That is, the radial derivative, $\\partial /\\partial r$, operated to the wave quantities was taken to be $ik$. Further, $k$ was assumed to be constant. This treatment of $k=$ const., however, was inadequate around the resonant point. If this inadequate treatment is properly corrected, we find that unlike in the previous papers, the resonant point is not the place where a local dispersion relation of the intermediate oscillations is satisfied with a constant $k$, but the radii of Lindblad resonance for intermediate oscillations (when resonance occurs by horizontal motions) or the radii where the intermediate oscillations are trapped in the vertical direction (when resonance occurs by vertical motions).  In this paper, we modify the analyses in our previous papers so that the above-mentioned inadequate treatment is corrected. We apply the mathematical techniques used by Meyer-Vernet and Sicardy (1987) in their study of resonant disk-satellite interaction. The results show that the stability criterion derived in the previous papers are unchanged. The expression for growth rate obtained in the previous papers is found to be still applicable, but in this paper the expression is changed in a form more suitable to numerical calculations.  We first summarize in section 2 the basic equations and relations necessary to the study of the present resonant excitation problem, although they are given in the previous papers. In section 3, the stability criterion and growth rate are derived to the case of pressure-less disks as it is instructive, and to the case of disks with pressure in section 4. The final section is devoted to a brief discussion on a meaning of the instability criterion. ",
        "conclusions": ""
    },
    "0808/0808.0608_arXiv.txt": {
        "abstract": "The first stars were key drivers of early cosmic evolution. We review the main physical elements of the current consensus view, positing that the first stars were predominantly very massive. We continue with a discussion of important open questions that confront the standard model. Among them are uncertainties in the atomic and molecular physics of the hydrogen and helium gas, the multiplicity of stars that form in minihalos, and the possible existence of two separate modes of metal-free star formation. ",
        "introduction": "We have learnt a great deal over the last decade concerning the formation of the first stars, the so-called population III or Pop.~III. A consensus has emerged regarding the properties of the protogalaxies (or `minihalos') in which the first stars formed and the major physical and chemical processes involved in their formation. Their masses remain uncertain, but are widely expected to be significantly larger than the characteristic mass for present-day star formation. Good summaries of the present state of the field can be found in \\citet{blar04}, \\citet{glo05} and \\citet{nm08}. Nevertheless, there remain some important open questions. In this review, we discuss four of the most important of these issues and the efforts being made to resolve them. These issues are the impact of chemical rate coefficient uncertainties on the accuracy of our models of Pop.~III star formation; the perennial question of whether we have identified all of the important physical processes responsible for cooling the gas; the number of population III stars that form in each minihalo; and the question of whether population III actually consists of two sub-populations (Pop.~III.1 and Pop.~III.2) with different characteristic masses. Two further important issues -- the question of what physical process terminates accretion onto the earliest protostars, and the impact of dark matter decay and annihilation on the formation of the first stars -- are not discussed here as they are covered in detail elsewhere in these proceedings (see e.g.\\ the contributions by Whalen, Tan, Freese \\& Iocco). ",
        "conclusions": ""
    },
    "0808/0808.2111_arXiv.txt": {
        "abstract": "{For an understanding of Galactic stellar populations in the SDSS filter system well defined stellar samples are needed. The nearby stars provide a complete stellar sample representative for the thin disc population. We compare the filter transformations of different authors applied to the main sequence stars from F to K dwarfs to SDSS filter system and discuss the properties of the main sequence. The location of the mean main sequence in colour-magnitude diagrams is very sensitive to systematic differences in the filter transformation. A comparison with fiducial sequences of star clusters observed in g',r',i' show good agreement. Theoretical isochrones from Padua and from Dartmouth have still some problems especially in (r-i)-colour.} ",
        "introduction": "}  The Sloan Digital Sky Survey (SDSS) with  sky coverage of almost 10,000 deg$^2$ has collected at present the largest and most homogeneous database comprising about $10^8$ stellar objects in the Milky Way. The SDSS photometric system was specifically designed for the survey to provide continuous coverage over the entire optical wavelength range in the five filters $u\\, g\\, r\\, i\\, z$. In the most recent sixth data release (Adelman-McCarthy et al. \\cite{Ad08}) the new 'ubercalibration' of the photometry was applied to provide an improved homogeneous relative calibration over the whole survey to 1\\% in $g\\, r\\, i\\, z$ and 2\\% in $u$ (Padmanabhan et al. \\cite{Pa08}). The zero-points are not exactly in the AB system, but for the $g\\, r\\, i$ filters, the corrections are below 1\\% and can be neglected. For the zero-points of stars brighter than 14.5\\,mag see also Chonis \\& Gascell (\\cite{Ch08}).  Since most standard stars are too bright to be directly observed with the 2.5\\,m SDSS telescope, it is a complicated matter to establish standards. One way is to measure the standard stars with a smaller telescope. The SDSS photometric system was originally defined at the 0.5\\,m photometric telescope (PT) with $u'\\, g'\\, r'\\, i'\\, z'$ (Fukugita et al. \\cite{Fu96}). A similar filter set is used in the USNO 1\\,m telescope for these calibration issues. Since these filters are slightly different from $u\\, g\\, r\\, i\\, z$ at the main 2.5\\,m telescope (Abazajian et al. \\cite{Ab03}), transformation formulas from $u'\\, g'\\, r'\\, i'\\, z'$ to $u\\, g\\, r\\, i\\, z$ must be established (Tucker et al. \\cite{Tu06}).  Despite the enormous wealth of information that the SDSS database provides the majority of current observational and theoretical knowledge of resolved stellar populations is based largely on the Johnson-Kron-Cousins $U\\,B\\,V\\,R_\\mathrm{C}\\,I_\\mathrm{C}$ photometric system and some other systems such as the Str\\\"omgren, DDO, Vilnius, and Geneva systems.  The best way of calibrating the theoretical isochrones is the observation of fiducial isochrones of simple stellar populations (SSP) with known age and metallicity in SDSS filters. Clem, VandenBerg, \\& Stetson (\\cite{Cl08}) obtained fiducial stellar population sequences for the $u'\\, g'\\, r'\\, i'\\, z'$ photometric system from a number of different Galactic star clusters, spanning a wide range in metallicity and magnitude. In order to achieve the required accuracy they have first established a new set of secondary standard stars in selected open and globular star clusters for this photometric system (Clem, VandenBerg, \\& Stetson \\cite{Cl07}).  For a physical understanding of observed stellar populations is essential to interpret the data via stellar evolution tracks or isochrones. Girardi et al. (\\cite{Gi04}) were the first, who provided tables of theoretical isochrones in $u\\, g\\, r\\, i\\, z$ derived from stellar spectra. These have however turned out to lack the necessary precision to correctly interpret different stellar populations from the SDSS database. Recently in the Dartmouth Stellar Evolution Program (Dotter et al. \\cite{Do07,Do08}) isochrones in $u\\, g\\, r\\, i\\, z$ are presented based on VandenBerg \\& Clem (\\cite{Va03}) and Ivezi\\'c et al. (\\cite{Iv07}).  The solar neighbourhood provides a complete stellar sample of the Milky Way disc with measured absolute magnitudes from Hipparcos data and the Catalogue of Nearby Stars (CNS4, Jahreiss \\& Wielen \\cite{jah97}). Therefore it is a unique place for investigating the properties of the stellar disc population testing the transformations to the SDSS filters. The purpose of this paper is to transform the CNS4 absolute luminosities in ($M_\\mathrm{B}\\,M_\\mathrm{V}\\,M_{\\mathrm{R_C}}\\,M_{\\mathrm{I_C}}$) of main sequence (MS) stars to the SDSS luminosities ($M_\\mathrm{g}\\, M_\\mathrm{r}\\, M_\\mathrm{i}$), and to derive the corresponding colour-magnitude relations.  In the next section we present the filter transformations of different authors. In Sect. \\ref{sec-stars} the selection of the stellar sample and the available colours are presented. In Sect. \\ref{sec-trafo} we discuss the differences of the various transformation formulas. In Sect. \\ref{sec-ms} we derive the MS properties for F-K stars and discuss the uncertainties due to different transformations. A comparison with observed fiducial isochrones is also presented and the local stellar number density in $g-r$ space is derived. In Sect. \\ref{sec-iso} a comparison with theoretical isochrones in the colour magnitude diagrams (CMD) is given. The results are summarized in Sect. \\ref{sec-disc}. ",
        "conclusions": "}  We transformed the MS of F to K dwarfs in the solar neighbourhood from Johnson Cousins $B\\,V\\,R_\\mathrm{C}\\,I_\\mathrm{C}$ to SDSS colours $g\\,r\\,i$ using the transformations of different authors. We used the complete sample of stars in a 25\\,pc sphere around the Sun from the Catalogue of Nearby Stars (CNS4) with Hipparcos distances to determine the mean MS properties in ($g-r,M_\\mathrm{g}$) with the transformation C08 of Chonis \\& Gascall (\\cite{Ch08}). With the subsample, where all colours are available, we determined the mean MS in the CMDs in ($g-r,M_\\mathrm{g}$) and ($r-i,M_\\mathrm{r}$) and the two-colour diagram ($g-r,r-i$). The systematic differences between the transformations of different authors in $M_\\mathrm{g}$, $M_\\mathrm{r}$, $g-r$ and $r-i$ are below 0.1\\,mag. In the cases, where $V-R_\\mathrm{C}$ is used in the transformations a considerable individual scatter up to 0.3\\,mag are introduced. This was not expected, when using more colour-terms in the transformations. The position of the mean MS in the CMDs is  very sensitive to systematic deviations in the transformations. The differences in $M_\\mathrm{g}(g-r)$ are up to 0.5\\,mag and in $M_\\mathrm{r}(r-i)$ up to 0.3\\,mag, which is significant for photometric distance estimations. Also the 'bright' MS of Juri\\'{c} et al. (\\cite{Ju08}) is consistent with the local stellar sample.  The distribution of the solar neighbourhood stars in the CMDs and the two-colour diagram agree reasonably well with fiducial colour magnitude sequences of star clusters observed in $g'\\,r'\\,i'$. The best agreement occurs with the transformation C08 of Chonis \\& Gascall (\\cite{Ch08}).  The comparison with theoretical isochrones, especially in ($r-i,M_\\mathrm{r}$) and ($g-r,r-i$), is less  satisfying. The systematic differences in shape and location require further calibration by observations in the SDSS filter system.  The solar neighbourhood provides a good testcase for the interpretation of mixed stellar populations. The results presented here can be used for detailed models of the Milky Way disc based on star counts, when extrapolated to the bright end in apparent magnitude. For this application we provide the local stellar number densities as function of colour $g-r$."
    },
    "0808/0808.0925_arXiv.txt": {
        "abstract": "{The $21\\,\\rm{cm}$ emission of neutral hydrogen is the most promising probe of the epoch of reionization (EoR). In the next few years, the SKA pathfinders will provide statistical measurements of this signal, such as its power spectrum. Within one decade, SKA should produce a full tomography of the signal. Numerical simulations predicting these observations are necessary to optimize the design of the instruments and help in the interpretation of the data. Simulations are reaching a reasonable level in terms of scale range, but still often rely on simplifications to compute the strength of the signal. The main difficulty is the computation of the spin temperature of neutral hydrogen which depends on the gas kinetic temperature and on the level of the local Lyman-$\\alpha$ flux (The Wouthuysen-Field effect). A $T_S \\gg T_{CMB}$ assumption is usual. However, this assumption does not apply early in the reionization history, or even later in the history  as long as the sources of X-rays are too weak to heat the intergalactic medium significantly. This work  presents the first EoR numerical simulations including, beside dynamics and ionizing continuum radiative transfer, a self-consistent treatment of the Ly-$\\alpha$ radiative transfer. This allows us to compute the spin temperature more accurately. We use two different box sizes, $20\\,h^{-1}\\,\\rm{Mpc}$ and $100\\,h^{-1}\\,\\rm{Mpc}$, and a star source model. Using the redshift dependence of average quantities, maps, and power spectra, we quantify the effect of using different assumptions to compute the spin temperature and the influence of the box size.  The first effect comes from allowing for a signal in absorption (i.e. not making the $T_S \\gg T_{CMB}$ approximation). The magnitude of this effect depends on the amount of heating by hydrodynamic shocks and X-rays in the intergalactic medium(IGM). With our source model we have little heating so regions seen in absorption survive almost until the end of reionization. The second effects comes from using the real, local, Lyman-$\\alpha$ flux. This effect is important for an average ionization fraction of less than $\\sim 10\\%$~: it changes the overall amplitude of the $21\\,\\rm{cm}$ signal, and adds its own fluctuations to the power spectrum. } ",
        "introduction": "The Epoch of Reionization (EoR) begins with the formation of the first sources of light resulting from the non-linear growth of primordial density fluctuations. When this event occurs, which could be anytime between $z\\sim 30$ and $z\\sim 15$, the intergalactic medium (IGM) is mostly cold, neutral and optically thick at many wavelengths. Our main probe of baryonic matter during this era is the $21\\,\\rm{cm}$ emission of neutral hydrogen, to which the IGM is optically thin (Madau et al. \\cite{Madau97}). To observe a $21\\,\\rm{cm}$ signal in emission or absorption against the cosmic microwave background (CMB), however, we need either a sufficient local flux of Ly-$\\alpha$ photons (Wouthuysen-Field effect, Wouthuysen \\cite{Wouthuysen52}, Field \\cite{Field58}) or a high enough baryon density (collisions), either of which will decouple the hydrogen spin temperature from the CMB temperature. The first condition requires a sufficient number of sources and the second a sufficient growth of the density fluctuations. During the EoR, the Wouthuysen-Field effect is prevalent as it becomes efficient everywhere with time. Where the $21\\,\\rm{cm}$ signal turns on, its intensity is determined, in addition to the spin temperature, by the neutral hydrogen density field which is the combination of the baryon density field and the complex distribution of growing ionized regions. This means that observing the $21\\,\\rm{cm}$ signal will teach us a lot about the reionization history. Presently, what is known about this history defines two constraints. The first is the value of $\\tau$, the optical depth of the Thomson scattering of CMB photons by free electrons, which is proportional to the column density of ionized hydrogen from us to the CMB. This  value is $0.09 \\pm 0.03$ for {\\sl WMAP3} cosmology (Spergel et al. \\cite{Spergel07}) and $0.087 \\pm 0.017$ in {\\sl WMAP5} cosmology (Komatsu et al. \\cite{Komatsu08}). In the over-simplified model of instantaneous reionization, this gives a reionization redshift of $z \\sim 11$. In realistic models, $\\tau$ provides a constraint on a redshift integral of the mean ionization fraction. The second constraint on reionization arises from the spectra of high-$z$ quasars. The Gunn-Peterson trough appears as a sharp feature in the quasar spectra when the neutral fraction in the IGM surrounding the quasar rises above $1 \\%$. From observations this happens at $z$ greater than $\\sim 6$ (Fan et al. \\cite{Fan06}). The reduced error bar on the value of $\\tau$ in the {\\sl WMAP5} results, together with the quasar constraint, strongly favor an extended reionization history.  Much information is coded in the $21\\,\\rm{cm}$ emission. First, the nature of the sources (Pop II, Pop III stars or  X-ray sources: QSO, SNe or X-ray binaries) and their formation history determine the evolution of the sky-averaged signal with redshift. It may also be possible to pinpoint specific events in the reionization history like an average ionization fraction of $0.5$ using the redshift dependence of the power spectrum (Mellema et al. \\cite{Mellema06b}, Lidz et al. \\cite{Lidz07b}). This type of information will be available in the next few years from pathfinders like LOFAR or MWA. Within one decade, the Square Kilometer Array (SKA) will deliver a tomography (maps at many different redshifts) of the signal at $1$ comoving Mpc resolution. If the different physics can be separated in the signal (Barkana \\& Loeb \\cite{Barkana05a}, McQuinn et al. \\cite{McQuinn06}), we will be able to extract information about the growth rate of large scale structures directly. If not, it is possible to use them as independent source planes for gravitational lensing by lower $z$ clusters and still deduce information on the structure growth rate (Metcalf \\& White \\cite{Metcalf07}).   Numerical simulations of the $21$ cm signal are useful to explore the different possible source models and derive constraints for the design of the future instruments in terms of frequency range, bandwidth, angular resolution and sensitivity. Comparison with observations will then enable us to discriminate between source models and formation histories. Predicting the $21\\,\\rm{cm}$ signal requires the knowledge of three local quantities~: the baryonic density field, the ionization fraction of hydrogen and the Ly-$\\alpha$ flux (to compute the spin temperature). Moreover, high resolution is critical to obtain realistic, non-spherical ionized regions arising from the high photon consumption in dense small scale structures, and large simulation boxes ($> 100\\,h^{-1}\\,\\rm{Mpc}$) are necessary to correctly sample the large-scale clustering in the source distribution and the abundance of rare sources (Barkana \\& Loeb \\cite{Barkana04}, Iliev et al. \\cite{Iliev06}). With the development of 3D radiative transfer codes (Gnedin \\& Ostriker \\cite{Gnedin97}, Gnedin \\& Abel \\cite{Gnedin01}, Nakamoto et al. \\cite{Nakamoto01}, Maselli et al. \\cite{Maselli03}, Razoumov \\& Cardall \\cite{Razoumov05}, Mellema et al. \\cite{Mellema06a}, Rijkhorst et al. \\cite{Rijkhorst06}, Susa \\cite{Susa06}, Whalen \\& Norman \\cite{Wahlen06}, Trac \\& Cen \\cite{Trac07}, Pawlik \\& Schaye \\cite{Pawlik08}, Altay et al. \\cite{Altay08})  and increases in computational power, it has become possible in the last few years to predict the $21\\,\\rm{cm}$ signal. But some compromises still have been necessary. The most common approximation is to assume $T_S \\gg T_{CMB}$, where $T_S$ is the hydrogen spin temperature and $T_{CMB}$ is the CMB temperature. In this regime the $21\\,\\rm{cm}$ brightness temperature is independent of the spin temperature~: no information about the gas temperature and local Ly-$\\alpha$ flux is needed. Obviously, this approximation fails if either the coupling of $T_S$ to $T_K$, the gas kinetic temperature, by Ly-$\\alpha$ is weak (early and\\slash or far from the sources), then $T_S \\sim T_{CMB}$, or if the gas is not significantly heated in the voids by X-rays. In other words, this approximation removes any possible absorption regime. However, the signal seen in absorption has the potential to be stronger than in emission (Furlanetto et al. \\cite{Furlanetto06a}).  Kuhlen et al. (\\cite{Kuhlen06}), Shapiro et al. (\\cite{Shapiro06}), Gnedin \\& Shaver (\\cite{Gnedin04}) and Santos et al. (\\cite{Santos07}) have probed the absorption regime in numerical simulations.  In Kuhlen et al. and Shapiro et al., this is done taking only the role of collisions into account in determining the $21\\,\\rm{cm}$ signal, i.e. ignoring the role of Ly-$\\alpha$ pumping or prior to the existence of any Ly-$\\alpha$ source. Gnedin \\& Shaver and Santos et al., in turn, add the effect of a Ly-$\\alpha$  flux in their analysis, but  still use a drastic simplification~: a homogeneous Ly-$\\alpha$ flux changing with redshift. Considering a simple $r^{-2}$ profile around the sources to compute the Ly-$\\alpha$ flux fluctuations,  Barkana \\& Loeb (\\cite{Barkana05b}) have shown that clustering and Poisson noise in the source distribution modify the $21\\,\\rm{cm}$ power spectrum. Furthermore, several recent works  (Semelin et al. \\cite{Semelin07}, Chuzhoy \\& Zheng \\cite{Chuzhoy07}, Naoz \\& Barkana \\cite{Naoz08}) have shown that 3D radiative transfer effects modify the expected $r^{-2}$ dependence of the flux in a homogeneous medium, and even more so in an inhomogeneous density field. The present work includes a full modeling of the Ly-$\\alpha$ radiative transfer and examines how the resulting $21\\,\\rm{cm}$ signal compares to the one using a homogeneous Ly-$\\alpha$ flux.  Often working with the $T_S \\gg T_{CMB}$ approximation, the early simulations used box sizes  equal to or smaller than $20\\,h^{-1}\\,\\rm{Mpc}$ (Ciardi \\& Madau \\cite{Ciardi03}, Gnedin \\& Shaver \\cite{Gnedin04}, Furlanetto et al. \\cite{Furlanetto04}, Salvaterra et al. \\cite{Salvaterra05}, Vald\\'es et al. \\cite{Valdes06}). Recent works put the effort on simulating larger boxes, usually $\\sim 100\\,h^{-1}\\,\\rm{Mpc}$, while trying to preserve a good enough resolution (Mellema et al. \\cite{Mellema06b}, Zahn et al. \\cite{Zahn07}, Lidz et al. \\cite{Lidz07a}, McQuinn et al. \\cite{McQuinn07}, Lidz et al. \\cite{Lidz07b}, Iliev et al. \\cite{Iliev08}). For the same resolution, radiative transfer simulations are usually more demanding than dark matter dynamical simulations. Consequently, a common strategy is to run very high resolution DM simulations, derive the baryonic density field with a constant bias, and compute a clumping factor within each larger scale radiative transfer cell. This is a method to include the effect of minihalos (Ciardi et al. \\cite{Ciardi06}) as subgrid physics. In this method, the least robust assumption is the use of a constant baryonic to DM density ratio~: while correct on a large scale, it fails on the scale of clusters where hydrodynamics and heating\\slash cooling processes play an important role. In this work, we rely on self-consistent hydrodynamic simulations and do not use a clumping factor.  Section 2 presents the radiative transfer code used, Section 3 describes the simulation runs, which are analyzed in Section 4. Section 5 summarizes the main results. ",
        "conclusions": "In this work, we have presented simulations of the $21\\,\\rm{cm}$ emission from the EoR. The first hydrodynamic simulations were run in cosmological boxes of sizes $20\\,h^{-1}\\,\\rm{Mpc}$ and $100\\,h^{-1}\\,\\rm{Mpc}$ with $2 \\times 256^3$ particles. Then, using our adaptive Monte-Carlo code LICORICE, we ran, in post-treatment, radiative transfer simulations for the ionizing continuum. Finally cosmological radiative transfer simulations in the Lyman-$\\alpha$ line were performed.  The latter are necessary to compute the neutral hydrogen spin temperature without making any assumption about the strength of the coupling. This approach has not been followed before. We model the sources of both ionizing continuum and Lyman-$\\alpha$  photons as stars, using a Salpeter IMF cut off at $100$ M$_\\odot$. The $100\\,h^{-1}\\,\\rm{Mpc}$ simulation star formation history is calibrated on the $20\\,h^{-1}\\,\\rm{Mpc}$. The photon escape fraction is calibrated independently in both simulations to produce a Thomson scattering optical depth of CMB photons within $1\\sigma$ of the WMAP3 value, and a complete reionization at redshift $z\\sim 6$. We have not included X-ray sources. If we had they would heat up the IGM and reduce the contribution of absorption vs emission in the signal.  We first computed the evolution of a number of box-averaged quantities as functions of the average ionization fraction. We find that, with our source modeling, the average Lyman-$\\alpha$ coupling coefficient $x_\\alpha$ reaches 1 for small values of the ionization fraction (a few $10^{-3}$). However full coupling ($x_\\alpha \\gg 1$) is achieved only for ionization fractions greater than $10^{-1}$. This is in line with expectations. Next we plotted the average spin temperature and differential brightness temperature.  It was shown that the results are very different depending on how the average is computed~: directly on the spin$\\backslash$brightness temperature of the particles, or from the volume-averaged values of the gas kinetic temperature and coupling coefficients. The latter is often used, but the former is more relevant. The direct average tends to erase the signal as regions in absorption have a much stronger signal than regions in emission. Maps of the $21\\,\\rm{cm}$ emission confirm that the different assumptions made in computing the spin temperature lead to very different levels of the signal. The $T_S \\gg T_{CMB}$ assumption is not relevant at all for our choice of source modeling. It would be more relevant at not-too-early redshifts if we included X-ray sources. As expected the full coupling assumption, $x_\\alpha = \\infty$, is reasonable except in the early phase of reionization. In this early phase, we illustrate how computing the local values of $x_\\alpha$ produces large fluctuations in the brightness temperature compared to using a redshift-dependent uniform value.  Finally  the 3D dimensional power spectrum of the differential brightness temperature was plotted, with different assumptions in computing the spin temperature, at various stages of reionization and for the two box sizes. The first important (and expected) result is that including the possibility of a signal seen in absorption adds a large amount of amplitude to the power spectrum. This obviously depends on the choice of the source model. Then, in the early phase of reionization ($\\langle x_H \\rangle < 0.1$), using the local value of the Lyman-$\\alpha$ coupling changes the spectrum by transferring power from large scales to small scales. We reproduce the behaviour shown by Mellema et al. (\\cite{Mellema06b}) and Lidz et al. (\\cite{Lidz07b})~: a maximum of the amplitude at $\\langle x_H \\rangle=0.5$ for the scales around $10\\,h^{-1}\\,\\rm{Mpc}$. However, when  absorption and the exact computation of the Lyman-$\\alpha$ coupling are included, this maximum is shifted to smaller ionization fractions.  As mentioned several times, the first modification to our model that will be investigated is a modification of the source model. Including even a limited number of X-ray sources such as quasars would more efficiently preheat the regions outside of the ionization front, reducing the strength of the absorption regime, or even turning it to emission. This will be investigated  in a forthcoming paper.  Another interesting development of this work is to include the anisotropic effect of the peculiar velocities on the $21\\,\\rm{cm}$ emission (Barkana \\& Loeb \\cite{Barkana05a}). Using the anisotropy, it is possible to separate the contribution of the velocities to the $21\\,\\rm{cm}$ fluctuations from the other contributions and to derive constraints on the cosmological parameters. Since this effect is important mostly in the early reionization phase and on small scales, it is a logical follow-up of including the full modeling of the Lyman-$\\alpha$ coupling.  The implication of this work for the design of the future radio interferometers is simply to emphasize the potential for the detection of the signal during the early phase of reionization when the fluctuations in the signal are strong. Unfortunately, translating \"early phase\" into a numerical value for the redshift depends crucially on the star formation history at these high redshifts, of which very little is known. From the value of the Thomson scattering optical depth, it is likely, however, that this early phase occurs at $z > 12$."
    },
    "0808/0808.1262_arXiv.txt": {
        "abstract": "We study a particular nonlinear dispersion relation $\\omega_p(k_p)$ --- a series expansion in the physical wavenumber $k_p$ --- for modeling first-order corrections in the equation of motion of a test scalar field in a de Sitter spacetime from trans-Planckian physics in cosmology. Using both a numerical approach and a semianalytical one, we show that the WKB approximation previously adopted in the literature should be used with caution, since it holds only when the comoving wavenumber $k\\gg aH$. We determine the amplitude and behavior of the corrections on the power spectrum for this test field. Furthermore, we consider also a more realistic model of inflation, the power-law model, using only a numerical approach to determine the corrections on the power spectrum. ",
        "introduction": "Inflationary models provide answers to many problems in standard big bang cosmology, in particular the origin of density fluctuations and the spectrum of cosmic microwave background anisotropies. The basis of the whole mechanism is the stretching of quantum fluctuations generated at sub-Hubble scales due to the exponential expansion of the spacetime during inflation. This model, however, has a serious ``problem'': if we consider a plain scalar-field-driven inflationary model --- say, chaotic inflation --- the period of inflation lasts so long that the wavelengths of the fluctuations which at present correspond to cosmological scales were sub-Planckian at the begining of the inflationary phase. Therefore, the evolution of fluctuations at such scales is supposed to follow different rules from those provided by the standard theory of cosmological perturbations --- which is based on quantum field theory and general relativity. The set of rules expected to hold above this energy scale is the so-called trans-Planckian physics (TPP from now on). The TPP could lead to deviations from standard predictions on the cosmic microwave background radiation (CMBR), which probes the scales we mentioned before. The question to be asked is ``whether the predictions of the standard cosmology are insensitive to effects of TPP''. This is the precise statement of the trans-Planckian problem \\cite{rhb99,first}.  In the present paper we adopt, to mimic TPP, a series-expansion expression for the nonlinear dispersion relation $\\omega_p$ (Eq.~(\\ref{dr}) below), first suggested by \\cite{w_phys}, to modify the equation of motion of a test scalar field in a de Sitter spacetime. We argue that, in this framework, the WKB approximation is valid only for $k \\gg a H$ (or $k_p \\equiv k/a \\gg H$ with $k_p$ the physical wavenumber and $H$ the Hubble factor). As a consequence,  the perturbative approach based on the WKB approximation, used in \\cite{rhb} to tackle this issue, should be used with caution.  The outline of the paper is the following. In Sec.~\\ref{std} we present the standard approach for first-order cosmological perturbations and then we restrict ourselves to the case of a test scalar field in a de Sitter spacetime. We also introduce the WKB approximation in this section. In Sec.~\\ref{tp} we adopt a particular way to approach TPP: a modification of the equation of motion due to a nonlinear dispersion relation. The initial conditions for the test scalar field are set in Sec.~\\ref{ic}. In the following sections we present a numerical calculation, a 3-piecewise approximation (adopted in Ref.~\\cite{rhb}) and a semianalytical approach to solve the problem. In Sec. \\ref{power} we investigate a realistic model of inflation, namely power-law inflation, using the same nonlinear dispersion relation. We then conclude in Sec. \\ref{Conc}. ",
        "conclusions": "\\label{Conc}  In this paper we have shown that the nonlinear d.r. (\\ref{dr}) presents nonlinear results, i.e, which are not shown in the perturbative approaches used in previous papers in the subject. In particular, the correction $C_{\\alpha\\beta}$ can be almost as large as 10 depending on the value of the parameters $\\alpha$, $\\beta$. Therefore, it can hardly be considered a perturbation.  Our results agree, in part, with the literature on the effects on the power spectrum obtained using different approachs to mimic TPP. In fact, also using different approachs, the effects are always of the order of $H/M_{pl}$ (see \\cite{rhb} and references cited therein) which in our case is ${\\cal O} (1)$.   We have compared previously used approximations and introduced a new one, which allowed a semianalytical calculation.  The fundamental point behind such a large correction factor lies in the break of the WKB approximation at early (conformal) times. If the d.r. at this moment is highly nonlinear, as it happens here, the corrections are bound to be large \\cite{dan}.  However, from Fig. \\ref{num_exact_D} we can see that for particular value of $z$ and $\\Delta$ we obtain no correction to the power spectrum. Namely, that we could still have no correction in spite of finite $\\alpha$ and $\\beta$. In other words, this nonlinear d.r. yields no correction whatsoever to the power spectrum if the parameters happen to be around those values.  In the last section we have  used a power-law model of inflation as the background in order to understand what features were particular to the de Sitter (exponential) expansion. It is clear that one would always get a scale-free spectrum in a de Sitter background, but what would happen to the spectral index $n_s$ in a different one? We have shown that this particular d.r. (\\ref{dr}) yields a even smaller $n_s$ as compared to power-law inflation with a linear d.r., and therefore, it is strongly disfavored by observational results. On the other hand, this result suggests that another nonlinear d.r. may yield also a different correction (perhaps in the opposite direction) to the value of $n_s$ in a power-law background.  We have not addressed the question on back-reactions which may play an important role, since the corrections to the power spectrum are not perturbative ones. Such implications are beyond the scope of this work and deserve a separate study, which will be published elsewhere.    {\\bf Acknowledgments:} We would like to thank R. Ansari, A. Sarkar and M. Schulze for useful discussions and Professor R.H. Brandenberger for useful correspondence. S.E.J. thanks Professor Duval for her help on the Heun functions and acknowledges financial support from ICTP and from CNPq."
    },
    "0808/0808.3117_arXiv.txt": {
        "abstract": "{HH~30 is an edge-on disk around a young stellar object. Previous imaging with the {\\itshape Hubble Space Telescope} has show morphological variability that is possibly related to the rotation of the star or the disk. We report the results of two terrestrial observing campaigns to monitor the integrated magnitude of HH~30. We use the Lomb-Scargle periodogram to look for periodic modulation with periods between 2 days and almost 90 days in these two data sets and in a third, previously published, data set. We develop a method to deal with short-term correlations in the data. Our results indicate that none of the data sets shows evidence for significant periodic photometric modulation.}  \\resumen{HH~30 es un disco visto casi de canto alrededor de un objeto estelar joven. Im\u00e1genes previas del {\\itshape Hubble Space Telescope} muestran una variabilidad morfol\u00f3gica que posiblemente est\u00e9 relacionada con la rotaci\u00f3n de la estrella o el disco. Reportamos los resultados de dos campa\u00f1as observacionales realizadas con un telescopio terrestre para monitorear la magnitud integrada de HH~30. Usamos el periodograma de Lomb-Scargle para buscar modulaciones peri\u00f3dicas con periodos entre 2 y casi 90 d\u00edas en estos dos conjuntos de datos y en un tercer conjunto de datos previamente publicado. Desarrollamos un m\u00e9todo para mitigar los efectos de las  correlaciones de periodo corto en los datos. Nuestros resultados indican que ninguno de los conjuntos de datos muestra evidencia de una modulaci\u00f3n peri\u00f3dica en su photometr\u00eda.}  \\addkeyword{accretion, accretion disks} \\addkeyword{circumstellar matter} \\addkeyword{stars: individual (HH~30)} \\addkeyword{stars: pre-main sequence}  \\begin{document} ",
        "introduction": "High-resolution images show that HH~30 is a compact bipolar reflection nebula bisected by a dark lane (Burrows et al.\\ 1996). Its location in the L1551 molecular cloud and similarity to the model images of Whitney \\& Hartman (1992) led immediately to the conclusion that HH~30 is an optically-thick circumstellar disk seen almost edge-on around a young stellar object.  An interesting aspect of HH~30 is the prominent morphological variability (Burrows et al.\\ 1996; Stapelfeldt et al.\\ 1999; Cotera et al. 2001; Watson \\& Stapelfeldt 2007). This variability includes changes in the contrast between the brighter and fainter nebulae over a range of more than one magnitude; changes in the lateral contrast between the two sides of the brighter nebula over a range of more than one magnitude; and changes in the lateral contrast between the two sides of the fainter nebula over a range of about half a magnitude. It appears that the central source is acting as a lighthouse, preferentially illuminating different parts of the disk.  Two mechanisms have been suggested for the lighthouse. Wood \\& Whitney (1998) suggested non-axisymmetric stellar accretion hot-spots. Stapelfeldt et al.\\ (1999) suggested voids or clumps in the inner disk. AA Tau seems to be a prototype for both mechanisms, apparently possessing both inclined hot spots and occulting inner-disk warps, both presumably the result of an inclined stellar magnetic dipole (Bouvier et al.\\@ 1999; M\u00e9nard et al.\\@ 2003; O'Sullivan et al.\\@ 2005).  The two mechanisms are likely to be periodic, as they are expected to be tied to stellar rotation and orbital motions. Therefore, there is a hope that we might see a periodic modulation in the integrated photometry of HH~30, as one might expect the nebulae to be observed to be brighter when the lighthouse beam is pointing towards the observer. In this work we report an unsuccessful attempt to detect such a periodic modulation in three data sets. ",
        "conclusions": "\\subsection{No Significant Periodic Photometric Variability}  Our analysis indicates that HH~30 shows photometric variability in $V$, $R$, and $I$ (as previously reported), but that periodograms show no significant evidence for a periodic photometric signal between periods of 2 and 87 days. This result is in disagreement with Wood et el.\\ (2000); we suggest that correctly accounting for short-term correlations explains this difference. Of course, this result does not mean that there is no periodic photometric signal present. Rather, it simply means that any periodic signal must be sufficiently weak that it is hidden in the non-periodic noise.       \\subsection{Origin of the Photometric Variability}  The large amplitude of the variability in $I$ (0.8, 1.1, 1.2, and 1.4 magnitudes in data sets 1, 2, and 3 and in the observations reported by Watson \\& Stapelfeldt (2007) along with the lack of a detected period suggests that the photometric variability in HH~30 is related to Type II variability seen in other young stellar objects (Herbst et al.\\ 1994). This is most common in classical T Tauri stars; Lamm et al.\\ (2004) found that 61\\% of the stars in NGC~2264 show this sort of irregular variability whereas only 31\\% show significant periodic variability. Type II variability is thought to be caused by a variable accretion luminosity. This is consistent with the presence of strong collimated jets in HH~30 (Mundt et al.\\ 1990; Burrows et al.\\ 1996; Ray et al.\\ 1996).  \\subsection{Simultaneous HST Imaging}  Watson \\& Stapelfeldt (2007) observed HH~30 with the WFPC2 camera of the {\\it Hubble Space Telescope} on 1999 February 3, coincidentally during the period in which data set 1 was obtained. At this epoch, HH~30 showed a strong lateral asymmetry in the upper nebula. The photometry of data set 1 shows that the magnitude of HH~30 was close to the minimum of its range at this time but rose to the maximum six days later. However, in the absence of evidence for a periodic photometric variability, we are not sure of the significance of these events."
    },
    "0808/0808.1330_arXiv.txt": {
        "abstract": "We show that density-weighted moments of the dissipation rate, $\\epsilon_l$, averaged over a scale $l$, in supersonic turbulence can be successfully explained by the She and L\\'{e}v\\^{e}que model [Phys.\\ Rev.\\ Lett.\\ {\\bf 72}, 336 (1994)]. A general method is developed to measure the two parameters of the model, $\\gamma$ and $d$, based directly on their physical interpretations as the scaling exponent of the dissipation rate in the most intermittent structures ($\\gamma$) and the dimension of the structures ($d$). We find that the best-fit parameters ($\\gamma=0.71$ and $d=1.90$) derived from the $\\epsilon_l$ scalings in a simulation of supersonic turbulence at Mach 6 agree with their direct measurements, confirming the validity of the model in supersonic turbulence. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.1106_arXiv.txt": {
        "abstract": "Cosmological simulations of galaxy clusters typically find that the weight of a cluster at a given radius is not balanced entirely by the thermal gas pressure of the hot ICM, with theoretical studies emphasizing the role of random turbulent motions to provide the necessary additional pressure support. Using a set of high-resolution, hydrodynamical simulations of galaxy clusters that include radiative cooling and star formation and are formed in a Cold Dark Matter universe, we find instead that in the most relaxed clusters rotational support exceeds that from random turbulent motions for radii, $0.1- 0.5\\, r_{500}$, while at larger radii, out to $0.8r_{500}$, they remain comparable.  We also find the X-ray images of the ICM flatten prominently over a wide radial range, $0.1-0.4\\, r_{500}$.  When compared to the average ellipticity profile of the observed X-ray images computed for 9 relaxed nearby clusters, we find that the observed clusters are much rounder than the relaxed CDM clusters within $\\approx 0.4r_{500}$.  Moreover, while the observed clusters display an average ellipticity profile that does not vary significantly with radius, the ellipticity of the relaxed CDM clusters declines markedly with increasing radius, suggesting that the ICM of the observed clusters rotates less rapidly than that of the relaxed CDM clusters out to $\\approx 0.6r_{500}$. When these results are compared to those obtained from a simulation without radiative cooling, we find a cluster ellipticity profile in much better agreement with the observations, implying that over-cooling has a substantial impact on the gas dynamics and morphology out to larger radii than previously recognized. It also suggests that the 10\\%-20\\% systematic errors from non-thermal gas pressure support reported for simulated cluster masses, obtained from fitting simulated X-ray data over large radial ranges within $r_{500}$, may need to be revised downward. These results demonstrate the utility of X-ray ellipticity profiles as a probe of ICM rotation and over-cooling which should be used to constrain future cosmological cluster simulations. ",
        "introduction": "\\label{intro}  Clusters of galaxies, the largest virialized structures in the universe, play a critical role in understanding the formation and evolution of large scale structure, and determining fundamental cosmological parameters. It is now well-known that the dark matter, the dominant component, makes up about 80\\% of the total gravitational mass of a typical galaxy cluster. The baryonic matter in galaxies accounts for about 3 --- 5\\% of the total mass. The remaining mass is made up by hot gas distributed between galaxies, with temperatures over $10^7$ K. This hot gas, called the ``intracluster medium'' (ICM), emits X-rays through thermal bremsstrahlung radiation, and makes galaxy clusters the second most luminous X-ray sources in the universe, after quasars. In the era of precision cosmology, it becomes imperative to measure accurately cluster properties, such as its gravitational mass, to a precision of a few percent.  The X-ray measurement of the density and temperature profiles of the hot ICM provides one of the most accurate methods to determine the cluster total gravitational mass. For clusters that have not been disturbed recently by a major merger, the hot ICM should obey hydrostatic equilibrium to a good approximation, representing a balance between the gravitational force and thermal gas pressure.  For the special case of spherical symmetry the equation of hydrostatic equilibrium may be written, \\begin{equation} \\frac{dP_g}{dr} = -\\rho_g\\frac{GM(r)}{r^2}, \\label{eqn.he} \\end{equation} where $\\rho_g$ is the gas density, $P_g$ is the thermal gas pressure, and $M(r)$ is the mass enclosed within radius $r$. Applications of the X-ray method to constrain cosmological parameters include, e.g., measurements of the gas mass fraction (Allen et al.~2007, and references therein), virial mass function (e.g., Stanek et al.~2006, and references therein) and the relationship between the virial concentration and mass (Buote et al.~2007).  It is necessary to use cosmological hydrodynamical N-body simulations to test the key assumptions underlying equation [\\ref{eqn.he}]; i.e., spherical symmetry and negligible non-thermal pressure support of the hot ICM (e.g., Evrard et al.~1996; Thomas et al.~1998).  Early studies showed that it is possible to recover the cluster mass within a factor of $\\sim 2$ with X-ray telescopes such as {\\sl EINSTEIN} and {\\sl EXOSAT} (e.g., Tsai et al.~1994), where isothermal gas had to be assumed in the mock X-ray analysis. Most recent studies, which involve mock observations with the most advanced telescopes such as {\\sl Chandra}, allow the radial gas temperature profile to be measured, and use simulations that include a wider range of physical processes. They find that the global cluster mass, typically computed within a radius of fixed over-density (e.g., $r_{500}$), can be recovered more accurately, with systematic errors typically ranging from 5\\% to 20\\% (see, e.g., Nagai et al.~2007).  \\begin{figure*}[t] \\begin{center} \\includegraphics[angle=0]{f1} \\end{center} \\caption{Top panel: mock {\\sl Chandra} X-ray flux maps of CL7 from $x$ projection (left), $y$ projection (middle), and $z$ projection (right). Bottom panel: slice of the velocity field in the center of cluster CL7, $y-z$ plane (left), $x-z$ plane (middle), $x-y$ plane (right). The bottom color bar is for the velocity field, and in units of $\\rm km\\ s^{-1}$. Each box has a size of $1h^{-1}\\times1h^{-1}\\rm\\ Mpc^2$.} \\label{image} \\end{figure*}  \\begin{figure*}[t] \\begin{center} \\includegraphics[angle=0]{f2} \\end{center} \\caption{Top panel: mock {\\sl Chandra} X-ray flux maps of CL11 from $x$ projection (left), $y$ projection (middle), and $z$ projection (right). Bottom panel: slice of the velocity field in the center of cluster CL11, $y-z$ plane (left), $x-z$ plane (middle), $x-y$ plane (right). The bottom color bar is for the velocity field, and in units of $\\rm km\\ s^{-1}$. Each box has a size of $1h^{-1}\\times1h^{-1}\\rm\\ Mpc^2$.} \\label{image11} \\end{figure*}  Over the past decade simulators have emphasized how much additional pressure support from random turbulent gas motion contributes to systematic errors in the X-ray method (e.g., Bryan \\& Norman~1998; Kay et al.~2004; Faltenbacher et al.\\ 2005; Dolag et al.~2005; Rasia et al.~2006; Hallman et al.~2006; Vazza et al.~2006; Nagai et al.~2007; Brunetti \\& Lazarian~2007).  It has also been noticed in such high-resolution, hydrodynamic simulations that while the velocity dispersion tensor of the ICM is approximately isotropic in the outskirts of clusters, it becomes increasingly tangential at smaller radii, especially for the most relaxed systems (Rasia, Tormen, \\& Moscardini 2004; Lau et al.~2008, in preparation); Recently Lau et al.\\ (2008) have shown the ICM velocity dispersion becomes increasingly tangentially anisotropic as one moves inward from $r_{500}$, such that for their relaxed clusters the anisotropy parameter falls to $\\beta\\approx -0.3$ near $0.2r_{500}$.  These authors consider the non-thermal pressure support of this random turbulent motion (both radial and tangential) on estimates of $M_{500}$ assuming hydrostatic equilibrium. However, rotational support of the gas, and its observable signatures, is not addressed.  In this paper, we show that for the relaxed clusters studied by Lau et al. (2008; Nagai et al.\\ 2007) support of the ICM from rotational and streaming motions is comparable to the support from the random turbulent pressure out to $\\approx 0.8 r_{500}$.  While the overall magnitude of the rotational motion ($\\sim$ a few hundred $\\rm km\\ s^{-1}$) is not large enough to be detected directly through Doppler shifts of emission lines in X-ray spectra, even with the most advanced X-ray telescopes we have today (e.g., Sunyaev et al.~2003; Inogamov \\& Sunyaev~2003; Schuecker et al.~2004;Br{\\\"u}ggen et al.~2005; Pawl et al.~2005; Chepurnov \\& Lazarian~2006; but see Dupke \\& Bregman~2006 for a recent mmeasurement), simulated clusters with large-scale rotation are significantly flat, and this translates to observable large ellipticities of the X-ray isophotes. By comparing the ellipticities of the X-ray isophotes of the relaxed simulated clusters to those of nine observed clusters, we show that the observed clusters are, on average, much rounder and have a distinctly different radial variation in ellipticity.  This demonstrates the utility of X-ray ellipticity profiles as a constraint for future cosmological cluster simulations.  This paper is organized as follows. In \\S2 we study the importance of ICM rotation in the 16 simulated clusters of Lau et al. (2008; Nagai et al.\\ 2007), focusing our discussion using the examples of one relaxed and one disturbed cluster.  In \\S3 we present ellipticity profiles of nine clusters obtained from X-ray observations with {\\sl Chandra} and {\\sl ROSAT} and compare them to the simulated clusters. The last section is devoted to summary and discussion. ",
        "conclusions": "Over the past decade theoretical studies of cosmological hydrodynamical simulations have emphasized the important role of random turbulent motions in providing pressure support of the ICM in galaxy clusters. Using the set of high-resolution CDM simulations of Nagai et al.\\ (2007) we have shown instead that rotational support of the cluster weight exceeds that of random turbulence from $0.1-0.5\\, r_{500}$ and remains comparable out to $\\approx 0.8r_{500}$ for the clusters classified as the most relaxed in the simulations.  When we compared the average ellipticity profile of the relaxed CDM clusters to that obtained for a set of 9 real nearby clusters observed by {\\sl Chandra} and {\\sl ROSAT}, we found the simulated CDM clusters are significantly flatter within $\\approx 0.4r_{500}$ and the profile shapes differ interior to $\\approx 0.6r_{500}$. By comparing simulations with and without radiative cooling and star formation, we conclude that this flatness is mainly caused by gas over-cooling and large-scale rotation.  The substantial ICM rotation present in the relaxed simulated CDM clusters very likely indicates that a classical ``rotating cooling flow'' (Nulsen, Stewart, \\& Fabian 1984; Kley \\& Mathews 1994; Brighenti \\& Mathews 1996) operates in those systems, implying large mass deposition rates which are not observed.  At very small radii ($<0.1 r_{500}$) it is well-known that cosmological hydrodynamical simulations fail to reproduce observed ICM properties because of over-cooling, but since the ICM rotation we have discussed extends to much larger radii, the over-cooling problem apparently does as well.  We stress that the Nagai et al.\\ (2007) simulations are not unusual in regards to displaying signatures of ICM rotation. Recently, Jeltema et al.\\ (2008) compare X-ray images of simulated and observed clusters using power ratios (Buote \\& Tsai 1995). They find that $P_2/P_0$, which is similar to ellipticity, computed within a 0.5~Mpc aperture for their relaxed simulated CDM clusters is slightly larger than for the observed clusters. They also find that distribution of $P_3/P_0$, which is sensitive to asymmetry but not ellipticity, is consistent for the simulated and observed clusters, implying that it is very unlikely that the $P_2/P_0$ (and hence ellipticity) discrepancy can be attributed to the chosen cosmology.  Since 0.5~Mpc corresponds roughly to $0.5r_{500}$ for clusters, the relatively weak discrepancy is consistent with our results (see Figure \\ref{rosat}).  Because the unphysical ICM rotation in the relaxed CDM clusters extends out to a large radius ($\\approx 0.6r_{500}$), ICM quantities determined within these regions in the simulations need to be reconsidered. For example, it has been noted by several authors that X-ray measurements of $M(<r_{500})$ for relaxed clusters assuming hydrostatic equilibrium underestimate the actual mass by $5-20\\, \\%$ (e.g., Rasia et al.\\ 2006; Nagai et al.\\ 2007; Burns et al.~2007). But since the X-ray analysis involves fitting gas density and temperature profiles of the ICM over a large radial range interior to $r_{500}$, the region of unphysical ICM rotation is included in the analysis. (Note that excluding large radial regions can strongly bias X-ray mass estimates -- Gastaldello et al.\\ 2007.) While we do not expect these X-ray mass underestimates to change radically for relaxed simulated CDM clusters without unphysical ICM rotation, given that $M_{rot}\\sim M_{turb}$ out to $\\approx 0.8r_{500}$ in the simulated clusters we have studied, it is not unreasonable to expect the mass underestimate to be cut in half (see Section 2.4). This would be of considerable importance to studies using X-ray determinations of cluster masses for precision cosmology.  {\\it Acknowledgments:} We thank Daisuke Nagai for providing the sample of simulated clusters and a careful reading of our manuscript. We also thank the referee for helpful comments. We thank Renyue Cen, Fabio Gastaldello, Andrey Kravtsov, Luca Zappacosta for helpful discussions. We gratefully acknowledge partial support from NASA through {\\sl Chandra} Award Number GO6-7118X, GO6-7120X, GO6-7671X, and NNG04GE76G, issued through the Office of Space Sciences Long-Term Space Astrophysics Program."
    },
    "0808/0808.3045_arXiv.txt": {
        "abstract": "{There are many unknowns in the formation of subdwarf B stars. Different formation channels are considered to be possible and to lead to a variety of helium-burning subdwarfs. All seismic models to date, however, assume that a subdwarf B star is a post-helium-flash-core surrounded by a thin inert layer of hydrogen.} {We examine an alternative formation channel, in which the subdwarf B star originates from a massive ($>$ $\\sim$2 M$_{\\odot}$) red giant with a non-degenerate helium-core. Although these subdwarfs may evolve through the same region of the $\\log g-T_{\\rm eff}$ diagram as the canonical post-flash subdwarfs, their interior structure is rather different. We examine how this difference affects their pulsation modes and whether it can be observed.} {Using detailed stellar evolution calculations we construct subdwarf B models from both formation channels. The iron accumulation in the driving region due to diffusion, which causes the excitation of the modes, is approximated by a Gaussian function. The pulsation modes and frequencies are calculated with a non-adiabatic pulsation code. } {A detailed comparison of two subdwarf B models from different channels, but with the same $\\log g$ and $T_{\\rm eff}$, shows that their mode excitation is different. The excited frequencies are lower for the post-flash than for the post-non-degenerate subdwarf B star. This is mainly due to the differing chemical composition of the stellar envelope. A more general comparison between two grids of models shows that the excited frequencies of most post-non-degenerate subdwarfs cannot be well-matched with the frequencies of post-flash subdwarfs. In the rare event that an acceptable seismic match is found, additional information, such as mode identification and $\\log g$ and $T_{\\rm eff}$ determinations, allows us to distinguish between the two formation channels. } {} ",
        "introduction": "Commonly, subdwarf B (sdB) stars are identified as extreme horizontal branch (EHB) stars, and they are believed to be post-He-core-flash products with core masses $\\sim$0.5 M$_{\\odot}$ surrounded by a very thin inert H-envelope \\citep{heber1986,saffer1994}. From a single stellar evolution point of view, this can be explained by enhanced mass loss of stars close to He-ignition with very lightly bound envelopes \\citep{d'cruz1996}, i.e. stars with degenerate cores near the tip of the red giant branch (RGB). However, as they are frequently observed in binaries (e.g.~\\citealt{allard1994,morales2006}), binary interactions most likely play an important role in their formation. \\citet{han2002} explored the main binary evolution channels that can produce sdB stars: common-envelope ejection (CEE), stable Roche lobe overflow (RLOF), and helium white dwarf mergers. They found that the sdB mass distribution may be much broader than previously thought, $0.3-0.8$ M$_{\\odot}$ instead of $0.4-0.5$ M$_{\\odot}$. The sdB stars with non-canonical masses follow from mergers or massive ($>$ $\\sim$2 M$_{\\odot}$) progenitors that ignite helium quiescently, where the latter can be a subchannel of either the CEE channel or the stable RLOF channel. Binary population synthesis shows that the massive progenitors do not contribute significantly to the sdB population \\citep{han2003}. But one should keep in mind that it is asumed in such studies that CE evolution is described by the $\\alpha$ formalism, i.e. that the CE ejection is driven by the orbital energy. Because the physics of the CE phase is poorly understood, other scenarios should not be excluded a priori. For example, the $\\gamma$-formalism proposed by \\citet{nelemans2000}, based on the angular momentum equation rather than the energy equation, provides an alternative description. In this case, the massive red giants cannot be ruled out as possible progenitors of post-CE sdB stars \\citep{hu2007}. We therefore want to explore the possibility of this neglected class of progenitors in a different manner, by using the seismic properties that have been observed in some sdB stars.  Although the post-flash and the post-non-degenerate sdB stars can appear in the same $\\log g-T_{\\rm{eff}}$ region, their interior structure is quite different. In particular, the chemical composition profiles differ greatly depending on whether helium ignited in a flash or quiescently. For example, the canonical post-He-flash sdB star has a very narrow He$-$H transition zone, while the sdB star created from a more massive progentior has a much broader H-profile. This is a direct result of the differing chemical compositions between low-mass and high-mass stars on the RGB, owing to the different convective regions during the main-sequence and RGB evolution. We examine whether this difference in the interior structure will result in observable differences in the pulsation modes.  The sdB pulsators consist of two classes, the short-period variable EC 14026 stars \\citep{kilkenny1997}, and the long-period variable PG 1716 stars \\citep{green2003}. The rapid oscillations in EC 14026 stars are interpreted in terms of low-order $p$-modes \\citep{charpinet1996}, driven by the $\\kappa$-mechanism operating in the iron opacity bump. The same mechanism has been shown to excite long-period, high-order $g$-modes in the cooler models \\citep{fontaine2003}. The local iron enhancement necessary in the driving region around $\\log T\\approx 5.3$ is due to the competing diffusion processes of radiative levitation and gravitational settling. It is well-known that the opacities play an important role in the study of the pulsations. \\citet{seaton2004} showed that the iron opacity bump is situated at slightly higher temperatures using OP opacities \\citep{seaton1994,badnell2003} compared with OPAL opacities \\citep{iglesias1996}. \\citet{jeffery2006} found that, using OP opacities and nickel enhancement in addition to iron, the theoretical instability strip of $g$-mode sdB oscillators is more consistent with observations. For our purposes it is sufficient to use OPAL opacities and iron enhancement, since we are interested in the relative differences between two types of sdB stars. We acknowledge the importance of including the effect of OP opacities and nickel enhancement in further detailed studies.  The details of the computations are given in \\S \\ref{computations}. The results are presented in \\S \\ref{results}. In \\S \\ref{structure}, we compare the detailed physical characteristics of two reference models with different formation histories. In \\S \\ref{grid} we compare the frequency characteristics globally between two grids of models. The results and conclusions are discussed in \\S \\ref{discussion}. ",
        "conclusions": "\\label{discussion}  We studied the so far neglected, post-non-degenerate sdB stars and compared their physical and seismic characteristics with those of canonical post-flash sdB stars, both formed in the CEE channel. The results presented here are a first step in distinguishing these two kinds of sdB stars on the basis of their observed oscillation character, which is necessary if seismic modelling is to achieve reliable mass determination. Furthermore, the observation of a post-non-degenerate sdB star in a post-CE binary would give strong constraints on the CE evolution. We plan to continue such investigations with an application to the sdB pulsator in the post-CE, eclipsing binary \\object{PG\\,1336$-$018} which we started in \\citet{hu2007} and \\citet{vuckovic2007}.  We find that, in principle, a post-non-degenerate sdB star may appear as an EC 14026 star with similar pulsation frequencies as the canonical post-He-flash sdB star, although it is not likely. Additional observables, such as spectroscopic $\\log g$ and $T_{\\rm eff}$ determinations and/or empirical mode identification from observables enable us to distinguish the two types of sdB stars more decisively. The frequency range of the unstable modes is also an important discriminator between the two formation channels. In general, for the same $\\log g$ and $T_{\\rm eff}$ values, the excited frequencies of the post-non-degenerate sdB star are higher than the excited frequencies of the post-flash star. This is a direct result of the differing interior structures. Thus, special attention must be paid when observed frequencies are matched with theoretically predicted ones of modes that are not excited.  Up to now, there have not been any evolutionary models of sdB stars available that include the coupling between diffusion and evolution consistently. This is a deficiency, since iron accumulation due to radiative levitation is responsible for the pulsational instability in these stars \\citep{charpinet1996}. Also, it has been shown by \\citet{fontaine2006} that the iron accumulation changes the frequencies significantly. In our study, we have parametrized the iron accumulation, so that we can, at least in an approximative manner, simultaneously take into account the effects of iron enhancement and evolution on the pulsation modes.  Here we have not considered the influence of the other diffusive processes, i.e.~diffusion due to gradients of pressure, temperature, and concentration. To a certain extent this can affect our results, because one of the main differences between the two types of sdB stars is the chemical composition of the stellar envelope. Specifically, we find in the envelope of the post-non-degenerate sdB star an H-mass fraction of $X=0.18$, while the post-flash sdB star has $X=0.66$ there. Normally, it is assumed that sdB stars have H-rich or even pure H-envelopes, caused by gravitational settling. While this is true for the outermost layers, diffusion is not expected to work efficiently at depths $\\log q\\gtrsim-3$ \\citep{richard2002, michaud2007}. Since the envelopes of the post-non-degenerate sdB stars extend to $\\log q\\gtrsim -2$  (i.e.~$T\\gtrsim10^7$ K), we do not expect diffusion to wash away all the qualitative differences in the chemical profiles, although the differences may be less pronounced. Diffusion, however, will significantly change the surface abundances of our models and likely will bring them into agreement with the observed values. Spectroscopic line profile analysis has shown that the majority of sdB is He-deficient, and only a few are He-rich \\citep{edelmann2003,lisker2004}. Stellar evolution models that include diffusion coupled to reliable atmosphere models are needed to assess whether the two different formation channels will be distinguishable via a spectroscopic abundance analysis. We are currently computing such evolutionary sdB models including diffusion due to gradients of pressure, temperature, and concentration. Our preliminary results indeed agree with our expectations, i.e.~the H-surface abundance increases on a very short timescale, but the chemical profiles at deeper layers are not affected. The pulsational properties of these improved models will be discussed in detail in a forthcoming paper.  We have made a modest grid of models that is sufficient for our comparative study. Detailed seismic modelling of an observed star, however, will require a finer grid. For now, we have chosen not to make sdB models above 0.47 M$_{\\odot}$, since this is the maximum mass the degenerate He-core of a red giant with $Z=0.02$ can have before experiencing the He-flash. A metallicity of $Z=0.004$ allows the He-core to grow up to 0.48 M$_{\\odot}$ on the RGB. However, we find that, in order to excite modes in these low metallicity stars, an iron enhancement greater than a factor 10 is required. This was to be expected, since \\citet{charpinet1996} found unstable pulsation modes for models with uniform $Z\\geq0.04$ in the H-rich envelope. We have therefore not pursued these models further. The question whether post-flash sdB stars can have masses $>0.47$ M$_{\\odot}$ is also closely related to the input physics (e.g.~convective overshooting) and the physics of the He-flash, and needs to be examined further.  In this paper, we have focused on the short-period $p$-mode sdB pulsators. The case of the long-period $g$-mode sdB pulsators is, although challenging from an observational point of view, an additional very interesting theoretical case study.  The $p$-modes only probe the outermost layers, and hence are less affected by the differing composition gradients than the $g$-modes, as they propagate deeper into the star. The long-period sdB pulsators are interpreted as cooler sdB models with much thicker hydrogen-envelopes than the short-period sdB pulsators \\citep{fontaine2003, jeffery2006}. Since the $g$-modes are deep interior modes, full evolutionary models including iron accumulation, as developed here, are required to model these stars. At present, these are not available yet. We are currently developing a similar approach to the one presented here to study the long-period sdB pulsators."
    },
    "0808/0808.2796_arXiv.txt": {
        "abstract": "{Recent works with improved model atmospheres, line formation, atomic and molecular data, and detailed treatment of blends, have resulted in a significant downward revision of the solar oxygen abundance.} {Considering the importance of the Sun as an astrophysical standard and the current conflict of standard solar models using the new solar abundances with helioseismological observations we have performed a new study of the solar oxygen abundance based on the forbidden [\\ion{O}{i}] line at 5577.34 \\AA, not previously considered.} {High-resolution (R $>$ 500 000), high signal-to-noise (S/N $>$ 1000) solar spectra of the [\\ion{O}{i}] 5577.34 \\AA\\ line have been analyzed employing both three-dimensional (3D) and a variety of 1D (spatially and temporally averaged 3D, Holweger \\& M\\\"uller, MARCS and Kurucz models with and without convective overshooting) model atmospheres.} { The oxygen abundance obtained from the [\\ion{O}{i}] 5577.3 \\AA\\ forbidden line is almost insensitive to the input model atmosphere and has a mean value of $\\log \\epsilon_{\\rm O} = 8.71 \\pm 0.02$ ($\\sigma$ from using the different model atmospheres). The total error (0.07 dex) is dominated by uncertainties in the log $gf$ value (0.03 dex), apparent line variation (0.04 dex) and uncertainties in the continuum and line positions (0.05 dex). } { The here derived oxygen abundance is close to the 3D-based estimates from the two other [\\ion{O}{i}] lines at 6300 and 6363\\,\\AA , the permitted \\ion{O}{i} lines and vibrational and rotational OH transitions in the infrared. Our study thus supports a low solar oxygen abundance  ($\\log \\epsilon_{\\rm O} \\approx$ 8.7), independent of the adopted model atmosphere. } ",
        "introduction": "In recent years a significant reduction in the solar oxygen abundance has been proposed from $\\log \\epsilon_{\\rm O} \\equiv \\log (N_{\\rm O}/N_{\\rm H}) + 12 = 8.93$ (Lambert 1978; Sauval et al. 1984; Grevesse et al. 1984; Anders \\& Grevesse 1989) down to  $\\log \\epsilon_{\\rm O} \\approx$ 8.7 or even lower (Allende Prieto et al. 2001, 2004; Holweger 2001; Asplund et al. 2004; Melendez 2004; Socas-Navarro \\& Norton 2007; Caffau et al. 2008; but see Ayres et al. 2006, Ayres 2008, and Centeno \\& Socas-Navarro 2008 for a departing conclusion). This large revision in the oxygen abundance has an important impact in many areas of astrophysics. In particular, solar models computed with the low oxygen abundance do not agree with precise helioseismological measurements (see Basu \\& Antia 2008 and references therein).  If the revised solar oxygen abundance is correct, this would imply that the standard solar structure model may be missing (or have a too simplified treatment of) important physical processes, which would carry over to errors in the calculations of stellar evolution at large. On the other hand, if the fault lies with the new solar photospheric abundance analyses it raises concerns about how well we understand stellar atmospheres and spectral line formation in general with wider implications for Galactic chemical evolution studies relying on accurate stellar abundance measurements. Indeed, the oxygen content in metal-poor stars is currently under debate (see Asplund \\& Garc{\\'{\\i}}a P{\\'e}rez 2001; Mel\\'endez et al. 2001, 2006, and references therein).  Since the Sun is used as a fundamental reference in most areas of astrophysics, it is important to study all spectral features available for abundance analysis. Recent studies have mainly used the  [\\ion{O}{i}]  6300, 6363 \\AA\\ forbidden lines, the \\ion{O}{i} 7777 \\AA\\ triplet, pure rotation OH lines as well as the fundamental vibration-rotation OH lines (Allende Prieto et al. 2001, 2004; Asplund et al. 2004), from which $\\log \\epsilon_{\\rm O} = 8.66 \\pm 0.05$ have been recommended (Asplund et al. 2004). Mel\\'endez (2004) presented a study of the first-overtone OH lines, which are unfortunately very weak for a precise determination of the oxygen abundance, but within the uncertainties these lines also support a low solar oxygen abundance.  Another independent way to estimate the solar oxygen abundance is using the forbidden line at 5577.34 \\AA. This line has been disregarded in the past due to blending with C$_2$ lines (e.g. Altrock 1968; Lambert 1978). Here, we present the first detailed study of the [\\ion{O}{i}] 5577 \\AA\\ line, carefully taking into account the blends by C$_2$ lines. The advantage of this feature is that the C$_2$ blends introduce a large asymmetry in the profile, allowing thus to readily constrain their contribution. Another advantage is that there are other nearby C$_2$ lines that can be used to calibrate the blending C$_2$ lines, and therefore to determine a reliable solar oxygen abundance.  \\begin{figure} \\includegraphics[width=\\hsize]{o5575_5580.eps} \\caption{Observed solar spectrum (Delbouille et al. 1973) between 5575 and 5580\\,\\AA\\ (red circles) and the best fit employing the $<$3D$>$ model (solid blue line). The locations of the most prominent C$_2$ lines used to estimate the C abundance are denoted with vertical lines, while the position of the [\\ion{O}{i}] line is marked with an arrow. The resulting solar C abundance is  $\\log \\epsilon_{\\rm C} = 8.42$ with which the 5577.3\\,\\AA\\ feature implies a solar O abundance of $\\log \\epsilon_{\\rm O} = 8.73$. } \\label{f:5575-5580} \\end{figure}   \\begin{figure} \\includegraphics[width=\\hsize]{o5577small.eps} \\includegraphics[width=\\hsize]{onoiron.eps} \\caption{{\\em Upper panel:} An enlargement of the observed (red circles, Delbouille et al. 1973) and $<$3D$>$-based theoretical profiles (solid blue line) shown in Fig. \\ref{f:5575-5580}. The dotted lines show the separate contribution of [\\ion{O}{i}] and other lines labeled accordingly. The 5577.54 {\\AA} P$_3$25 C$_2$ line is blended with an unknown feature, assumed here to be \\ion{Fe}{I}. {\\em Lower panel:} Same as above but assuming instead that the red feature is entirely due to P$_3$25 C$_2$. Applying this 0.03\\,dex higher C abundance to the [\\ion{O}{i}]+C$_2$ blend implies a smaller left-over contribution for [\\ion{O}{i}] and thus a 0.13\\,dex lower O abundance at the expense of a significantly poorer profile fit. } \\label{f:oi} \\end{figure} ",
        "conclusions": "A low oxygen abundance is obtained from the [\\ion{O}{i}] 5577.3 \\AA\\ line, almost independent of the adopted model atmosphere: $\\log \\epsilon_{\\rm O} = 8.71 \\pm 0.02 \\pm 0.07$ (0.02 dex is the $\\sigma$ from using different models, and 0.07 dex is the error due to uncertainties in the the log $gf$ value (0.03 dex), apparent line variation (0.04 dex) and uncertainties in the continuum and line positions (0.05 dex)). This value is close to the results of Asplund et al. (2004) and Mel\\'endez (2004) for other solar oxygen abundance indicators (\\ion{O}{i}, [\\ion{O}{i}] and OH lines). Including those abundances and employing six different 3D and 1D model atmospheres, we estimate a solar O abundance of $\\log \\epsilon_{\\rm O} = 8.70$ ($\\sigma$ = 0.02) with only a small sensitivity to the employed model atmosphere.  The low solar oxygen abundance that we propose here as well as the downward revision of other solar chemical abundances (Asplund  2000, 2005, 2007; Asplund et al. 2000, 2005, 2006) pose a challenge to standard solar models in light of helioseismological observations (e.g. Antia \\& Basu 2005, 2006; Bahcall et al. 2004, 2005, 2006; Basu \\& Antia 2004, 2008; Chaplin et al. 2007; Guzik et al. 2005; Yang \\& Bi 2007)."
    },
    "0808/0808.0270_arXiv.txt": {
        "abstract": "{The physical properties of almost any kind of astronomical object can be derived by fitting synthetic spectra or photometry extracted from theoretical models to observational data.} {We want to develop an automatic procedure to perform this kind of fittings to a relatively large sample of members of a stellar association and apply this methodology to the case of Collinder 69.} {We combine the multiwavelength data of our sources and follow a work-flow to derive the physical parameters of the sources. The key step of the work-flow is performed by a new VO-tool, VOSA. All the steps in this process are done in a VO environment.} {We present this new tool, and provide physical parameters such as T$_{\\rm eff}$, gravity, luminosity, etc. for $\\sim$170 candidate members to Collinder 69, and an upper-limit for the age of this stellar association.} {This kind of studies of star forming regions, clusters, etc. produces a huge amount of data, very tedious to analyse using the traditional methodology. Thus, they are excellent examples where to apply the VO capabilities.} ",
        "introduction": "\\label{intro}  The determination of physical parameters of astronomical objects from observational data is frequently linked with the use of theoretical models as templates. In the case of stars, it is very common to estimate gravities by fitting different model/templates to gravity--sensitive lines (and other features, \\citealt{Decin04}) within the wavelength coverage of the observed spectra. Another method in the same framework is the comparison of the shape of the Spectral Energy Distributions (SEDs hereafter) constructed from observational photometric data with the corresponding synthetic ones derived using large grids of theoretical spectra (see for example \\citealt{Robitaille07}).  The use, -in the traditional way-, of these methodologies can easily become tedious and even unfeasible when applied to large amount of data. Nowadays, astronomers deal with very large databases, including not only empirical/observational entries, but also theoretical ones (for instance, models from di\\-ffe\\-rent groups that need to be combined in order to reproduce the observations).  The Virtual Observatory (VO) philosophy (see Section~\\ref{TVOT}) seems to be a reasonable approach to solve some intrinsic problems generated by this new way of working in astronomy. VO-tools\\footnote{ http://ivoa.net/twiki/bin/IVOA/IvoaApplications} now allow astronomers to deal with their own observational or theoretical databases and compare them (in an efficient way) with others already published in the ``on-line literature''.  Here we present a work-flow to estimate the effective temperature (T$_{\\rm eff}$), gravity and luminosity (and, when comparing with theoretical isochrones and evolutionary tracks, age and mass) of members of an stellar association (in this particular case, Collinder 69, \\citealt{Barrado04}) based on their photometry. This work-flow has been built mainly with steps that are ``VO-compliant''.  In Section~\\ref{SC} we explain the scientific case that caused the development of VOSA\\footnote{http://www.laeff.inta.es/svo/theory/vosa/}, the key VO-tool used in this work-flow. In Section~\\ref{TVOT} we describe the philosophy of the Virtual Observatory and some of its most known applications (tools that we have used in several steps of the process). In that section we also explain all the capabilities of the previously mentioned VO-tool VOSA (developed by the Spanish Virtual Observatory\\footnote{ https://laeff.inta.es/svo/} specially for this work). In Section~\\ref{TM} we explain step by step the work-flow followed to determine the physical parameters and properties of our sources. In Sections ~\\ref{Results} and ~\\ref{C}, we present the results and conclusions of this work, respectively. Finally, in Appendix A, we focus on technical issues concerning the computation of synthetic photometry  and in Appendix B we provide a graphical scheme of the work-flow. ",
        "conclusions": "\\label{C}  We have presented a new VO-tool, VOSA, developed by the Spanish Virtual Observatory team for this particular scientific case (but with a wider applicability range) to easy the process of calculating synthetic photometry and performing $\\chi^{2}$ tests to large sets of fittings as well as to interpolate among collections of isochrones and evolutionary tracks.  When applied to the case of Collinder 69, we have been able to perform the following studies: \\begin{enumerate} \\item We have estimated different physical parameters (effective temperatures, gravity, bolometric luminosity, and in most of the cases, mass and age) for 167 candidate members.% \\item We have independently confirmed the classification from \\citet{Barrado07} of non-members for 16 of the sources of the sample, and we have added a new possible non-member to this list (LOri162). \\item We have derived an upper-limit for the age of 12.3-16 Myr consistent with previous estimations in the literature. \\end{enumerate}    We must note that VOSA is a very efficient tool; as an example, it only takes $\\sim$ 20 minutes to reproduce the whole workflow presented in this paper."
    },
    "0808/0808.0793_arXiv.txt": {
        "abstract": "Using new, high-resolution interferometric observations of the CO and HCN molecules, we directly compare the molecular and ionized components of the interstellar medium in the center of the nearby spiral galaxy IC\\,342, on spatial scales of $\\approx$ 10\\,pc. The morphology of the tracers suggests that the molecular gas flow caused by a large-scale stellar bar has been strongly affected by the mechanical feedback from recent star formation activity within the central 100\\,pc in the nucleus of the galaxy. Possibly, stellar winds and/or supernova shocks originating in the nuclear star cluster have compressed, and likely pushed outward, the infalling molecular gas, thus significantly reducing the gas supply to the central 10\\,pc. Although our analysis currently lacks kinematic confirmation due to the face-on orientation of IC\\,342, the described scenario is supported by the generally observed repetitive nature of star formation in the nuclear star clusters of late-type spiral galaxies. ",
        "introduction": "Any ``activity'' found in a galactic nucleus -- be it due to massive star formation, an accreting super-massive black hole, or a combination of both -- requires the inflow of gas into the central few pc. On scales of kpc, the gas flow is regulated by the response to asymmetries in the gravitational potential, e.g. stellar bar or tidal interactions \\citep{sak99,she05}. Numerical models \\citep[e.g.][]{ath92,sel93} have demonstrated that large-scale stellar bars are especially efficient in moving gas into the central kpc. However, no clear picture is yet available for the gas flow inside a few 100 pc from the nucleus \\citep{gar03,wad04}.  Arguably the simplest galaxies to observationally test these models are the latest-type spirals, because they have bulge-less disks. However, even these objects have rather complex central star formation histories. About 75\\% of late-type disks host distinct nuclear star clusters \\citep{boe02}, most of which have experienced multiple discrete star formation events in their history \\citep{ros06,wal06}. So far, it is unclear whether the repetitive star formation in these nuclear clusters is due to variability of the gravitational potential (e.g. dissolution/formation of a stellar bar), the clumpy nature of the molecular gas (e.g. inflow of discrete giant molecular complexes -- GMCs), or whether nuclear massive star formation itself is disrupting the gas supply.  The Scd spiral \\objectname{IC342} is one of the nearest examples of a late-type spiral harboring a well-studied nuclear star cluster. The nuclear cluster has a stellar mass of $\\sim 10^7\\,\\msun$ and its luminosity is dominated by a recent, 4--30~Myr~old, short-lived star formation event \\citep[after correction for a revised distance of 3.0 Mpc;][]{boe97,boe99,fin07}. Previous studies of the molecular gas in IC342, with single dish and interferometric millimeter instruments \\citep{tur92,tur93,ish90,dow92,mei01}, lack sufficient spatial resolution (58--72~pc) to resolve the structure of the GMCs, thus hampering accurate comparisons with the nuclear star formation. The spiral geometry of the molecular gas in the central 500pc is due to the response of the gas to the large-scale stellar bar \\citep{ish90}. A comprehensive study \\citep{mei05} of the molecular chemistry in IC342 has demonstrated that the molecular gas in the central few hundred pc is subject to a number of different excitation mechanisms. Recent OVRO observations of the $^{12}$CO(2-1) line showed evidence for molecular gas being associated with the nuclear star cluster itself \\citep{sch03}. ",
        "conclusions": ""
    },
    "0808/0808.2105_arXiv.txt": {
        "abstract": "We present the results of EVN+MERLIN VLBI polarization observations of 8 Broad Absorption Line (BAL) quasars at 1.6 GHz, including 4 LoBALs and 4 HiBALs with either steep or flat spectra on VLA scales. Only one steep-spectrum source, J1122+3124, shows two-sided structure on the scale of 2~kpc. The other four steep-spectrum sources and three flat-spectrum sources display either an unresolved image or a core-jet structure on scales of less than three hundred parsecs. In all cases the marginally resolved core is the dominant radio component. Linear polarization in the cores has been detected in the range of a few to 10 percent. Polarization, together with high brightness temperatures (from 2$\\times10^9-5\\times10^{10}$K), suggest a synchrotron origin for the radio emission. There is no apparent difference in the radio morphologies or polarization between low-ionization and high-ionization BAL QSOs nor between flat- and steep-spectrum sources. We discuss the orientation of BAL QSOs with both flat and steep spectra, and consider a possible evolutionary scenario for BAL QSOs. In this scenario, BAL QSOs are probably the young population of radio sources, which are Compact Steep Spectrum or GHz peaked radio source analog at the low end of radio power. ",
        "introduction": "About~$10$~-~$20$\\% of optically selected quasars exhibit Broad Absorption Line (BAL) troughs up to $\\sim 0.1c$ blueward of their corresponding emission lines (Weymann et al. 1991; Weymann 2002; Tolea et al. 2002; Hewett \\& Foltz 2003; Reichard et al. 2003). These BALs are produced via resonant scattering in partially ionized outflows. BAL quasars can be further divided into high- and low-ionization sub-classes (HiBALs and LoBALs). HiBAL quasars are those objects which show BALs only in highly ionized species such as CIV and NV, while LoBAL quasars also show BALs in Mg II or Al III. Two different scenarios have been proposed to explain the BAL phenomena. In the first scenario, the BAL region (BALR) is present with a small covering factor in every QSO. The different appearance of BAL and non-BAL QSOs is solely due to different lines of sight. In a BAL QSO our line of sight intercepts the BALR while it does not in non-BAL QSOs. In this scenario, the fraction of BAL QSOs is interpreted as the covering factor of BALRs. This paradigm has also been generalized to interpret the dichotomy of LoBAL and HiBAL as an orientation effect, in which low-ionization outflow covers small fraction sky with the BAL outflow. This scenario is consistent with a number of statistical properties, such as the similarity between the UV continuum and emission line spectra for both types of QSOs (e.g., Weymann et al. 1991), their similar millimeter and far-infrared luminosities (e.g., Willott et al. 2003), as well as their similar large scale environments (Shen et al. 2008). In the second scenario BALs exist only in a relative short phase of QSO activity, which is very likely in their early phase of evolution, at least for LoBAL QSOs (Briggs, Turnshek \\& Wolfe 1984). Support for the latter scenario includes the excess fraction of LoBAL quasars among infrared selected quasars (Boroson \\& Meyers 1992), and a special locus in black hole mass versus accretion rate space (Boroson 2002).  In the first scenario the BALR is probably located in a preferred direction with respect to the system$^{\\rm '}$s symmetry axis, e.g., in the equatorial or polar direction. Higher polarization in the optical continuum in comparison to non-BAL QSOs suggests that BAL QSOs are seen nearly edge-on (Goodrich \\& Miller 1995; Hines \\& Wills 1995; Cohen et al. 1995; Wang, Wang \\& Wang 2005). Equatorial outflows are also preferred from theoretical considerations for radiatively accelerated winds from an accretion disc or evaporated gas from a putative dusty torus (Punsly 2006).  However, there are indications for both polar and equatorial outflows in radio-loud BAL QSOs.  It is generally believed that the radio jet is aligned with the symmetry axis of the accretion disk. Thus it can be taken as an indicator of the system inclination. According to the unification scheme of radio quasars, a radio-loud quasar would show a flat radio spectrum and a core-dominated morphology when viewed along the radio jet, due to relativistic boosting of the optically thick core. It would appear as a steep-spectrum radio source with double-lobed morphology when viewed side-ways (e.g., Urry \\& Padovani 1995). Thus the radio morphology and spectrum can serve as a surrogate for the inclination of the system: in an equatorial outflow model a radio-loud BAL QSO would appear as a steep-spectrum radio source with double-lobed radio morphology, while the opposite situation will be observed in a polar-outflow model. Becker et al. (2000) found that BAL QSOs in the FIRST Bright Quasar Survey (FBQS) display both steep and flat radio spectra (see also Menou et al. 2001), indicating both polar and equatorial outflows. Other studies suggest a similar result. On the one hand, nearly a dozen double-lobed FR~II quasars were found, mostly from the SDSS and FIRST samples (Wills, Brandt, \\& Loar 1999; Gregg et al. 2000; Brotherton et al. 2002; Zhou et al. 2006; Gregg, Becker \\& de Vries 2006). On the other hand, evidence for polar outflows has also been found in a small number of radio variable BAL QSOs (Zhou et al 2006; Ghosh \\& Punsly 2007), in which the radio variability suggests that the radio jet is beamed towards the observer. The presence of both polar and equatorial outflows can be better interpreted in the context of an evolutionary scenario, rather than the geometrical unification scheme.  Further studies of radio-loud BAL QSOs are necessary in order to fully understand the geometry of outflows in their population. Note that most radio-loud BAL QSOs with steep radio spectra show compact structures on VLA scales, different from classical radio-loud quasars on which the unification scheme is based. Thus it is not clear whether the radio spectrum can be used as an indicator for inclination or not. Powerful FR~II radio BAL QSOs are extremely rare among radio-loud BAL QSOs, and most radio BAL QSOs are only moderately strong in the radio. Morphological studies of these radio intermediate BAL quasars can reveal representative properties of the majority radio population of BAL QSOs. Currently, less than a handful of such BAL QSOs have been studied at sufficiently high angular resolution and they have revealed a complex situation.  In previous EVN observations, Jiang \\& Wang (2003) observed 3 BAL quasars at L band using the phase reference technique. Among these three sources, only one (J1312+2319) has compact two-sided structure, and the other two are still unresolved. The jet orientation might be near the line of sight if the compact component is the jet base in these two unresolved sources. In this case the orientation model may be problematic, at least for these 3 sources.  In this paper, we report our EVN+MERLIN observations at 1.6 GHz of eight BAL quasars, including equal numbers of HiBAL and LoBAL quasars and then give some discussion of the results. $\\S$ 2 presents the observations as well as data reduction and the results are described in $\\S$ 3. The discussion is contained in $\\S$ 4, and in the last section $\\S$ 5 we give the conclusions. Throughout the paper we adopt the spectral index convention $\\rm f_{\\nu}\\propto\\nu^{-\\alpha}$ and a cosmology with $\\rm H_{0}=70 \\rm {~km ~s^ {-1}~Mpc^{-1}}$, $\\rm \\Omega_{M}=0.3$, $\\rm \\Omega_{\\Lambda} = 0.7$. All values of luminosity used in this paper are calculated with our adopted cosmological parameters. ",
        "conclusions": "We present the results of EVN plus MERLIN polarization observations of 8 BAL quasars at 1.6 GHz, including 4 LoBALs and 4 HiBALs with either steep or flat spectra on VLA scales. The main conclusions are summarized as follows: \\begin{itemize} \\item{Only one steep-spectrum source, J1122+3124, shows two-sided structure on the scale of 2~kpc. The other four steep-spectrum sources and three flat-spectrum sources display either an unresolved image or a core-jet structure on scales of less than three hundred parsecs, well within the galaxy size. In all cases, the marginally resolved core is the dominant radio component. Making use of the phase-reference technique, celestial positions are derived from our observations which are more accurate than those from the VLA.}  \\item{Linear polarization has been detected in the core in the range of a few to 10 percent. Polarization, together with high brightness temperatures (from 2$\\times10^9-5\\times10^{10}$K), suggest a synchrotron origin for the radio emission. There is no apparent difference in the radio morphologies or polarization between low-ionization and high-ionization BAL quasars nor between flat- and steep-spectrum sources.}  \\item{We considered compact steep-spectrum or GHz peaked radio source at the low end of radio power as the most likely explanation for these radio BAL QSOs. Therefore, they are probably a population of young radio sources.}  \\end{itemize}"
    },
    "0808/0808.3723_arXiv.txt": {
        "abstract": "Massive stars played a key role in the early evolution of the Universe. They formed with the first halos and started the re-ionisation. It is therefore very important to understand their evolution. In this paper, we describe the strong impact of rotation induced mixing and mass loss at very low $Z$. The strong mixing leads to a significant production of primary \\el{N}{14}, \\el{C}{13} and \\el{Ne}{22}. Mass loss during the red supergiant stage allows the production of Wolf-Rayet stars, type Ib,c supernovae and possibly gamma-ray bursts (GRBs) down to almost $Z=0$ for stars more massive than 60 $M_\\odot$. Galactic chemical evolution models calculated with models of rotating stars better reproduce the early evolution of N/O, C/O and \\el{C}{12}/\\el{C}{13}. We calculated the weak s-process production induced by the primary \\el{Ne}{22} and obtain overproduction factors (relative to the initial composition, $Z=10^{-6}$) between 100-1000 in the mass range 60-90. ",
        "introduction": "Massive stars ($M \\gtrsim 10\\ M_\\odot$) started forming about 400 millions years after the Big Bang and ended the dark ages by re-ionising the Universe. They therefore played a key role in the early evolution of the Universe and it is important to understand the properties and the evolution of the first stellar generations to determine the feedback they had on the formation of the first cosmic structures. It is unfortunately not possible to observe the first massive stars because they died a long time ago but their chemical signature can be observed in low mass halo stars (called EMP stars), which are so old and metal poor that the interstellar medium out of which these halo stars formed are thought to have been enriched by one or a few generations of massive stars. Since the re-ionisation, massive stars have continuously injected kinetic energy (via various types of supernovae) and newly produced chemical elements (by both hydrostatic and explosive burning and s and r processes) into the interstellar medium of their host galaxies. They are thus important players for the chemo-dynamical evolution of galaxies. Most massive stars leave a remnant at their death, either a neutron star or a black hole, which often lead to the formation of pulsars or X-ray binaries.  The evolution of stars is governed by three main parameters, which are the initial mass, metallicity ($Z$) and rotation rate. The evolution is also influenced by the presence of magnetic fields and of a close binary companion. For massive stars around solar metallicity, mass loss plays a crucial role, in some cases removing more than half of the initial mass. Internal mixing, induced mainly by convection and rotation also significantly affect the evolution of stars. The properties of non-rotating low-Z stars are summarised in Sect. 2 of Hirschi et~al. (2008). In short, at low Z, stars are more compact, usually have weaker winds (see however \\cite{Pu08}) and are more massive (Bromm \\& Loeb 2003, \\cite{SOIF06}). The fate of non-rotating massive single stars at low Z is summarised in \\cite{HFWLH03} and several groups have calculated the corresponding stellar yields (\\cite{HW02,CL04,TUN07}). In this paper, after discussing the impact of rotation induced mixing and mass loss at low $Z$, we present the implications for the nucleosynthesis and for galactic chemical evolution in the context of extremely metal poor stars. ",
        "conclusions": "The inclusion of the effects of rotation changes significantly the simple picture in which stellar evolution at low Z is just stellar evolution without mass loss. A strong mixing is induced between the helium and hydrogen burning layers leading to a significant production of primary \\el{N}{14}, \\el{C}{13} and \\el{Ne}{22}. Rotating stellar models also predict a strong mass loss during the RSG stage for stars more massive than 60 $M_\\odot$. These models predict the formation of WR and type Ib/c SNe down to almost Z=0. The chemical composition of the stellar winds is compatible with the CNO abundance observed in the most metal-poor star known to date, HE1327-2326. GCE models including the stellar yields of these rotating star models are able to better reproduce the early evolution of N/O, C/O and \\el{C}{12}/\\el{C}{13} in our galaxy. These new stellar evolution models predict a large neutron release during core He burning and we present here the corresponding s-process production during He burning.  Large surveys of EMP stars (SEGUE, OZ surveys), of GRBs and SNe (Swift and GLAST satellites) and of massive stars (e.~g. VLT FLAMES survey) are underway and will bring more information and constraints on the evolution of massive stars at low Z. On the theoretical side, more models are necessary to fully understand and study the complex interplay between rotation, magnetic fields, mass loss and binary interactions at different metallicities. Large grids of models at low Z will have many applications, for example to study the evolution of massive stars and their feedback in high redshift objects like Lyman-break galaxies and damped Ly-alpha systems."
    },
    "0808/0808.1726_arXiv.txt": {
        "abstract": "Recent studies have indicated that the emission of gravitational waves at the merger of two black holes gives a kick to the final black hole. If the supermassive black hole at the center of a disk galaxy is kicked but the velocity is not large enough to escape from the host galaxy, it will fall back onto the the disk and accrete the interstellar medium in the disk. We study the X-ray emission from the black holes with masses of $\\sim 10^7\\: M_\\sun$ recoiled from the galactic center with velocities of $\\sim 600\\rm\\: km\\: s^{-1}$. We find that their luminosities can reach $\\ga 10^{39}\\rm\\: erg\\: s^{-1}$, when they pass the apastrons in the disk. While the X-ray luminosities are comparable to those of ultra-luminous X-ray sources (ULXs) observed in disk galaxies, ULXs observed so far do not seem to be such supermassive black holes. Statical studies could constrain the probability of merger and recoil of supermassive black holes. ",
        "introduction": "Recently, studies in numerical general relativity have shown that merged binary black holes can have a large recoil velocity through anisotropic emission of gravitational waves \\citep[e.g.][]{gon07,cam07}. The maximum velocity would reach $\\sim 4000\\rm\\: km\\: s^{-1}$, although the actual distribution of kick velocities is very uncertain. The discovery of such recoiled black holes is important for studies about the growth of black holes as well as the general relativity.  Supermassive black holes ($\\ga 10^6\\: M_\\sun$) at the centers of disk galaxies would be kicked through this mechanism. Although they would be bright just after the kick because of the emission from the accretion disk carried by the black holes, they would soon get dim as the disk is consumed by the black holes \\citep{loe07}. However, if the kick velocity of a black hole is not large enough to escape from the host galaxy, it will eventually fall back onto the galactic disk. If the time-scale of dynamical friction is large enough, it will revolve around the galactic center many times before it finally spirals into the galactic center. When the black hole passes the galactic disk, it will accrete the interstellar medium (ISM) in the disk \\citep{ble08}.  The accretion of the surrounding gas onto an isolated black hole has been studied by several authors \\citep*[][and references therein]{fuj98,ago02,mii05,map06}. Most of the previous studies focused on stellar mass ($\\sim 10\\: M_\\sun$) or intermediate mass black holes (IMBHs; $\\sim 10^3\\: M_\\sun$). The accretion rate and thus the luminosity of a black hole depend on the mass of the black hole and the density of the gas surrounding it (see equation~[\\ref{eq:dotm}]). Since the mass of a stellar-mass black hole is small, its luminosity becomes large enough to be observed only when it plunges into a high density region such as a molecular cloud. On the other hand, in this letter, we show that a recoiled supermassive black hole can shine even in the ordinary region of a galactic disk because of its huge mass. ",
        "conclusions": "We have shown that a supermassive black hole ejected from the center of the host disk galaxy will return to the galactic disk, if the initial velocity is smaller than the escape velocity of the galaxy. The black hole accretes the surrounding ISM and the resultant X-ray luminosity can reach $\\ga 10^{39}\\:\\rm erg\\: s^{-1}$, when it passes the apastrons in the disk. Although the luminosity of a recoiled supermassive black hole is comparable to that of ultra-luminous X-ray sources (ULXs), it is unlikely that many of the observed ULXs are the supermassive black holes."
    },
    "0808/0808.3515_arXiv.txt": {
        "abstract": "We generate an explicit four-fold infinity of physically acceptable exact perfect fluid solutions of Einstein's equations by way of conformal transformations of physically unacceptable solutions (one way to view the use of isotropic coordinates). Special cases include the Schwarzschild interior solution and the Einstein static universe. The process we consider involves solving two equations of the Riccati type coupled by a single generating function rather than a specification of one of the two metric functions. ",
        "introduction": "Perhaps the simplest of all procedures that one can think of for generating new exact solutions of Einstein's equations is the use of conformal transformations. Unfortunately, when applied to vacuum, no new solutions emerge via this procedure \\cite{exact}. However, when considering static fluid spheres in isotropic coordinates, the seed metric is not vacuum, but an unphysical fluid with pressure but zero energy density and viable new solutions do indeed emerge. We now have a rather vast array of methods for generating static fluid spheres \\cite{visser}, with many of the successful procedures relying on what amounts to the development of a linear equation of first order, for example \\cite{es1} and \\cite{es2}. In isotropic coordinates linear equations do not emerge directly \\cite{lvisser}. Rather, specifying one of the two metric functions leads to a differential equation of the Riccati type. Here we do not specify either metric function but rather solve two equations of the Riccati type coupled by a single generating function. Whereas we are able to solve this system for a variety of generating functions, we have found only one class of generating functions that gives rise to tractable and physically interesting solutions of Einstein's equations. ",
        "conclusions": "An explicit four-fold infinity of new physically acceptable exact perfect fluid solutions of Einstein's equations have been generated by solving two equations of the Riccati type coupled by a single generating function rather than specifying one of the metric functions. Special cases of these solutions include the Schwarzschild interior solution and the Einstein static universe. The solutions are qualitatively similar to the Tolman IV solution (see for example \\cite{es2}) and so should be of interest for the study of internal properties of neutron stars \\cite{lat}.  \\bigskip"
    },
    "0808/0808.1510_arXiv.txt": {
        "abstract": "We present Smoothed Particle Hydrodynamics simulations of protoclusters including the effects of the stellar winds from massive stars. Using a particle--injection method, we investigate the effect of structure in close proximity to the wind sources and the short--timescale influence of winds on protoclusters. We find that structures such as disks and gaseous filaments have a strong collimating effect on winds. By a different technique of injecting momentum from point sources into our simulations, we compare the large--scale and long--term effects of isotropic and intrinsically--collimated winds on protoclusters and find them to be similar, although the collimated winds take longer to exert a significant influence. We find that both types of wind are able to dramatically slow the global star formation process, but that the timescale on which they can expel significant quantities of mass from the cluster is rather long (approaching ten freefall times). Clusters may then experience rapid star formation very early in their lifetimes, before switching to a mode where gas is gradually expelled, while star formation proceeds much more slowly over many freefall times. This complicates any conclusions regarding slow star formation derived from measuring the star--formation efficiency per freefall time. We find that estimates of the efficacy of winds in dispersing clusters derived simply from the total wind momentum flux may not be very reliable. This is due to material being expelled from deep within stellar potential wells, ofen to velocities well in excess of the cluster escape velocity, and also to the loss of momentum flux through holes in the gas distribtuion. Winds have little effect on the accretion--driven stellar initial mass function (IMF) except at the very high--mass end, where they serve to prevent some of the most massive objects accreting more material. Feedback does not result in the formation of further massive stars through the monolithic collapse of massive cores. We also find that the morphology of the gas, the rapid motions of the wind sources and the action of large--scale accretion flows prevent the formation of bubble--like structures. These effects may make it difficult to discern the influence of winds on very young clusters.\\\\ ",
        "introduction": "Elucidating the effects of stellar feedback on embedded star clusters is a problem of key importance in astrophysics. $\\sim90\\%$ of newborn stellar systems do not survive their embedded phase to become open clusters (\\cite{2003ARA&A..41...57L}). In addition, star formation at the scale of molecular clouds and embedded clusters is not a very efficient process (e.g. \\cite{1991IAUS..147..275F}). There are two possible reasons for these observations. If star--forming complexes are initially gravitationally bound (as is usually assumed), some process must expel most of the molecular gas before it can be incorporated into stars, while also usually disrupting the system. Alternatively, if such complexes are not initially bound, the low star formation efficiency and low survival probability are a natural consequence of this, although bound stellar systems can be left behind in some cases (\\cite{2004MNRAS.347L..36C}, \\cite{2005MNRAS.361....2C}).\\\\ \\indent If star--forming systems are initially bound, an obvious candidate for a means of unbinding them is feedback from their own stars. In embedded clusters massive enough to form O--type stars (i.e. those with masses $>10^{3}$ M$_{\\odot}$, the ionizing radiation, winds and eventual supernova (SN) explosions of these stars are likely to dominate the combined effect of the (much more numerous) lower mass objects. Photoionization and winds are particularly important in this respect, since they act for the duration of the O--star's lifetime, whereas SNe cannot occur for at least $5$ Myr after the onset of star formation. This corresponds to the freefall time of a molecular cloud of number density $\\sim10^{2}$ cm$^{-3}$, approximately the minimum density required to form molecular gas able to cool efficiently and form stars (\\cite{2001ApJ...562..852H}). If it is the case that star formation in, and disruption of, molecular clouds occurs on timescales of around the local freefall time (\\cite{2000ApJ...530..277E}, but see \\cite{2007ApJ...654..304K}), SNe may occur too late, particularly as the density of a molecular cloud generally increases considerably (with a corresponding decrease in the freefall time) before any massive stars are able to form.\\\\ \\indent In this paper, we consider the effects of the winds from massive stars on the cluster in which they form. In reality, winds and photoionization act in concert and it is not clear which will dominate, nor if one will help or hinder the other (see \\cite{2001PASP..113..677C} for a discussion of this subject). In this paper, we model winds alone -- we leave a comparative study including winds and ionization to a later date.\\\\ \\indent The interaction of a stellar wind with the interstellar medium (ISM) has been modeled analytically and numerically by several authors, although usually in one--dimensional geometries. \\cite{1975ApJ...198..575S} and \\cite{1977ApJ...218..377W} introduced the idea of dividing the wind bubble into several concentric zones. The innermost zone in their models consists of freely--flowing wind and the outermost zone of undisturbed ISM. The number and contents of the zones in between is the subject of considerable debate. \\cite{1975ApJ...198..575S} have a single zone between the free--flowing wind and the undisturbed ISM, consisting of a shocked shell driven by the momentum of the wind. It is not difficult to show that this assumption produces a wind bubble whose radius $R(t)$ increases with time as $R(t)\\propto t^{\\frac{1}{2}}$ \\citep[e.g.][]{1988RvMP...60....1O,1999isw..book.....L}). \\cite{1977ApJ...218..377W} introduce two zones in this region -- an inner shell of hot shocked wind and an outer shell of shocked ISM separated by a contact discontinuity. The difference between the models lies in Steigman et al's assumption that the shocked wind can cool and that the bubble is momentum--driven, as opposed to Weaver et al's contention that the shocked wind cannot cool and that the bubble is thus pressure--driven (resulting in a radial evolution of the form $R(t)\\propto t^{\\frac{3}{5}}$). \\cite{2001PASP..113..677C} discuss at some length the validity of these assumptions, which essentially rests on whether the contact discontinuity between the shocked wind and the shocked ISM is stable, which governs how likely mixing and subsequent cooling of these phases is. Analytical work by \\cite{1983ApJ...274..152V} and simulations by \\cite{1996A&A...316..133G} suggest that the discontinuity may not be stable and the shocked wind may therefore be able to cool efficiently. In this case, the wind bubble is driven purely by ram pressure and $R(t)\\propto t^{\\frac{1}{2}}$.\\\\ ",
        "conclusions": "\\indent It is clear from these simulations that spherical stellar winds can be collimated by structures near the wind sources at a wide range of scales. Whether the winds are isotropic or intrinsically anisotropic (i.e. collimated by some structure very close to the source) has some influence on the effectiveness of the winds, but largely to delay the point at which winds become important in controlling the cluster dynamics and the star--formation process. We found that our intrinsically--collimated winds had similar effects to the isotropic winds, but delayed by $\\sim0.2$ freefall times, a relatively small fraction of the duration of our simulations. In either case, winds can have a strong effect on the evolution of protoclusters.\\\\ \\indent We have shown that winds can dramatically slow the process of conversion of molecular gas into stars. In our simulations, this was achieved largely by slowing or stopping accretion on existing stars, rather than preventing the formation of new stars, since few new stars formed in any of our simulations after the point at which wind sources were introduced ($\\sim1.8$ freefall times into the evolution of the system). In particular accretion onto the more massive wind sources themselves, which constitute a significant fraction ($\\sim20 \\%$) of the total stellar mass, was severely curtailed by isotropic or anisotropic winds. Since the accretion rates onto the intermediate-- and low--mass stars were not so strongly affected, we found that winds had little discernible effect on the cluster IMF except at the high--mass end, at and above a mass similar to the threshold mass we chose to determine whether a star was considered to be a wind source or not. It is possible that the (slight) changes in the IMF are therefore influenced by our choice of this parameter, so we do not attempt to draw strong conclusions on this subject, except that winds appear to have little influence on the generation of the IMF by competitive accretion. Although there do not appear to be any instances of star formation induced by winds in our calculations, this may only be a consequence of the fact that the clusters are only marginally bound and that the wind sources are turned on relatively late, after much of the bound gas has already fragmented. It s possible that, in clusters that are more strongly bound, winds acting earlier in the cluster's evolution may encourage fragmentation and induce star formation.\\\\ \\indent Additionally, the stellar winds from a few stars can come to dominate the dynamics of a whole cluster, so that the rate at which mass is being expelled from the cluster comes to exceed the rate at which mass is being converted to stars. In this case, the winds are beginning to unbind the cluster and an examination of the evolution of the kinetic, gravitational and total energies of our model clusters revealed that the action of winds is to stop and even reverse the decline in the cluster's total energy. However, the rates at which winds were able to expel gas from the clusters in our simulations were relatively low ($\\sim50$ M$_{\\odot}$ per freefall time, or $\\sim5$ percent of the cluster's total mass per freefall time). This amounts to $\\approx 3\\times10^{-4}$M$_{\\odot}$yr$^{-1}$. As can be seen by comparison with Figure \\ref{fig:wind_sourcelum}, this is much less than the momentum fluxes of the sources, if the gas is expelled at velocities comparable to the cluster escape velocity of a few kilometres per second. We identify three reasons for this discrepancy. Firstly, gas is generally expelled from regions near the massaive stars, where escape velocities are $\\sim100$km s$^{-1}$. Secondly, some of the gas leaving the cluster is moving at velocities in excess of the escape velocity. Thirdly, due to the inhomogenous gas distribution, some of the momentum flux from the sources is simply lost through holes in the gas. This result implies that the momentum flux of embedded wind sources alone may not be a good criteroin for determining whether a given protocluster can be dispersed by winds.\\\\ \\indent The slow rate of gas expulsion led to an interesting situation in which our wind--influenced clusters still possessed large quantities of gas which it would take the winds nearly ten freefall times to expel, but were converting gas to stars very slowly. The time for which the clusters were forming stars at a the rate dictated by purely isothermal evolution (as observed in the control run without winds) would then be only $\\sim20$ percent of the time for which they possessed molecular gas. If generally true, this would have important implications for calculations of star formation rates based on dividing the stellar mass of a system by an age derived from its oldest stars. If the cluster in question had spent a large fraction of its time forming stars very slowly because of the feedback effects observed in our winds runs, the inferred star formation rate would not be representative of the earliest stage of the cluster's evolution. More simulations are obviously required to determine the likelihood of this being true.\\\\ \\indent We note that the morphology of the gas in our wind--influenced simulations prevents the formation of more--or--less spherical wind bubbles. In our simulations, the high relative velocities of the wind sources to each other and to the gas, and the continuing large scale accretion flows present in the model clusters prevented the formation of such structures. It is actually very difficult to tell which of our simulations include winds and which do not from the gas morphology alone, except that the calculations including winds exhibit larger quantities of dense shocked gas. It may be true therefore that it is not always easy to see the effects of winds on young (i.e $<1$Myr old) clusters.\\\\ \\indent In this work, we have not included the effects of photoionising radiation. The wind sources in our calculations are massive enough to have strong ionising fluxes and these would also influence the evolution of the protocluster. It is not clear whether winds and photoionisation acting together help or hinder each other in unbinding stellar systems -- winds may help clear out dense material from around the massive stars, allowing ionising radiation to escape, but swept up shells of wind material may confine the HII regions at relatively small radii and prevent them having much effect. We leave this interesting question to a later paper.\\\\"
    },
    "0808/0808.1456_arXiv.txt": {
        "abstract": "{We have used spectra of hot stars from the RAVE Survey in order to investigate the visibility and properties of five diffuse interstellar bands previously reported in the literature.  The RAVE spectroscopic survey for Galactic structure and kinematics records CCD spectra covering the 8400-8800~\\AA\\ wavelength region at 7500 resolving power. The spectra are obtained with the UK Schmidt at the AAO, equipped with the 6dF multi-fiber positioner. The DIB at 8620.4~\\AA\\ is by far the strongest and cleanest of all DIBs occurring within the RAVE wavelength range, with no interference by underlying absorption stellar lines in hot stars. It correlates so tightly with reddening that it turns out to be a reliable tool to measure it, following the relation $E_{B-V} = 2.72 (\\pm 0.03)\\times E.W.$(\\AA), valid throughout the general interstellar medium of our Galaxy. The presence of a DIB at 8648~\\AA\\ is confirmed. Its intensity appears unrelated to reddening, in agreement with scanty and preliminary reports available in the literature, and its measurability is strongly compromised by severe blending with underlying stellar HeI doublet at 8649~\\AA. The two weak DIBS at 8531 and 8572~\\AA\\ do not appear real and should actually be blends of underlying stellar lines. The very weak DIB at 8439~\\AA\\ cannot be resolved within the profile of the much stronger underlying hydrogen Paschen~18 stellar line. ",
        "introduction": "\\begin{figure*} \\centering \\includegraphics[height=18.0cm,angle=270]{0232fig1.ps} \\caption{A sample of RAVE spectra of early type HD stars ordered according to reddening. The strongest stellar lines are identified. The arrows point to DIB 8620, the thick vertical marks to DIBs 8439, 8531, 8572, and the dashed line to DIB 8648~\\AA.} \\label{fig1} \\end{figure*}  Diffuse interstellar bands (DIBs) were first discovered by Heger (1922) as stable lines that did not follow the orbital motion seen in spectroscopic binaries. The first systematic studies began only with Merrill (1934, 1936), who noted similarities (occurrence, intensity and velocity) and differences (far wider and with diffuse edges) with respect to atomic interstellar lines. A few years later the first interstellar molecule (CH) was identified by means of absorption lines at 4300.3~\\AA\\ by Swings and Rosenfeld (1937). Since then, much progress has been recorded in understanding the absorption spectra of interstellar ions and molecules, while the origin of DIBs is still mysterious after almost a century after their discovery (Sarre 2006). The census of major DIBs in the optical region seems quite complete now, at least for those away from strong telluric absorption bands and stellar lines.  A compilation maintained by Jenniskens (2007) lists 281 DIBs over the wavelength range from 3980 to 9632~\\AA. DIBs are also observed in external galaxies (e.g. in the SMC by Cox et al. 2007b; in the LMC by Cox et al. 2006; in NGC 1448 by Sollerman et al. 2005; in M31 by Cordiner et al. 2008), in starbust complexes (Heckman and Lehnert 2000) and damped Ly-$\\alpha$ systems (Jukkarinen et al. 2004, York et al. 2006, Ellison et al. 2008)  The carriers of DIBs are still unknown. Some DIBs tend to show an appreciable correlation with reddening, even if others do not (e.g. Sanner et al. 1978, Krelowski et al. 1999), and this led to the hypothesis that they were produced on or in the interstellar dust grains. However, lack of polarization in DIB profiles (Cox et al. 2007a) argues against, and complex carbon-bearing molecules are generally considered as viable carriers (e.g. Fulara and Krelowski 2000), with fullerenes being subject of intensive laboratory studies (e.g. Herbig 1995, Leach 1995, Iglesias-Groth 2007). PAH (polycyclic aromatic hydrocarbons) have been frequently considered as promising carriers (van der Zwet and Allamandola 1985, Leger and Dhendecourt 1985), in particular for their success in explaining the unidentified infrared emission bands (UIR; Sarre et al. 1995). Laboratory studies suggest that only ionized and not neutral PAH could be viable carriers (e.g. Ruiterkamp et al. 2002, Halasinski et al. 2005), and this would agree with the observed insensitivity of DIB intensity to ambient electron density (Gnaciski et al. 2007). A search for correlations between different DIBs has been carried out as a criterion to guide possible identification (a strict correlation may imply a common carrier, whereas a lack of correlation indicates that different species are involved), but the degree of correlation cover the whole interval from good to very poor (e.g. Moutou et al. 1999). Finally, Holmlid (2008) has recently proposed a radically different mechanism for the formation of DIBs, namely doubly excited atoms embedded in the condensed phase named Rydberg matter.  The Radial Velocity Experiment (RAVE) is an ongoing spectroscopic survey of the whole southern sky at galactic latitudes $|b|$$\\geq$25$^\\circ$ for stars in the magnitude interval 9$\\leq$$I_{\\rm C}$$\\leq$12, with spectra recorded over the 8400-8800~\\AA\\ range at a resolving power around 7\\,500. The 150 fiber positioner 6dF is used at the UK Schmidt telescope of the Anglo-Australian Observatory. The main scientific driver of the project is the study of the stellar kinematics and metallicity of galactic populations away from the galactic plane (e.g. Smith et al. 2007, Seabroke et al. 2008, Veltz et al. 2008). RAVE begun operations in 2003 and has now reached the milestones of the first (Steinmetz et al. 2006) and second (Zwitter et al. 2008) data releases, with a current total of $\\sim$250\\,000 stars already observed.  In this paper we investigate the detectability, measurability and properties of the DIBs known to occur within the RAVE wavelength interval. In this range there is no resonant line from ions sufficiently abundant in the interstellar medium to produce detectable features in high-resolution ground spectra. Weak C$_2$ interstellar lines due to the (2,0) band of the Phillips system are seen longword of hydrogen Paschen~12 and around HeI 8777 stellar lines (Gredel and Muench 1986) in highly reddened stars, the strongest ones occurring at 8751.5, 8753.8, 8761.0, 8763.6, 8773.1 and 8780.0~\\AA\\ (Gredel et al. 2001). ",
        "conclusions": ""
    },
    "0808/0808.1984_arXiv.txt": {
        "abstract": "{Methanol masers at 6.7 GHz have been found exclusively towards high-mass star forming regions. Recently, some Class 0 protostars have been found to display conditions similar to what are found in hot cores that are associated with massive star formation. These hot corino sources have densities, gas temperatures, and methanol abundances that are adequate for exciting strong 6.7 GHz maser emission.} {This raises the question of whether 6.7 GHz methanol masers can be found in both hot corinos and massive star forming regions, and if not, whether thermal methanol emission can be detected.} {We searched for the 6.7 GHz methanol line towards five hot corino sources in the Perseus region using the Arecibo radio telescope. To constrain the excitation conditions of methanol, we observed thermal submillimeter lines of methanol in the NGC1333-IRAS~4 region with the APEX telescope.} {We did not detect 6.7 GHz emission in any of the sources, but found absorption against the cosmic microwave background in NGC1333-IRAS~4A and NGC1333-IRAS~4B. Using a large velocity gradient analysis, we modeled the excitation of methanol over a wide range of physical parameters, and verify that the 6.7~GHz line is indeed strongly anti-inverted for densities lower than 10$^{6}$~\\cmiii.  We used the submillimeter observations of methanol to verify the predictions of our model for IRAS~4A by comparison with other CH$_3$OH transitions. Our results indicate that the methanol observations from the APEX and Arecibo telescopes are consistent with dense (n $\\sim 10^6$~\\cmiii), cold ($T\\sim 15-30$~K) gas.} {The lack of maser emission in hot corinos and low-mass protostellar objects in general may be due to densities that are much higher than the quenching density in the region where the radiation field is conducive to maser pumping.} ",
        "introduction": "The $5_1-6_0$ A$^+$ line of methanol at 6.7 GHz is the strongest of Class II methanol masers. 6.7 GHz masers are unique compared to OH and H$_2$O masers in that they are associated exclusively with high-mass star forming regions. For example, targeted searches towards low-mass young stellar objects have not resulted in any detections of 6.7 GHz methanol masers \\citep{mini03,bour05,xu08}. Theoretical models suggest that the maser emission switches on at gas densities greater than $\\sim 10^4$~\\cmiii and methanol fractional abundance greater than $\\sim 10^{-7}$, along with the requirement of a far-infrared radiation field to excite methanol to torsionally excited states as part of the pumping cycle \\citep{sobo97, crag02, crag05}. These conditions are found in high-mass star forming regions where the far-infrared radiation field is generated by warm dust ($T_{D} \\sim 100-150$ K), and high methanol abundances are achieved by evaporating the mantles of dust grains, either by the passage of a shock or through the radiation field of the massive young stellar object. The molecules released from the grains further trigger formation of more complex molecules in warm gas chemistry (e.g. \\citealt{char92}).  Recently, there has been evidence of hot core chemistry in low-mass protostars. For example, complex molecules such as HCOOCH$_3$, HCOOH, and CH$_3$CN have been found in IRAS 16293--2422 \\citep{caza03} and NGC1333-IRAS~4A \\citep{bott04} along with high abundance of H$_2$O, CH$_3$OH, and H$_2$CO \\citep{cecc00a,cecc00b, scho02}. The abundance of methanol in the compact, warm regions in the immediate proximity of these embedded sources, known as hot corinos, can be as high as $7\\times 10^{-7}$ (\\citealt{mare05} and references therein). Moreover, the regions have warm temperature ($> 100$ K) and high density ($> 10^7$ \\cmiii; \\citealt{cecc00a}). This raises the question of whether hot corinos can excite maser emission in Class II methanol lines. A discovery of Class II methanol masers in hot corinos would force re-interpretation of the methanol maser data in the literature. In this paper, we report on the results of a search for 6.7 GHz methanol masers towards five hot corino sources in the Perseus region. ",
        "conclusions": "We have made the first conclusive detection of absorption of the 6.7 GHz line against the CMB towards hot corino sources NGC1333-IRAS~4A and NGC1333-IRAS~4B. A LVG analysis indicates that the 6.7 GHz line is indeed strongly anti-inverted for densities lower than $10^6$ \\cmiii~almost independent of temperature. Using observations of submillimeter $6_K-5_K$ lines of methanol using the APEX telescope, we model NGC1333-IRAS~4A, and are able to reproduce the observed 6.7 GHz absorption using cold ($T\\sim 15-30$ K), dense (n $\\sim 10^6$ \\cmiii) gas. The absence of maser emission at 6.7 GHz in hot corinos, and low-mass protostellar objects in general, is presumably due to the very high densities in the region where the radiation field would provide for maser pumping."
    },
    "0808/0808.0385_arXiv.txt": {
        "abstract": "{Clusters are potentially powerful tools for cosmology provided their observed properties such as the Sunyaev-Zel'dovich (SZ) or X-ray signals can be translated into physical quantities like mass and temperature.  Scaling relations are the appropriate mean to perform this translation. It is therefore, important to understand their evolution and their modifications with respect to the physics and to the underlying cosmology.}  {In this spirit, we investigate the effect of dark energy on the X-ray and SZ scaling relations. The study is based on the first hydro-simulations of cluster formation for diferent models of dark energy. We present results for four dark energy models which differ from each other by their equations of state parameter, $w$. Namely, we use a cosmological constant model $w=-1$ (as a reference), a perfect fluid with constant equation of state parameter $w=-0.8$ and one with $w = -1.2$ and a scalar field model (or quintessence) with varying $w$.}  {We generate $N$-body/hydrodynamic simulations that include radiative cooling with the public version of the Hydra code, modified to consider an arbitrary dark energy component. We produce cluster catalogues for the four models and derive the associated X-ray and SZ scaling relations.} { We find that dark energy has little effect on scaling laws making it safe to use the $\\Lambda$CDM scalings for conversion of observed quantities into temperature and masses.} {} ",
        "introduction": "In order to explain the current acceleration of the universe \\cite{riess1998, perlmutter1999} in the context of the theory of General Relativity, it is general procedure to introduce a new form of gravitational component with negative pressure -- dark energy. Various candidates such as a cosmological constant or a quintessence field have been proposed.  These models are characterised by their equation of state parameter and constrained by either using the observation of background quantities or the growth of cosmic structures. An alternative way to explain the acceleration of the universe is to allow for modifications of gravity. Many classes of models exist, for instance, a light scalar field coupled to matter leads to models of extended quintessence and more generally to scalar-tensor type theories. Further possibilities were studied in the context of braneworld models. Testing the Poisson equation on large scales may be a way to distinguish between all these alternative scenarios.  Regardless of its nature, dark energy as a dominant component, plays a role in structure formation and thus modifies the number of formed structures. The evolution of linear perturbations in the presence of a scalar field like quintessence and the effects on the abundance of collapsed structures and its dependence with redshift were widely explored and suggested as a tool to constrain the nature of dark energy and its evolution, e.g. \\cite{haiman2001, weller2001, weinberg2003, battye2003, wang2004, mohr2004}.  The properties of collapsed halos (density contrast and virial radius) depend strongly on the shape of the potential, the initial conditions, the time evolution of the dark energy equation of state and on its ability to collapse at the structure scale, e.g. \\cite{nunes2004}.  One can additionally investigate the effects of dark energy on the growth rate of structure and consequently study how dark energy affects the abundance of collapsed halos in dark energy models, e.g. \\cite{nunes2006, manera2006}.  Cluster number counts can potentially discriminate between dark energy models, in particular, the evolution of their equation of state. However, a large number of clusters with known redshifts from SZ, X-ray or optical surveys are needed, e.g. \\cite{bartelmann05}. Future Sunyaev-Zel'dovich (SZ) as well as X-rays observations will provide us with high quality data making galaxy clusters an efficient and powerful cosmological tool. Cluster physics, however, is complex. The presence of substructures, the possible contamination by radio and IR sources, the still imperfect knowledge of the relation between the halo mass and the clusters observed properties induce degeneracies between cluster physics and cosmological models. Scaling relations are key quantities in observational cosmology as they relate the observables in X-rays and SZ to the cluster properties, namely, masses and temperatures. The latter are then used in cluster number counts to constrain the cosmological models.  Due to their importance in translating observations to physical quantities, but most of all in probing the cluster formation, scaling relations have raised considerable attention.  Simple models of formation of virialised systems such as clusters predict that they exhibit self-similar behaviours \\cite{kaiser1986}, see also \\cite{ascasibar2006}. In the self-similar model, gravitational infall drives shock heating of the intra-cluster medium (ICM) and establishes the gas properties such that they scale with the halo mass giving rise to the scaling relations. However, it is now clear that additional physics is required to provide a more complete picture of the cluster formation and evolution, and to explain the deviations between observations and the predictions based on self-similar scaling as shown mostly by X-ray observations \\cite{edge91,allen98,markevitch98,nevalainen2000,finoguenov2001,ettori2004,henry2004,arnaud2005,rasia2005,balogh2006,maughan2006,morandi2007}. As the complexity of the physical description increases due to the additional gas physics (galactic winds and/or quasar outflows, radiative cooling, preheating) the use of numerical simulations appears the best option to compare predictions and observations and examine the role of those new ingredients in explaining the departures form self similarity (e.g. \\cite[]{evrard96,bryan98,bialek2001,thomas2001,babul2002, voit2002,borgani2004,rowley2004,muanwong2006,kay2007}).  The ICM can additionally be probed by the Sunyaev-Zel'dovich (SZ) effect, inverse Compton scattering of cosmic microwave background (CMB) photons off high energy electrons. The magnitude of the SZ effect is determined by the integrated gas pressure along the line-of-sight, called the Compton parameter, $y$.  Unlike the X-ray surface brightness, the SZ effect is not subject to cosmological dimming and can be used out to high redshift which makes SZ scaling relations particularly interesting and attractive tests. Nevertheless, SZ measurements, although gaining in quality and quantity, are at the moment still not sufficient to be fully used and this explains why investigations of SZ scaling relations are still limited from the observational point of view \\cite[]{cooray99, benson2004, mccarthy2003, morandi2007} as well as from the numerical point of view \\cite[]{dasilva2004,molt2005,bonaldi2007}.   Scaling laws relate the observed properties in X-rays and SZ to the cluster properties, namely, masses and temperatures which are then used to construct mass functions and number counts utilised to constrain the cosmological models. Understanding the possible biases in the scaling relations is thus essential for the use of clusters as cosmological probes. Previous studies, based on numerical simulations, have focused on the effects of additional gas physics in the scaling laws. In the present study, we explore the effects of the cosmological model on the scaling properties of galaxy clusters and their evolutions. We perform hydrodynamic numerical simulation of cluster formation and evolution assuming a simple radiative cooling model, rather than a more complete gas model, which allow us to single out only dark energy properties. We then investigate weather the X-ray and SZ scaling laws derived for different dark energy models depart from those obtained in the standard cosmological constant model (the $\\Lambda$CDM model) taken as our reference. In the next section we briefly present the dark energy models used for the first numerical simulation of cluster formation with hydrodynamics in dark energy dominated universes. We describe the simulation code and the procedure used to construct the X-ray and SZ cluster catalogues. In Sect. \\ref{sec:scal}, we present the X-ray and SZ scaling laws studied in the article together with the fitting procedure. Our results and conclusions are summarised in Sect. \\ref{sec:res} and \\ref{sec:conc} respectively. ",
        "conclusions": "\\label{sec:conc}  The abundance of clusters, their redshift distribution, as well as their clustering, are governed by the geometry of the universe and the power spectrum of the initial density perturbations. Gas physics related to cluster structure and evolution also enters through mapping of the cluster observable (SZ flux or X-ray luminosity) relative to the total mass of the cluster. As a result, galaxy cluster counts can be used as probes of cosmological and cluster properties. However, there are several requirements needed to achieve precise cosmological constraints: (i) advances in understanding the formation and evolution of cluster size halos; (ii) a good understanding of the selection function; (iii) robust observational proxies for the cluster mass.  The first condition relates to the conduction of large simulations. This is greatly achieved in the standard cosmological model with a cosmological constant. In the context of dark energy dominated universe, N-body simulations are becoming available for models different from the simple cosmological constant with constant or varying equation of state parameter. These simulations now provide us with a good understanding of the halo properties. They show that the halo mass function is well approximated by the Jenkins mass function. Simulations also indicate that the clusters halos are more concentrated in dark energy models, since structure grow earlier, than in lambda models and that concentrations are higher in models with varying than constant equation state. The second condition, understanding the cluster selection function, translates in understanding the limiting mass and the completeness of the surveys from realistic mock cluster catalogues. The mass (or temperature) selection function is directly linked with the cluster observed quantities through cluster scaling relations which is our third requirement (the need for a good proxy for the cluster mass).  In this work we have thus for the first time explored the scaling laws for both SZ and X-rays observations using hydrodynamic simulations of galaxy clusters in four dark energy models with constant or varying equation of states spanning a large class of models. We have studied the scaling properties at $z=0$ and their evolution with redshift. We have found that dark energy induces no modifications on the scaling laws at $z=0$ and presents very little differences from the cosmological constant model at higher redshifts.  While detailed simulations incorporating viable dark energy models remains a program in progress, it is reassuring that all models considered in this work predict similar scaling properties to the lambda model. The modeling of the cluster gas component appears to be nearly independent of the dark energy model.  Therefore, using the ``standard'' lambda model scaling relations for converting observable to masses and temperature in future surveys should not introduce any additional bias in the cosmological constraints derived from cluster counts. In this work, we have considered that dark energy does not cluster with dark matter. It would be interesting, however, to evaluate how our conclusions stand for numerical simulations in a scenario where dark energy is inhomogeneous and collapses along with dark matter during the formation of structure. This will be pursued in a forthcoming study.  \\begin{table*} \\begin{center} \\caption{\\label{tabela} Best fit values of the parameters $\\alpha$, $\\log A$ and $\\beta$ as well as their respective $1\\sigma$ errors. These values are valid within the redshift range $0<z<1.5$.  For the $L_{\\rm X}-T_{\\rm X}$ and $Y-L_{\\rm X}$ relations, the linear fit and the associated parameters are only valid in the range $0<z<0.75$ above which the linear fit is not a good approximation.} \\begin{tabular}{|l|l|rrrr|} \\hline Model    &     & $w=-1$ & $w=-0.8$ & 2EXP1 & $w=-1.2$  \\\\ \\hline $T_{\\rm X}-M$  & $\\alpha_{\\rm TM}$  & $0.620\\pm 0.029$ & $0.604\\pm 0.031$ & $0.581\\pm 0.033$ & $0.602\\pm 0.029$\\\\ & $\\log A_{\\rm TM}$ & $0.260\\pm 0.005$ & $0.267\\pm 0.004$ & $0.271\\pm 0.005$ & $0.258\\pm 0.005$ \\\\ & $\\beta_{\\rm TM}$ & $-0.228\\pm 0.020$ & $-0.249\\pm 0.017$ & $-0.264\\pm 0.023$ & $-0.230\\pm 0.023$  \\\\ \\hline $Y-M$  & $\\alpha_{\\rm YM}$ & $1.732\\pm 0.025$ & $1.730\\pm 0.025$ & $1.721\\pm 0.022$ & $1.752\\pm 0.024$ \\\\ & $\\log A_{\\rm YM}$ & $-5.910\\pm 0.004$ & $-5.906\\pm 0.004$ & $-5.902\\pm 0.005$ & $-5.910\\pm 0.003$ \\\\ & $\\beta_{\\rm YM}$  & $0.128\\pm 0.016$ & $0.116\\pm 0.016$ & $0.108\\pm 0.020$ & $0.135\\pm 0.013$  \\\\ \\hline $Y-T_{\\rm mw}$ & $\\alpha_{\\rm YT}$ & $2.922\\pm 0.100$ & $2.902\\pm 0.136$ & $2.838\\pm 0.072$ & $2.985\\pm 0.175 $\\\\ & $\\log A_{\\rm YT}$ & $-6.522\\pm 0.008$ & $-6.518\\pm 0.005$ & $-6.499\\pm 0.007$ & $-6.538\\pm 0.008$\\\\ & $\\beta_{\\rm YT}$  & $0.454\\pm 0.036$ & $0.443\\pm 0.022$ & $0.430\\pm 0.031$ & $0.517\\pm 0.036$\\\\ \\hline $L_{\\rm X}-T_{\\rm X}$  & $\\alpha_{\\rm LT}$ & $2.738\\pm 0.086$ & $2.691\\pm 0.089$ & $2.902\\pm 0.099$ & $2.796\\pm 0.146$ \\\\ & $\\log A_{\\rm LT}$ & $2.602\\pm 0.010$ & $2.629\\pm 0.007$ & $2.558\\pm 0.006$ & $2.589\\pm 0.011$ \\\\ & $\\beta_{\\rm LT}$ & $0.279\\pm 0.063$ & $0.027\\pm 0.042$ & $0.270\\pm 0.035$ & $0.348\\pm 0.070$  \\\\ \\hline $Y-L_{\\rm X}$  & $\\alpha_{\\rm YL}$ &  $1.063\\pm 0.028$ & $1.064\\pm 0.026$ & $1.076\\pm 0.026$ & $1.084\\pm 0.037$ \\\\ & $\\log A_{\\rm YL}$ & $-6.314\\pm 0.005$ & $-6.330\\pm 0.004$ & $-6.344\\pm 0.005$ & $-6.305\\pm 0.006$ \\\\ & $\\beta_{\\rm YL}$  & $0.668\\pm 0.033$ & $0.890\\pm 0.028$ & $0.770\\pm 0.034$ & $0.497\\pm 0.037$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}"
    },
    "0808/0808.2039_arXiv.txt": {
        "abstract": "{The presence of magnetic fields in O-type stars has been suspected for a long time. The discovery of such fields would explain a wide range of well documented enigmatic phenomena in massive stars, in particular cyclical wind variability, H$\\alpha$ emission variations, chemical peculiarity, narrow X-ray emission lines and non-thermal radio/X-ray emission.} {To investigate the incidence of magnetic fields in O stars, we acquired 38 new spectropolarimetric observations with FORS\\,1 (FOcal Reducer low dispersion Spectrograph) mounted on the 8-m Kueyen telescope of the VLT. } { Spectropolarimetric observations have been obtained at different phases for a sample of 13 O stars. 10 stars were observed in the spectral range 348--589\\,nm, HD\\,36879 and HD\\,148937 were observed in the spectral region 325--621\\,nm, and HD\\,155806 was observed in both settings. To prove the feasibility of the FORS\\,1 spectropolarimetric mode for the measurements of magnetic fields in hot stars, we present in addition 12 FORS\\,1 observations of the mean longitudinal magnetic field in $\\theta^1$\\,Ori\\,C and compare them with measurements obtained with the MuSiCoS, ESPaDOnS and Narval spectropolarimeters. } { Most stars in our sample which have been observed on different nights show a change of the magnetic field polarity, but a field at a significance level of 3$\\sigma$ has been detected only in four stars, HD\\,36879, HD\\,148937, HD\\,152408, and HD\\,164794. The largest longitudinal magnetic field, $\\langle$$B_z$$\\rangle$\\,=\\,$-$276$\\pm$88\\,G, was detected in the Of?p star HD\\,148937. We conclude that large-scale organised magnetic fields with polar field strengths larger than 1\\,kG are not widespread among O-type stars. } {} ",
        "introduction": "\\label{sect:intro} Massive stars usually end their evolution with a final supernova explosion, producing neutron stars or black holes. The initial masses of these stars range from $\\sim$9--10\\,M$_\\odot$ to 100\\,M$_\\odot$ or more, which correspond to spectral types earlier than about B2. Magnetic O stars with masses larger than 30\\,M$_\\odot$  and their WR descendants have been suggested as progenitors of magnetars (Gaensler et al.\\ \\cite{gaensler05}). Contrary to the case of Sun-like stars, the magnetic fields of stars on the upper main sequence (Ap/Bp stars), white dwarfs, and neutron stars are dominated by large spatial scales and do not change on yearly time scales. In each of these classes there is a wide distribution of magnetic field strengths, but the distribution of magnetic fluxes appears to be similar in each class, with maxima of $\\Phi_\\mathrm{max}=\\pi R^2B\\sim 10^{27-28}\\mathrm{G~cm^2}$ (Reisenegger \\cite{reisenegger01}, Ferrario \\& Wickramasinghe \\cite{ferrario05}), arguing for a fossil field whose flux is conserved along the path of stellar evolution. Braithwaite \\& Spruit (\\cite{braithwaite04}) confirmed through simulations that there are stable MHD configurations that might account for long-lived, ordered fields in these types of stars.  However, very little is known about the existence, origin, and role of magnetic fields in massive O and Wolf-Rayet stars. The lack of information is especially disturbing because magnetic fields may have paramount influence on the stellar evolution of high-mass stars. Maeder \\& Meynet (\\cite{maeder05}) examined the effect of  magnetic fields on the transport of angular momentum and chemical mixing, and they found that the potential influence on the  evolution of massive stars is dramatic.  Indirect observational evidence for the presence of magnetic fields are the many unexplained phenomena observed in massive stars, that are thought to be related to magnetic fields. One of the main indications that massive stars have magnetic fields is the cyclic behaviour on a rotational timescale observed in the UV wind lines (e.g.\\ Henrichs et al.\\ \\cite{henrichs:2005}). Other indications are variability observed in the H and He lines (Moffat \\& Michaud \\cite{moffat:1981}, Stahl et al.\\ \\cite{stahl:1996}, Rauw et al.\\ \\cite{rauw:2001}), narrow X-ray emission lines (Cohen et al.\\ \\cite{cohen:2003}, Gagn\\'e et al.\\ \\cite{gagne:2005b}) and the presence of non-thermal radio emission (Bieging et al.\\ \\cite{bieging:1989}, Scuderi et al.\\ \\cite{scuderi:1998}, Schnerr et al.\\ \\cite{schnerr:2006c}).  Direct measurements of the magnetic field strength in massive stars using spectropolarimetry to determine the Zeeman splitting of the spectral lines is difficult, as only fewer spectral lines are available for the measurements and which are usually strongly broadened by rapid rotation. So far a magnetic field has only been found in three O stars, $\\theta^1$\\,Ori\\,C, HD\\,155806 and HD\\,191612 (Donati et al.\\ \\cite{donati:2002},  Hubrig et al.\\ \\cite{hubrig07}, Donati et al.\\ \\cite{donati:2006}) and in a handful of early B-type stars (Henrichs et al.\\ \\cite{henrichs:2000a}, Neiner et al.\\ \\cite{coralie:2003a}, Neiner et al.\\ \\cite{coralie:2003b}, Neiner et al.\\ \\cite{coralie:2003c}, Hubrig et al.\\ \\cite{hubrig:2006}, Donati et al.\\ \\cite{donati:2006b}, Hubrig et al.\\ 2007). In this paper we present the results of the measurements of magnetic fields in 13 O type stars using FORS\\,1 at the VLT in spectropolarimetric mode. Our observations and the data reduction are described in Sect.~\\ref{sect:observations}, the obtained results in Sect.~\\ref{sect:results} and their discussion is presented in Sect.~\\ref{sect:discussion}. ",
        "conclusions": "\\label{sect:discussion}  Stellar magnetic fields have been discovered across a large range of spectral types (see Charbonneau \\& MacGregor \\cite{charbonneau:2001}). In late type stars, dynamos active in the convective layers are believed to be the origin of the observed magnetic fields. In earlier type stars, which have radiative envelopes, large scale magnetic fields of the order of a kilogauss have been discovered in Ap/Bp stars, but the exact origin of these fields is not yet known (Charbonneau \\& MacGregor \\cite{charbonneau:2001}, Braithwaite \\& Nordlund \\cite{braithwaite:2006}). About 10\\% of main-sequence A and B stars are slowly rotating, chemically peculiar, magnetic Ap and Bp stars and among their descendants, white dwarfs, 10\\% have high magnetic fields. The magnetic fields in magnetic white dwarfs could be fossil remnants from the main-sequence phase, consistent with magnetic flux conservation (Ferrario \\& Wickramasinghe \\cite{ferrario05}). If we assume that massive stars behave like Ap and Bp stars, then for a magnetic field detection probability of 10\\%, an O stars sample should consist of a larger number of unbiased targets, including  stars in clusters and the Galactic field at different ages, in different Galactic metallicity zones, and with different rotational velocities and surface composition. As we mention in Sec.\\,2, our sample is biased in the sense that it is restricted to O-type stars exhibiting wind variability, anomalous X-ray behaviour and brightness variations.  Still, this is the first time that magnetic field strengths were determined for such a large sample of stars, with an accuracy comparable to the errors obtained for the three previously known magnetic O-type stars, $\\theta^1$\\,Ori\\,C, HD\\,155806, and HD\\,191612. For the magnetic Of?p star HD\\,191612, Donati et al.\\ (\\cite{donati:2006}) measured a magnetic field of $\\langle B_\\mathrm {z}\\rangle=-220\\pm38$\\,G, averaging a total of 52 exposures obtained over 4 different nights. This is similar to our typical errors of a few tens of G. We have found four new magnetic O-type stars which have different spectral types, luminosity classes, and behaviour in various observational domains (see Appendix~\\ref{sect:app}). It opens the question whether O-type stars are magnetic in different evolutionary states. The study of the evolutionary state of one of the  Galactic Of?p stars, HD\\,191612, indicates that it is significantly evolved with an $\\sim$O8 giant-like classification (Howarth et al.\\ \\cite{how07}). The youth of the best studied magnetic O-type star $\\theta^1$\\,Ori\\,C and the older age of the Of?p star HD\\,191612 suggest that the presence of magnetic fields in O-type stars is not related to their evolutionary state. On the other hand, $\\theta^1$\\,Ori\\,C seems to possess a somewhat stronger magnetic field in comparison to that of HD\\,191612, indicating that the magnetic field could be a fossil remnant. Considering different size of radii due to the different evolutionary state and assuming conservation of the magnetic surface flux the magnetic field strength is expected to decrease by a factor proportional to the square of the ratio of their radii. Some support for the fossil magnetic field origin arises also from the recent comparative study of magnetic fields of early B-type stars at different ages by Briquet et al.\\ (\\cite{bri07}). This work revealed that the strongest magnetic fields appear in the youngest Bp stars, compared to weaker magnetic fields in stars at advanced ages.  It is notable, that the current analyses of Naz\\'e et al.\\ (\\cite{naze08}) and Naz\\'e et al.\\ (in preparation) suggest slower rotation than usually observed in O-type stars and nitrogen enhancement in both other Of?p stars, HD\\,108 and HD\\,148937. Their multiwavelength studies indicate rotation periods of $\\sim$55 yr for HD\\,108 and 7\\,d for HD\\,148937. Also, recent NLTE abundance analyses of early B-type stars by  Morel et al.\\ (\\cite{mor08}) and Hunter et al.\\ (\\cite{hunt08}) confirm that slow rotators often have peculiar chemical enrichments. It is remarkable that the observational data collected by Morel et al.\\ (\\cite{mor08}) strongly point to a higher incidence of magnetic fields in stars with nitrogen excess and boron depletion. Clearly, future studies are necessary to determine the efficiency of various indirect indicators for the presence of magnetic fields in the atmospheres of hot stars.  Yet, although it was possible to recognise a few hot magnetic stars as peculiar from their spectral morphology prior to their field detection (Walborn \\cite{walborn06}), the presence of a magnetic field can also be expected in stars of other classification categories. Our measurements of 13 O-type stars indicate that magnetic fields are possibly present in stars with very different behaviour in visual, X-ray and radio domains. Subsequent magnetic field measurements will test and constrain the conditions controlling the presence of magnetic fields in hot stars, and the implications of such fields on their mass-loss rate and evolution.  Since no longitudinal magnetic fields stronger than 300\\,G were detected in our study (apart from $\\theta^1$\\,Ori\\,C) we suggest that large scale, dipole like, magnetic fields with polar field strengths larger than 1\\,kG are not widespread among O type stars. Stars more massive than about 9~M$_\\odot$ end up as neutron stars or as black holes. A significant fraction of newborn neutron stars are strongly magnetized, with typical fields of $\\sim10^{12}$\\,G, and fields of up to $\\sim10^{15}$\\,G in the magnetars. Simple conservation of magnetic flux would imply field strengths of at least (5\\,R$_\\odot$/10\\ km$)^{-2}\\times10^{12}$\\,G $\\simeq 10^1$\\,G as a minimum for their progenitors. This is similar to the minimum field strength required to explain the wind variability observed in the UV (several $10^1$ G), as can be concluded from numerical simulations of wind behaviour in early type stars (ud-Doula \\& Owocki \\cite{asif:2002}). Our measurements have a typical accuracy of a few tens of G and it is quite possible that weak magnetic fields are present in the atmospheres of the other stars of our sample, but remain undetected as long as the measurement uncertainties are not significantly improved.  As we mention in Sec.\\,2, a new mosaic detector with blue optimised E2V chips was implemented in FORS\\,1 in 2007. To achieve the highest possible signal-to-noise (S/N) ratio -- as is required for accurate measurements of stellar magnetic fields -- the (200kHz, low, 1$\\times$1) readout mode can be used, which makes it possible to achieve a S/N  ratio of more than 1000 with only one single sub-exposure. Our recent tests show that using a sequence of 8--10 sub-exposures we can obtain much better accuracies down to 10--20\\,G. The typical uncertainties presented in Table 2 and 3 at the time of our observations before this CCD upgrade have been of the order of 40 to 70\\,G.  In the absence of further direct magnetic measurements, it is not clear yet whether more complex, smaller scale fields play a role in the atmospheres of O-type stars. In the case of a more complex magnetic field topology the longitudinal magnetic field integrated over the visible stellar surface will be smaller (or will even cancel) and not be easily detected with the low-resolution FORS\\,1 measurements, which allow to detect only magnetic fields which possess a large-scale organization. However, high resolution spectropolarimeters should be able to detect such complex field configurations using high signal-to-noise observations of the Zeeman effect in metal lines (e.g. Donati et al.\\ \\cite{donati:2006b}).  \\appendix"
    },
    "0808/0808.2513_arXiv.txt": {
        "abstract": "The helical magnetorotational instability of the magnetized Taylor-Couette flow is studied numerically in a finite cylinder. A distant upstream insulating boundary is shown to stabilize the convective instability entirely while reducing the growth rate of the absolute instability. The reduction is less severe with larger height. After modeling the boundary conditions properly, the wave patterns observed in the experiment turn out to be a noise-sustained convective instability. After the source of the noise resulted from unstable Ekman and Stewartson layers is switched off, a slowly-decaying inertial oscillation is observed in the simulation. We reach the conclusion that the experiments completed to date have not yet reached the regime of absolute instability. ",
        "introduction": "The magnetorotational instability (MRI) is probably the main source of turbulence and accretion in sufficiently ionized astrophysical disks \\citep{bh98}. Due to  this crucial role in astrophysics, substantial efforts have been spent worldwide to observe MRI in a laboratory setting \\citep{jgk01,gj02,npc02,sisan04,vils06}, but MRI has never been conclusively demonstrated in the laboratory.  Most experiments have been done in cylindrical geometry with a background flow that approximates the ideal Couette rotating profile: \\begin{equation} \\label{couette} \\Omega=a+b/r^{2} \\end{equation} where $a=(\\Omega_{2}r_{2}^{2}-\\Omega_{1}r_{1}^{2})/(r_{2}^{2}-r_{1}^{2})$ and $b=r_{1}^{2}r_{2}^{2}(\\Omega_{1}-\\Omega_{2})/(r_{2}^{2}-r_{1}^{2})$, $\\Omega_1$ and $\\Omega_2$ are the rotation speed of the inner and outer cylinder and $r_1$ and $r_2$ are the radius of the inner and outer cylinder, respectively (see Fig.~\\ref{boundary}). For axially periodic or infinite magnetized Taylor-Couette flow, MRI-like modes have been shown theoretically to grow at much reduced magnetic Reynolds number $\\Rm\\equiv\\Omega_1r_1(r_2-r_1)/\\eta$ and Lundquist number $S\\equiv V_{A,0}r_1(r_2-r_1)/\\eta$ in the presence of a combination of axial and current-free toroidal field % \\begin{equation} \\label{eq:backgroundfield} \\b{B}^{0}=B_z^{0}\\left(\\b{e}_z+\\beta r_1/r \\b{e}_\\varphi\\right) \\end{equation} than the standard MRI (SMRI) with purely axial magnetic field \\citep{hr05,rhss05}. Here the cylindrical coordinates $(r,\\varphi,z)$ are used. $B_z^{0}$ and $\\beta$ are constants. The Alfv\\'en speed is defined as $V_{A,0}\\equiv B_z^{0}/\\sqrt{4\\pi\\rho}$. $\\eta$ and $\\rho$ are the magnetic diffusivity and density of the fluid, respectively  (see Fig.~\\ref{boundary}).  The Potsdam ROssendorf Magnetic Instability Experiment (PROMISE) group claimed to have observed this kind of helical MRI (HMRI) experimentally \\citep{sgg06,rhsgg06,sgg07}. However we have shown that the wave pattern observed in PROMISE is not a global instability, but rather a transient disturbance somehow excited by the Ekman circulation and then transiently amplified as it propagates along the background axial Poynting flux with nonzero group and phase velocities, but is then absorbed once it reaches the jet formed at midheight between two neighboring Ekman cells \\citep{lgj07}. PROMISE group have accordingly updated the experimental facility to PROMISE II to allow for two split rings at both endcaps: the inner ring attached to the inner cylinder and outer ring attached to the outer cylinder. If the width of the inner ring is chosen appropriately $\\sim0.4(r_2-r_1)$, the magnetized Ekman circulation could be significantly reduced, therefore removing one of the possible disturbance sources, \\emph{i.e.} the unsteady jet \\citep{sj07}.  As with other examples in the literatires, such as drifting dynamo waves \\citep{tpk98, ptk00}, it is of vital importance to distinguish absolute instability from convective instability in a traveling wave experiment like PROMISE. It is also an essential ingredient of the threshold prediction for the Riga dynamo \\citep{ggglps08}. For a traveling wave the positivity of the growth rate implies only an amplification of the perturbation as it moves downstream. In one case, despite the movement of the wave packet, the perturbation increases without limit in the course of time at any point fixed in space; this kind of instability with respect to any infinitesimal perturbations will be called \\emph{absolute instability}. In the other case, the packet is carried away so swiftly that at any point fixed in space the perturbation tends to zero as $t\\rightarrow\\infty$; this kind will be called \\emph{convective instability} \\citep{ll87} (see the details of \\S\\ref{packet}). For PROMISE II, it appears that under the experimental conditions the second kind occurs. A recent preprint also highlights the importance of the distinction between absolute and convective instabilities in the context of the HMRI \\citep{pg08}.  In a Taylor-Couette experiment bounded by insulating endcaps, \\citet{tpk98} have pointed out that without any external disturbances except a small initial disturbance needed as a seed for the instability, the distant upstream insulating boundary acts as an ``absorbing\" boundary while the characteristics of the downstream endcap is unimportant. Due to this absorption the convective unstable state cannot be sustained by a uniform driving force, therefore this unstable mode eventually decays \\citep{tpk98}. This driving force is not the noise mentioned before, but the power to drive the instability, which in the usual Taylor-Couette experiments can be quantified by the magnetic Reynolds number $\\Rm$. This conclusion has been rigorously demonstrated in the very resistive limit in \\S II.C of \\citet{lghj06} using a perturbative approach and \\S II.D of \\citet{lghj06} using a modifed WKB analysis, showing that the insulating endcap entirely stabilizes the HMRI mode, which is a convective unstable mode given the parameters of the PROMISE experiment.  The absorbing boundary is essential to the development, regardless of how distant it may be. The larger height only defers the time when we have to wait for the boundary-induced dissipation to dominate \\citep{tpk98}. On the other hand, if $\\Rm$ exceeds a higher threshold $Re_{\\rm m,f}$, the driving force of the system overcome the dissipation and a globally unstable mode appears \\citep{tpk98}. Therefore in a bounded system the unstable mode appears at $Re_{\\rm m,f}$ rather than $Re_{\\rm m,c}$, where $Re_{\\rm m,c}$ is the critical magnetic Reynolds number for the onset of the convective unstable mode without the ``absorbing\" boundary. \\citet{tpk98} has showed that in the presence of an ``absorbing\" boundary and large $h$, a global unstable mode appears when \\[ \\Rm\\geqslant Re_{\\rm m,f}\\equiv Re_{\\rm m,a}+O(h^{-2})\\,, \\] where $Re_{\\rm m,a}$ is the critical magnetic Reynolds number corresponding to the onset of the absolute instability without the ``absorbing\" boundary.  This has raised a big obstacle for people to observe absolutely unstable HMRI in the laboratory. The advantage of HMRI itself, \\emph{i.e.}, unstable with a low critical Reynolds number ($3$ orders lower than the SMRI) conflicts with the necessarily high threshold of the onset of an absolute HMRI mode, \\emph{i.e.}, excited at a reasonably high critical magnetic Reynolds number, thus high Reynolds number, which would result in much more severe end-effects than people had expected. Moreover the fact that the critical Lundquist number must usually increase together with the magnetic Reynolds number and high ratio of toroidal-to-poloidal magnetic field requirement ($\\beta>1$) would even worsen the situation.  We also find that by nonlinear numerical simulation the insulating endcap reduces the growth rate of the absolute instability somewhat. The higher the height $h$ is, the less the growth rate is reduced (Table~\\ref{reduce}).  In a typical experiment, the experiment is, however, highly likely affected by small external noise either from a physical cause or experimental imperfection such as the misalignment of the cylinders. If the system is convectively unstable, \\emph{i.e.}, disturbances grow as they move downstream, noise would sustain structures in the system even if no global mode is unstable \\citep{dr87,ptk00}. In the present paper, we show by numerical simulations that the perturbations from the unstable magnetized residual Ekman layer and Stewartson layer at the upper endcap would play the role of ``noise\" generator, though this perturbation level is reduced with increasing axial magnetic field \\citep{lw08a}. What is observed in PROMISE II turns out to be a noise-sustained convective traveling wave, not the absolute unstable mode.   This paper is organized as follows: \\S\\ref{packet} presents the wave packet analysis in a unbounded cylinder, which is the basis of the following sections. We report the nonlinear simulation results with \\emph{partially} conducting boundary conditions of PROMISE II experiment in \\S\\ref{noise}. The final conclusions and implications to the HMRI experiments are given in \\S\\ref{discussion}. ",
        "conclusions": "\\label{discussion} In this paper, nonlinear simulations of the helical magnetorotational instability in a magnetized Taylor-Couette flow are performed. The geometry mimics PROMISE II experiment with endcaps split into two rings. The partially conducting boundary condition introduced in \\citet{lgj07} is used. The waves patters change with applied magnetic field as in the experiment. However via numerical tests, we find that the wave patterns observed in PROMISE II experiment are not due to a global instability, but rather a noise-sustained convective instability.  The importance of the distinction between absolute and convective instability in a bounded system with broken reflection symmetry is discussed. The addition of the toroidal magnetic field breaks the axial symmetry of the system. In such cases, the effects of distant upstream insulating boundaries on the absolute instability differs remarkably from the ones on the convective instability. The insulating endcap would only reduce the growth rate of the absolute instability, but would stabilize the convective instability entirely, however distant it may be. For the absolute instability, the more distant insulating endcap would less reduce the growth rate, while for the convective instability the more distant endcap would only have the system wait longer for the dissipation due to the ``absorption\" boundary to dominate. These discoveries cast great obstacles for people to observe the helical magnetorotational instability in the laboratory: An absolute HMRI is needed to observe the global unstable mode in the experiment.  Unfortunately it is not easy to derive the critical magnetic Reynolds number $Re_{\\rm m,f}$ of the absolute HMRI analytically in a bounded system. However we can get a rough estimate of $Re_{\\rm m,a}$, \\emph{i.e.}, the critical magnetic Reynolds number of the absolute HMRI in an unbounded system, by wave packet analysis (\\S\\ref{packet}) and the approximate dispersion relation from Fig.~\\ref{fig:global}. From Fig.~\\ref{fig:global}, we derive the group velocity $V_g\\sim1.08\\cm\\s^{-1}$, $\\gamma^0\\sim0.31\\s^{-1}$, $k_z^0\\sim0.52\\cm^{-1}$ and $\\sigma\\sim2.29\\cm^2\\s^{-1}$. Therefore $\\gamma^0-V_g^2/2\\sigma\\sim0.05\\s^{-1}>0$, which corresponds to an absolute HMRI instability with $Re_{\\rm m, a}\\sim0.07$. We therefore conjecture that $Re_{\\rm m,f}\\equiv Re_{\\rm m,a}+O(h^{-2})\\gtrsim0.07$ in PROMISE II. The critical magnetic Reynolds number is somehow one order of magnitude lower than the standard MRI, but still requires Reynolds number $Re\\sim10^5$. Therefore we need to rotate the cylinder typically with more than one hundred $\\unit{rpm}$. Such rotation rates are of course achievable, however with such a Reynolds number the advantage of HMRI with much lower Reynolds number, thus much lower end-effects, is not so great as people had expected.  Moreover in most HMRI unstable modes $\\beta>1$ is preferred, this suggests a toroidal magnetic field typically at $\\sim1,000\\G$, which requires axial currents $>10^4\\unit{A}$ inside the inner cylinder. This is a big engineering challenge in itself. The technical constraints prevent us to try to find the threshold by nonlinear numerical simulations such as: (1)~the current code can not afford large Reynolds number $(\\sim10^{5})$, which is required for HMRI to enter the absolutely unstable regime; (2)~from the global linear calculation, the HMRI mode is stabilized if an artificially low Reynolds number like $\\sim10^3$, which could be afforded by the current code, is employed.  The non-axisymmetric $m=1$ modes are observed in the experiments \\citep{sgg06,rhsgg06,sgg07}. Unfortunately since the simulations presented in this paper are all axisymmtric, the possibility to study this important mode is excluded. The extension of the current work to 3D will be the subject of the future study. Ruediger et al have already done some excellent work on this issue and found that given PROMISE parameters the nonaxisymmetric HMRI modes are always harder to be excited than the symmetric mode \\citep{rhss05,rs08}."
    },
    "0808/0808.0807_arXiv.txt": {
        "abstract": "In the dense supernova core, self-interactions may align the flavor polarization vectors of $\\nu$ and $\\overline\\nu$, and induce collective flavor transformations. Different alignment ansatzes are known to describe approximately the phenomena of synchronized or bipolar oscillations, and the split of $\\nu$ energy spectra. We discuss another phenomenon observed in some numerical experiments in inverted hierarchy, showing features akin to a low-energy split of $\\overline\\nu$ spectra. The phenomenon appears to be approximately described by another alignment ansatz which, in the considered scenario, reduces the (nonadiabatic) dynamics of all energy modes to only two $\\nu$ plus two $\\overline\\nu$ modes. The associated spectral features, however, appear to be fragile when passing from single- to multi-angle simulations. ",
        "introduction": "Supernova (SN) neutrinos continue to be a subject of great interest in astroparticle physics \\cite{Raff07}. In particular, renewed attention is being paid to collective features of flavor transformations induced by $\\nu$ ($\\overline\\nu$) self-interactions (see \\cite{Duan08} and refs.\\ therein). The observed collective phenomena of synchronized \\cite{Past02} and bipolar \\cite{Bip1,Bip2} oscillations, and the split of $\\nu$ spectra \\cite{Split1,Split2} can be largely understood, in flavor space, in terms of various alignment (or antialignment) ansatzes for the Bloch polarization vectors of $\\nu$ ($\\mathbf{P}$) and $\\overline\\nu$ ($\\overline\\mathbf{P}$) at different energy $E$. Conversely, lack of alignment indicates flavor decoherence \\cite{Sigl}.  In this work we deal with a minor---yet interesting---phenomenon previously observed in the numerical experiments of \\cite{Foglilast} and then of \\cite{Dighe,Volpe}, akin to a low-energy ``antineutrino spectral split'' in inverted hierarchy. Differently from the neutrino case, this feature appears to have a nonadiabatic origin. We show that the $\\overline\\nu$ spectral split observed in \\cite{Foglilast} can be approximately described in terms of $\\mathbf{P}$ and $\\overline\\mathbf{P}$ alignment along four global modes (at low and high energy, for $\\nu$ and $\\overline\\nu$). We also discuss a scenario with different input spectra, where the $\\overline\\nu$ split is slightly more evident, at least in the approximation of averaged trajectories. The effects described herein might play a role, in principle, in future low-energy SN neutrino observations via $\\overline\\nu_e +p \\to n+ e^+$ scattering.  \\begin{figure}[t] \\centering \\label{fig1} \\vspace*{-7mm} \\hspace*{2mm} \\epsfig{figure=Fig_01.eps,width =.87\\columnwidth,angle=0} \\vspace{-4mm} \\caption{Radial profiles adopted for the matter  ($\\lambda$) and self-interaction ($\\mu$) potentials, in the range $r\\in[10,\\,250]$~km.} \\end{figure}  \\vspace*{-2mm} ",
        "conclusions": "\\vspace*{-2mm}  We have studied some low-energy features of the $\\overline\\nu$ energy spectrum, focusing on the spectral split phenomenon emerging (in inverted hierarchy) in single-angle simulations. We have related its origin to nonadiabatic aspects of the $\\nu$ and $\\overline\\nu$ evolution, which can be simply modeled through an alignment ansatz with four energy modes. However, the $\\overline\\nu$ split feature appears to be fragile when passing to multi-angle simulations. Further studies are required to deepen its analytical understanding, as well as the conditions for its observability.  \\begin{figure}[b] \\centering \\vspace*{-7mm} \\hspace*{-6mm} \\epsfig{figure=Fig_06.eps,width =1.2\\columnwidth} \\vspace*{-8mm} \\caption{As in Fig.~5, but in multi-angle simulation. \\label{fig6}} \\end{figure}  \\vspace*{1mm} {\\bf Acknowledgments.} This work is supported in part by the Italian INFN and MIUR (Astroparticle Physics project), and by the E.U.\\ (ENTApP network). In Munich, A.M.\\ is supported by an INFN postdoc fellowship.  \\vspace*{-2mm}"
    },
    "0808/0808.2033_arXiv.txt": {
        "abstract": "Proton-rich material in a state of nuclear statistical equilibrium (NSE) is one of the least studied regimes of nucleosynthesis. One reason for this is that after hydrogen burning, stellar evolution proceeds at conditions of equal number of neutrons and protons or at a slight degree of neutron-richness. Proton-rich nucleosynthesis in stars tends to occur only when hydrogen-rich material that accretes onto a white dwarf of neutron star explodes, or when neutrino interactions in the winds from a nascent proto-neutron star or collapsar-disk drive the matter proton-rich prior to or during the nucleosynthesis. In this paper we solve the NSE equations for a range of proton-rich thermodynamic conditions. We show that cold proton-rich NSE is qualitatively different from neutron-rich NSE. Instead of being dominated by the Fe-peak nuclei with the largest binding energy per nucleon that have a proton to nucleon ratio close to the prescribed electron fraction, NSE for proton-rich material near freeze-out temperature is mainly composed of \\nuclei{56}{Ni} and free protons. Previous results of nuclear reaction network calculations rely on this non-intuitive high proton abundance, which this paper will explain. We show how the differences and especially the large fraction of free protons arises from the minimization of the free energy as a result of a delicate competition between the entropy and nuclear binding energy. ",
        "introduction": "\\label{sec:intro} Recently, a new nucleosynthesis process was invented to explain the production of the proton-rich isotopes \\nuclei{92,94}{Mo} and \\nuclei{96,98}{Ru} \\citep{frohlich06}. In this so called $\\nu p$-process, matter ejected from the surface layers of a proto-neutron star is exposed to a strong neutrino flux which results in the following weak interactions: \\begin{eqnarray} \\nu_e + n  &\\leftrightharpoons& p + e^- \\label{eq:one} \\\\ \\bar{\\nu}_e + p  &\\leftrightharpoons& n + e^+ \\label{eq:two} \\end{eqnarray} The mass difference between the neutron and the proton causes the neutrino interactions to be dominated by the forward reaction of eq.~(\\ref{eq:one}), which drives the nuclear matter proton-rich. The matter then expands and assembles into NSE, which is not only proton-rich (i.e. $Y_e > 0.5$) but actually contains a large mass fraction of free protons. The forward reaction of eq.~(\\ref{eq:two}) converts some of the free protons to neutrons, which allows nuclear matter to move past the ``bottle neck\" nucleus $\\nuclei{64}{Ge}$ via the fast $\\nuclei{64}{Ge}(n,p)\\nuclei{64}{Ga}$ reaction and the aforementioned proton-rich isotopes are synthesized during the freeze out phase from the proton-rich NSE state \\citep{frohlich06}. \\citet{meyer_1994_aa}, \\citet{jordan_2004_aa}, \\citet{pruet05,pruet06} and \\citet{wanajo06} calculate nucleosynthesis in similar environments.  We show that the NSE mass fractions exhibit a great degree of symmetry across the line $Y_e=0.5$ for high temperatures where the composition is dominated by free nucleons and $\\nuclei{4}{He}.$ For colder temperatures, ($T_9 \\leq 6.0$ and $\\rho \\sim 10^7 \\gcc$), there are hardly any free neutrons present in the NSE state for $0.4 < \\ye \\leq 0.5$. A na\\\"ive guess based on symmetry of NSE abundances would therefore not lead one to expect many free protons during freeze out for $\\ye > 0.5$ either. A large number of free protons is however observed to occur in nuclear reaction networks calculations. We show that for a relatively cold, near freeze-out temperature NSE state, there is indeed a qualitative difference in the mass fraction trends for $\\ye > 0.5$ and $\\ye < 0.5$. Restricting ourselves to the Ni isotopic chain and free nucleons, we explain how this difference arises as a result of a competition between the temperature dependent entropy term and the nuclear binding energy term contribution to the free energy. ",
        "conclusions": "We have presented mass fraction trends for NSE in proton-rich environments. We have explained the, at first sight peculiar, high free proton mass fractions for $\\ye>0.5$ and large mass fraction of \\nuclei{56}{Ni} even out to $\\ye=0.6$ by considering a simplified model. Restricting ourselves to the Ni isotopic chain, we have explicitly shown that for $\\ye>0.5$ ${\\mathcal F}$ is minimized by a state consisting of a mixture of free protons and the self-conjugate \\nuclei{56}{Ni}, whereas for $\\ye<0.5$ a state with a pure composition made up of the Ni isotope that has $Z/A$ closest to the prescribed $\\ye$ is preferred. In reality, other isotopic chains are of course accessible.  Other Fe-peak nuclides (especially the slightly neutron-rich even isotopes of Fe) compete with Ni isotopes for the most tightly bound nucleus with $Z/A$ near $\\ye$ (see fig. \\ref{fig:lowT}), giving the familiar NSE abundance pattern."
    },
    "0808/0808.0493_arXiv.txt": {
        "abstract": "We derive the cosmic star formation history (CSFH) out to $z = 1.3$ using a sample of $\\sim350$ radio-selected star-forming galaxies, a far larger sample than in previous, similar studies.  We attempt to differentiate between radio emission from AGN and star-forming galaxies, and determine an evolving 1.4\\,GHz luminosity function based on these VLA-COSMOS star forming galaxies.  We precisely measure the high-luminosity end of the star forming galaxy luminosity function (SFR $\\ga 100\\,M_{\\sun}\\,{\\rm yr}^{-1}$; equivalent to ULIRGs) out to $z = 1.3$, finding a somewhat slower evolution than previously derived from mid-infrared data.  We find that more stars are forming in luminous starbursts at high redshift. We use extrapolations based on the local radio galaxy luminosity function; assuming pure luminosity evolution, we derive $L_* \\propto (1+z)^{2.1 \\pm 0.2}$ or $L_* \\propto (1+z)^{2.5 \\pm 0.1}$, depending on the choice of the local radio galaxy luminosity function.  Thus, our radio-derived results independently confirm the $\\sim 1$ order of magnitude decline in the CSFH since $z \\sim 1$. ",
        "introduction": "Studies based on different galaxy star formation indicators (UV, optical, FIR, radio) agree that the cosmic star formation history (i.e.\\ the total star formation rate per unit co-moving volume; CSFH hereafter) has declined by about an order of magnitude since $z\\sim1$ (for a compilation see e.g.\\ \\citealt{hopkins04}). One of the major difficulties of UV/optical based tracers is the significant model-dependent dust-obscuration correction that needs to be imposed on the data. This `dust-obscuration problem' may be overcome using longer wavelengths, such as the IR and radio regimes.  However, in these cases a multi-wavelength approach is essential as redshift information and a reliable identifier of star forming (SF) galaxies is required (e.g.\\ \\citealt{caputi07, smo08}; S08 hereafter).  In this context the {\\em radio} star formation tracer provides an important complementary view of the CSFH. First, radio emission is a dust-insensitive tracer of recent star formation (not affected by old stellar populations; see \\citealt{condon92} for a review). Second, interferometric radio observations with $\\sim1''$ resolution allow more reliable identifications (compared to FIR and sub-mm data) with objects detected at other wavelengths.  The dust-unbiased total CSFH has been constrained recently using MIR (24/8\\mic) selected samples obtained by deep small area surveys (CDFS, GEMS, GOODS; \\citealt{lefloc'h05, zheng06, zheng07, caputi07, bell07}). Small area surveys, however, are subject to cosmic variance. Moreover, they do not observe a large enough comoving volume in order to fairly sample rare high-luminosity galaxies.   In this paper we use the 2\\sqdeg\\ COSMOS field \\citep{scoville07a}, and its 1.4~GHz radio observations \\citep{schinnerer07}, to derive the cosmic star formation history. In such a large field cosmic variance is significantly reduced as 2\\sqdegs\\ ($1.4^\\circ\\times1.4^\\circ$) sample comoving volumes in the early universe ($z\\sim1$) comparable to the largest survey in the local universe (SDSS -- DR1; see Fig.~1 in \\citealt{scoville07a}). The physical angular size sampled by 1.4$^\\circ$ at redshifts 0.2 -- 1.1 roughly corresponds to a factor of 3 to 8 of the typical cluster scale length ($\\sim5$~Mpc). Thus at all redshifts, such a field fairly samples relevant structures in the universe (see \\citealt{scoville07a, scoville07b, mccracken07} for a more detailed discussions on cosmic variance in the COSMOS field).  In the last decade several radio surveys have been utilized to independently derive the cosmic star formation history. However, their results are based on  small observed areas, non-uniform rms in the final map, as well as a fairly non-uniform selection of SF galaxies.  The first derivation of the CSFH based on radio data has been performed by \\citet{haarsma00}. They combined three radio frequency observations of the Hubble Deep Field ($66$~arcmin$^2$; \\citealt{richards98}), SSA13 ($7$~arcmin$^2$; \\citealt{windhorst95}), and V15 ($86$~arcmin$^2$; \\citealt{fomalont91, hammer95}) fields reaching $5\\sigma$ radio depths of 9, 8.8, and 16~$\\mu$Jy at the field centers, respectively. Of the total number of their radio-selected sources (77) only 37 were securely classified as star forming galaxies (based on morphology and/or optical spectroscopy). Their sources reach out to $z\\sim3$, 23 have spectroscopic redshifts, and the redshifts for the remaining 14 have been estimated using only I or H and K' bands (see \\citealt{haarsma00} for details). More recently, \\citet{seymour08} used radio observations of the 13$^\\mathrm{H}$ XMM-Newton/Chandra Deep field (0.196\\sqdeg ; $4\\sigma\\sim30~\\mu$Jy ) to derive the CSFH. Out of a total of 449 radio sources they find 269 galaxies which they classify as star forming based on a number of criteria applicable to sub-samples of their objects (see their Tab.~1). About half of these galaxies have a spectroscopic redshift ($z\\lesssim3$).  Here we utilize the 1.4~GHz VLA observations of the COSMOS 2\\sqdeg\\ field \\citep[VLA-COSMOS Large Project;][]{schinnerer07} to overcome the above mentioned biases. The final mosaic has a resolution of $1.5'' \\times 1.4''$ and a typical $rms$ of $\\sim 10.5~(15)$~$\\mu$Jy/beam in the central 1 (2) \\sqdegs\\ making this survey to date the largest radio deep field at this sensitivity and angular resolution. Given the COSMOS panchromatic data set, \\citet{smo08} have developed a novel method to select star forming and AGN galaxies based only on NUV--NIR photometry. This yielded a robust statistical selection of 340 SF galaxies out to $z=1.3$ (see \\s{sec:sample} ). One particular advantage of the VLA-COSMOS survey are the large comoving volumes observed at all redshifts, thus allowing one to study a statistically significant sample of rare objects, such as the most intensely star forming galaxies.  In this paper we focus on the derivation of the CSFH, with emphasis on the evolution of galaxies with high star formation rates ($\\gtrsim100$~\\Msol~yr$^{-1}$), using the 1.4~GHz VLA-COSMOS Large Project.   We report magnitudes in the AB system, adopt a standard concordance cosmology ($H_0=70,\\, \\Omega_M=0.3, \\Omega_\\Lambda = 0.7$), and define the radio synchrotron spectrum as $F_{\\nu} \\varpropto \\nu^{-\\alpha}$, assuming $\\alpha = 0.7$. ",
        "conclusions": ""
    },
    "0808/0808.4086_arXiv.txt": {
        "abstract": "We describe  a new method for measuring the extragalactic background light (EBL) through the detection of $\\gamma$-ray  inverse Compton (IC) emission due to  scattering of the EBL photons  off relativistic electrons in the  lobes of  radio galaxies. Our method has no free physical parameters and is a powerful tool when the lobes are  characterized by a  high energy sharp break or cutoff in their  electron energy distribution (EED). We show that such a feature will produce  a high energy  IC  `imprint'  of the  EBL spectrum in which the  radio lobes are embedded, and show how this  imprint  can be used to derive the EBL.  We apply our method to  the bright nearby radio galaxy Fornax A, for which we demonstrate, using WMAP and EGRET observations,  that the EED of its lobes is characterized by a  conveniently located cutoff, bringing the IC EBL emission into  the {\\sl Fermi} energy range. We show that {\\sl Fermi} will set upper limits to the optical EBL and measure the more elusive infrared EBL. ",
        "introduction": "}   The EBL  that permeates the Universe in the optical-IR is closely connected to the history of structure formation. It  is composed  by the cosmic infrared background (CIB)  peaking at  $\\lambda\\sim 100 \\mu m$ ( dust-reprocessed starlight) and the cosmic optical background  (COB)  peaking at $\\sim 1 \\mu m$ (starlight), and its  level  is still unknown to a factor of a few, due to the dominance of foregrounds such as the interplanetary dust emission  and Galactic  emission \\citep{hauser01,kashlinsky05}. Robust lower limits   to the EBL level  have been set by galaxy counts (e.g. Dole et al. 2006). Useful upper limits in the 1-10 $\\mu$m range  can be obtained by using the TeV blazar emission  to set limits on the EBL by modeling its attenuation due to pair production with EBL photons  (e.g. Stecker, de Jager, Salamon 1992). This method assumes that the intrinsic blazar TeV spectrum cannot be arbitrarily hard (e.g. Aharonian et al. 2006). However, this can  be circumvented  \\citep{katarzynski06, stecker07, aharonian08}, significantly relaxing  the EBL limits derived (Mazin \\& Raue 2007).   We propose  to measure the EBL  by detecting the high energy emission produced  when EBL photons IC-scatter off the relativistic electrons of extragalactic sources.  IC scattering of the cosmic microwave background (CMB) has  already been detected in the X-rays  from the lobes of a handful of radio galaxies (e.g. Erlund et al. 2008). In fact,  the first confirmation for the presence of the  CMB  in an extragalactic location  came from the ROSAT X-ray detection of  % the lobes of Fornax A by Feigelson et al. (1995). The idea we propose  is based on the fact that, although the EBL energy density is $\\sim$  only a few percent that of the CMB, the CIB and COB  photons  have  energies  correspondingly  $\\sim$ 10 and 1000 times  higher than those of the  CMB. For a source with a given EED, its electrons  will  IC-scatter  the CMB and EBL photons, and the IC emission will  have a spectral shape related  to that  of the CMB and EBL, shifted in frequency by $\\gamma_{\\rm max}^2$, where $\\gamma_{\\rm max}$ is the maximum Lorentz factor of the EED. Therefore, the IC emission from  such a source  will consist of a powerful component due to CMB seed photons and two weaker  components with a power of a few percent that of the  CMB one, but shifted in energy by a factor $\\sim 10$ and $1000$. The cleanest spectral separation of these  IC components results from a power-law EED characterized by a  high energy cutoff.  If this IC  emission can be detected, we will observe the high energy  imprint  of the CMB+EBL at the location of our  source.  This imprint  will provide us with  the level and spectral shape of the  CIB+COB  at the source. Here we show that the lobes of  radio galaxies, if characterized by  an EED  with  a suitably located high energy cutoff, can be used to measure the EBL at their location  through {\\sl Fermi} observations, and we demonstrate our method on the  radio lobes of the nearby bright radio galaxy Fornax A. ",
        "conclusions": "}  We have presented a novel method for measuring the EBL through   the detection of  the IC emission resulting from the upscattering of EBL photons  by relativistic electrons in the lobes of radio galaxies. The requirements that the sources  must fulfill (extended, high Galactic latitude, X-ray detection of the IC emission due to CMB seed photons, sharp break of the EED, no significant AGN GeV emission) quickly narrow down the number of  candidate  sources.  Using existing multiwavelength data (radio, WMAP, X-ray, EGRET), we show that in the case of  Fornax A  the  EBL imprint is detectable by {\\sl Fermi}, even for the lowest expected EBL level.   Our method is parameter free, in the sense that  all  physical parameters are directly derived from observations.  The radiating electron production mechanism (e.g. electron shock acceleration and/or electrons created in hadronic processes)  is not relevant for our work, since we derive the EED, regardless of it origin,  from multiwavelength data. {\\sl Fermi} is observing in scanning mode, so observations of Fornax A are guaranteed. Additional multiwavelength observations in the future can help refine further the derived EBL. We also outline a procedure for solving the inverse problem of going from the upcoming  {\\sl Fermi} observations to the EBL determination. Our method will determine the CIB (and set upper limits on the COB),  finally measuring  this cosmologically important quantity."
    },
    "0808/0808.2985.txt": {
        "abstract": "We calculate the decay rate of bottomonium to two-charm quark jets $\\Upsilon \\to c \\bar c$ at the tree level and one-loop level including color-singlet and color-octet $b \\bar b$ annihilations. We find that the short distance coefficient of the color-octet piece is much larger than the color-singlet piece, and that the QCD correction will change the endpoint behavior of the charm quark jet. The color-singlet piece is strongly affected  by the one-loop QCD correction. In contrast, the QCD correction to the color-octet piece is weak. Once the experiment can measure the branching ratio and energy distribution of the two-charm quark jets in the $\\Upsilon$ decay, the result can be used to test the color octet mechanism or give a strong constraint on the color-octet matrix elements. ",
        "introduction": "It is commonly believed that the heavy quark pair production and annihilation decay can be described by perturbative Quantum Chromodynamics (pQCD) since the heavy quarkonium mass provides a scale that is much larger than $\\Lambda_{QCD}$.  Due to its nonrelativistic nature, the heavy quarkonium annihilation decay is expected to be described in an effective theory, non-relativistic Quantum Chromodynamics (NRQCD)\\cite{Bodwin:1994jh}. In the NRQCD factorization formalism, the decay of heavy quarkonium is described by a series of annihilations of the heavy quark pair states and corresponding long-distance matrix elements, which are scaled by the relative velocity $v$ of quark and antiquark in the quarkonium rest frame. The heavy quark pair states can have not only the same quantum numbers as those of the quarkonium, but also other different quantum numbers in color and angular momentum. In particular, the heavy quark pair can be in a color-octet state.  The color-octet scenario seems to acquire some significant successes in describing heavy quarkonium decay and production. But recently, several next-to-leading order (NLO) QCD corrections for the inclusive and exclusive heavy quarkonium production in the color-singlet piece are found to be large and significantly relieve the conflicts between the color-singlet model predictions and experiments. It may imply, though inconclusively, that the color-octet contributions in the production processes are not as big as previously expected, and the color-octet mechanism should be studied more carefully.  The current experimental results on inelastic $J/\\psi$ photoproduction at HERA are adequately described by the NLO color singlet piece~\\cite{Kramer:1995nb}. The DELPHI data favor the NRQCD formalism for $J/\\psi$ production in $\\gamma  \\gamma \\rightarrow J/\\psi X$, rather than the color-singlet model\\cite{Klasen:2001cu,Abdallah:2003du}. The large discrepancies in $J/\\psi$ production via double $c\\bar c$ in $e^+e^-$ annihilation at B factories between LO theoretical predictions \\cite{Braaten:2002fi, Liu:2002wq,Hagiwara:2003cw,Yuan:1996ep,Liu:2003zr,Liu:2003jj} and experimental results ~\\cite{Abe:2002rb,Aubert:2005tj} are probably resolved by including the higher order corrections: NLO QCD corrections and relativistic corrections \\cite{Zhang:2005ch,Zhang:2006ay,Zhang:2008gp, Gong:2008ce,Gong:2007db,He:2007te,bodwin06}. The NLO QCD corrections in $J/\\psi$ and $\\Upsilon$ production at the Tevatron and LHC are calculated including the color singlet piece~\\cite{Campbell:2007ws,Artoisenet:2007xi} and the color octet piece~\\cite{Gong:2008ft}. The QCD corrections to polarizations  of $J/\\psi$ and $\\Upsilon$ at the Tevatron and LHC are also calculated ~\\cite{Gong:2008hk,Gong:2008sn,Gong:2008ft}. The experimental data of polarizations at the Tevatron seem to favor the NLO QCD corrections of the color singlet piece rather than the color octet piece. Recent developments and related topics in quarkonium physics can be found in Refs.~\\cite{Brambilla:2004wf,Lansberg:2006dh,Lansberg:2008zm}.         In order to further test the color octet mechanism, in this paper we calculate the rate of bottomonium decay into a charm quark pair $\\Upsilon \\to c \\bar c$. There have been some works on bottomonium decays and the color octet mechanism. Fritzsch and Streng calculated the decay rate of $\\Upsilon$ into charm at leading order in $\\alpha_s$, $\\Upsilon \\to g gg^* \\to g g c\\bar c$\\cite{Fritzsch:1978ey}. Bigi and Nussinov have taken into account the contribution of $\\Upsilon \\to g g^*g^* \\to g c\\bar c$ \\cite{Bigi:1978tj}. Barbieri, Caffo, and Remiddi have calculated the decay rates of the $P$-wave bottomonium states into charm at leading order in $\\alpha_s$ \\cite{Barbieri:1979gg}. Maltoni and Petrelli calculated the effects of color-octet contributions on the radiative $\\Upsilon$ decay \\cite{Maltoni:1998nh}. Recently, Bodwin, Braaten and Kang calculated the inclusive decay rate of $\\chi_b$ into charmed hadron in the NRQCD framework\\cite{Bodwin:2007zf}. Gao, Zhang and Chao calculated the bottomonium radiative decays to charmonium and light mesons\\cite{Gao:2006bc,Gao:2007fv}, as well as $\\Upsilon$ radiative decay to light quark jet to test the color octet mechanism\\cite{Gao:2006ak}. The S-wave quarkonium decay to light hadrons was calculated up to order $v^4$ and $\\alpha_s^3$ \\cite{Petrelli:1997ge,Bodwin:2002hg}. The exclusive double charmonium production from $\\Upsilon$ decay was calculated by Jia \\cite{Jia:2007hy}. Kang, Kim, Lee and Yu have calculated the inclusive charm production in $\\Upsilon(nS)$ decay \\cite{Kang:2007uv}. And the invariant-mass distribution of $c \\bar c$ in $\\Upsilon(1S) \\to c\\bar c + X$ was also calculated by Chung, Kim and Lee\\cite{Chung:2008yf}. And  $\\eta_b$  inclusive charm decay was calculated by Hao, Qiao and Sun \\cite{Hao:2007rb}. As to experiment, ARGUS Collaboration searched for charm production in direct decays of the $\\Upsilon(1S)$, and found $B^{dir}[\\Upsilon(1S) \\to D^*(2010) ^\\pm + X] < 0.019$\\cite{Albrecht:1992ap}. Very recently CLEO has searched for the $D^0$ production in direct decays of the $\\chi_{bJ}(nS)$ (n=1,2) states \\cite{Briere:2008cv}. The present investigation for the $\\Upsilon$ decay to $c\\bar c$ pair will hopefully add a new contribution to the test of color-octet mechanism in heavy quarkonium decays.       This paper is organized as follows. In Sec.~\\ref{sec:TheoFram}, we present the theoretical framework for the decay of $\\Upsilon \\to c \\bar c$. In Sec.~\\ref{sec:singlet}, we estimate the color-singlet contributions. In Sec.~\\ref{sec:octet}, we include the color-octet contributions. In Sec.~\\ref{sec:matrixelement}, we discuss the NRQCD matrix elements e.g. $\\left\\langle\\Upsilon|{ \\cal O}({}^3S_{1,8})| \\Upsilon \\right\\rangle$ and give a numerical estimation of the color-octet contributions. %In Sec.~\\ref{sec:exclusive}, we discuss %the exclusive decay of the $\\Upsilon$ to $D^+ D^{*-}$ and $D^{*+} %D^{*-}$ through the $B$ factory data and fragmentation functions. Summary and discussion are presented in Sec.~\\ref{sec:summ&dis}. The detailed and lengthy intermediate steps and formulas in the calculation will be given in the appendices. ",
        "conclusions": "%------------------------------------------------ \\label{sec:summ&dis} We calculate the decay rate of bottomonium to two-charm quark jets $\\Upsilon \\to c \\bar c$ at the tree level and one-loop level including color-singlet and color-octet $b \\bar b$ annihilations. We find that the short distance coefficient of the color-octet piece is much larger than the color-singlet piece, and that the QCD correction will change the endpoint behavior of the charm quark jet. The color-singlet piece is strongly affected  by the one-loop QCD correction. In contrast, the QCD correction to the color-octet piece is weak. Once the experiment can measure the branching ratio and energy distribution of the two-charm quark jets in the $\\Upsilon$ decay, the result can be used to test the color octet mechanism or give a strong constraint on the color-octet matrix elements. \\newline  After our work was completed \\cite{zhang:2007phd}, a paper appeared \\cite{Kang:2007uv}, in which Kang, Kim, Lee, and Yu calculated the inclusive charm production in $\\Upsilon(nS)$ decay. They focused on the inclusive charm production of the color-singlet piece at leading order in the strong coupling constant $\\alpha_s$. We focused on the $c \\bar c$ final state and the color-octet mechanism. The $c \\bar c$ final state is essentially the two charm-jet process. And we have calculated the next-to-leading order QCD corrections in both color singlet and color octet pieces. Our leading order result of $\\Upsilon \\to \\gamma^* \\to c \\bar c$ is consistent with their result \\cite{Kang:2007uv}.  %--------------------------------------------------------------------"
    },
    "0808/0808.2339_arXiv.txt": {
        "abstract": "We report on Suzaku and Chandra observations of the young supernova remnant CTB37B, from which TeV $\\gamma$-rays were detected by the H.E.S.S. Cherenkov telescope. The 80~ks Suzaku observation provided us with a clear image of diffuse emission and high-quality spectra. The spectra revealed that the diffuse emission is comprised of thermal and non-thermal components. The thermal component can be represented by an NEI model with a temperature, a pre-shock electron density and an age of 0.9$\\pm0.2$ keV, 0.4$\\pm0.1$~cm$^{-3}$ and 650$^{+2500}_{-300}$ yr, respectively. This suggests that the explosion of CTB37B occurred in a low-density space. A non-thermal power-law component was found from the southern region of CTB37B. Its photon index of $\\sim$1.5 and a high roll-off energy ($\\gtsim$15~keV) indicate efficient cosmic-ray acceleration. A comparison of this X-ray spectrum with the TeV $\\gamma$-ray spectrum leads us to conclude that the TeV $\\gamma$-ray emission seems to be powered by either multi-zone Inverse Compton scattering or the decay of neutral pions. The point source resolved by Chandra near the shell is probably associated with CTB37B, because of the common hydrogen column density with the diffuse thermal emission. Spectral and temporal characteristics suggest that this source is a new anomalous X-ray pulsar. ",
        "introduction": "Supernova Remnants (SNRs) are one of the most promising acceleration sites of cosmic rays. In fact, ASCA detected synchrotron X-ray emission from the shell of SN~1006, which unambiguously indicates the acceleration of electrons up to $\\sim$100~TeV \\citep{1995Natur.378..255K}. Following this discovery, the synchrotron X-ray emission has been discovered from a shell of a few more SNRs, such as RX~J1713.7--3946 \\citep{1997PASJ...49L...7K} and RCW~86 \\citep{2000PASJ...52.1157B}.  On the other hand, TeV $\\gamma$-rays have also been detected from some non-thermal shell-type SNRs. The radiation of TeV $\\gamma$-ray is explained by either (1)~Inverse-Compton scattering~(IC) of cosmic microwave background photons by the same high energy electron giving rise to the X-ray synchrotron emission or (2)~the decay of neutral pions that are generated by collisions between high energy protons and dense interstellar matter. The ratio of fluxes between the TeV $\\gamma$-ray and the X-ray provides the magnetic field intensity as long as one assumes that the TeV $\\gamma$-ray is produced through the IC mechanism. Utilizing this characteristic, \\citet{2007PASJ...59S.199M} found that the TeV $\\gamma$-ray from HESS~J1616$-$508 is likely the result of the proton acceleration, because the non-detection of X-ray using the Suzaku XIS provides much weaker magnetic field than the interstellar average.  Although the evidence of particle acceleration has accumulated rapidly, our knowledge is still limited on what sort of conditions are necessary for SNRs to accelerate particles. A breakthrough may be brought about by searching SNRs from which the TeV $\\gamma$-ray emission is already detected for thermal emission systematically, since the thermal emission provides us with a lot of information on the environment such as temperature, density, and age of the plasma.  CTB37B locates at ($l$, $b$) = ($\\timeform{348D.7}$, $+\\timeform{0D.3}$) with a distance of $10.2\\pm 3.5$~kpc \\citep{1975A&A....45..239C}. This region is one of the most active regions in our galaxy where star burst activities, a number of shell structures probably associated with recent SNRs \\citep{1991ApJ...374..212K}, and OH maser sources \\citep{1996AJ....111.1651F} are detected in radio band. TeV~$\\gamma$-ray emission is also detected by the H.E.S.S. observation \\citep{2007A&A...472..489A}. In spite of the evidence of the high activities in other wave bands, X-ray observations have been relatively poor. Only ASCA \\citep{1994PASJ...46L..37T} has detected a part of CTB37B at the edge of the field of view of the Gas Imaging Spectrometer (GIS) \\citep{1996PASJ...48..157O,1996PASJ...48..171M} in the course of the galactic plane survey (\\cite{2001ApJS..134...77S}; Yamauchi et al. 2008). Although the statistics are limited and the response of the GIS is not qualified at the pointing position of CTB37B, Yamauchi et al.~(2008) represents that the fit of a power law to the GIS spectrum results in a steep photon index of $\\sim$4.1, whereas fit of an optically thin thermal plasma model requires a high temperature of $\\sim$1.6~keV. These results strongly suggest that the X-ray spectrum is a mixture of a non-thermal power law and an optically thin thermal plasma emission. In addition, \\citet{2008arXiv0803.0682H} resolved a bright point source located near the shell of CTB37B from the diffuse emission by Chandra, although its spectral parameters are not constrained very well because of short exposure time.  In order to take an image and high quality spectra of CTB37B, we have carried out an observation of CTB37B with Suzaku. We also refer to the Chandra archival data to include the spatial structure and to compare them to our Suzaku data. In \\S~2, we present the observation log and data reduction method. Image analysis is presented in \\S~3. Spectrum analysis and timing analysis are shown in \\S~4 and \\S~5, respectively. We have really detected both the thermal and non-thermal power-law components from CTB37B. Discussions are made in \\S~6 on the nature of the thermal and non-thermal component as well as the point source. Finally we summarize our results in \\S~7. ",
        "conclusions": "We have obtained with Suzaku the images and the high quality spectra of the supernova remnant CTB37B. The X-ray diffuse emission region coincides with that of radio and TeV $\\gamma$-ray. The X-ray emission consists of thermal and non-thermal diffuse components as well as a point source resolved by Chandra. CTB37B is the third SNR from which thermal and non-thermal X-ray emissions as well as TeV $\\gamma$-ray emission are detected all together, and the fifth SNR that is accompanied by non-thermal emission both in X-ray and TeV $\\gamma$-ray bands.  The diffuse thermal emission can be best described by a non-equilibrium collisional ionization plasma model (NEI model) with a temperature, an ionization parameter~($n_{\\rm e}t$~[cm$^{-3}$s]), and the abundances of 0.9$\\pm0.2$~keV, 3.5~$^{+13}_{-1.1}\\times10^{10}$, and $\\simeq$ 0.5$Z_{\\odot}$~(Mg, Si), respectively. The image size and the observed emission measure provides the number density of the thermal electrons before the shock to be 0.2--0.4~cm$^{-3}$, which is significantly lower than that of the Galactic plane. This suggests that the supernova explosion associated with CTB37B took place at a low density space. From the ionization parameter and the number density of the thermal electron, the age of the plasma is found to be $\\sim$650~$^{+2500}_{-300}$~yr. This is consistent with the tentative identification of CTB37B with SN393 within the error.  In contrast, the diffuse component occupying southern part of CTB37B (region 2) is non-thermal and represented by a power-law model or a srcut model. The photon index of 1.5 is significantly smaller than any other non-thermal SNR, but is consistent with that of a typical non-thermal SNR in the radio band. The srcut model fit with its normalization set free to vary therefore results in a high roll-off energy of $>$15~keV. Under the assumption that TeV $\\gamma$-ray was emitted by 1-zone IC scattering, there are no solution for magnetic field strength that can reproduce the observed synchrotron spectrum in X-ray. This suggests that TeV $\\gamma$-ray is produced by multi-zone IC scattering, or by the decay of neutral pions generated by the high energy proton impacts.  Owing to the high spatial resolution of Chandra, a point source is resolved from the brightest part of the Suzaku image of CTB37B (region~1). Its association to the diffuse thermal emission indicated by $N_{\\rm H}$, the photon index of $\\sim$3, the X-ray luminosity of order $10^{34}$~erg~s$^{-1}$, and the long term flux variation evident from the Chandra and Suzaku observations all indicate that the point source is a new anomalous X-ray pulsar. A high speed photometric observation is encouraged.  \\bigskip The authors are greatful to all of the Suzaku and H.E.S.S. team members. We also thank H. Yamaguchi, T. Tanaka and J. Vink for useful comments. This work was supported in part by Grant-in-Aid for Scientific Research of the Japanese Ministry of Education, Culture, Sports, Science and Technology, No.~19$\\cdot$4014 (A.~B.), No.~18740153, No.~19047004 (R.~Y.)."
    },
    "0808/0808.2425_arXiv.txt": {
        "abstract": "We propose a class of models in which the $\\eta$-problem of supersymmetric hybrid inflation is resolved using a Heisenberg symmetry, where the associated modulus field is stabilized and made heavy with the help of the large vacuum energy during inflation without any fine-tuning. The proposed class of models is well motivated both from string theory considerations, since it includes the commonly encountered case of no-scale supergravity K\\\"ahler potential, and from the perspective of particle physics since a natural candidate for the inflaton in this class of models is the right-handed sneutrino which is massless during the inflationary epoch, and subsequently acquires a large mass at the end of inflation. We study a specific example motivated by sneutrino hybrid inflation with no-scale supergravity in some detail, and show that the spectral index may lie within the latest WMAP range, while the tensor-to-scalar ratio is very small. ",
        "introduction": "The inflationary paradigm is extremely successful in solving the horizon and flatness problems of the standard Big Bang cosmology, and at the same time in explaining the origin of structure in the observable Universe~\\cite{Liddle:2000cg}. However the problem of how to incorporate inflation into a concrete model of high energy particle physics remains unclear. On the observational side, the currently available data is not precise enough to select a particular model, whereas on the theoretical side we still lack the full understanding of the dynamics of inflation. Among the several mechanisms proposed for inflation, hybrid inflation \\cite{Linde:1991km, Copeland:1994vg} continues to be a very well motivated possibility \\cite{Linde:1997sj,Jeannerot:1997is, BasteroGil:1999fz}, especially from the point of view of constructing a model of inflation based on a supersymmetric (SUSY) extension of the Standard Model (SM)~\\cite{Dvali:1994ms}. The main advantage of hybrid inflation is that it involves small field values below the Planck scale, thereby allowing a small field expansion of the K\\\"ahler potential in the effective supergravity (SUGRA) theory, facilitating the connection with effective low energy particle physics models. In this paper we will be concerned with two problems confronting SUSY hybrid inflation, namely the $\\eta$-problem and the moduli stabilization problem, and show that their joint resolution seems to favor a particular class of models which are well motivated from both string theory and particle physics.  To be consistent with recent observations \\cite{Komatsu:2008hk, Hinshaw: 2008 ap}, implementing inflation in a quantum field theoretical context typically requires a scalar field (at least at the effective field theory level), dubbed inflaton, whose potential is extremely flat. Global supersymmetry provides a plethora of additional scalar fields, i.e.\\ the bosonic superpartners of the SM fermions. However in order to have a flat enough potential and be a viable candidate for being the inflaton, either one (or several) of the scalar fields has to be a gauge singlet or a flat direction in the multi-dimensional field space. One promising candidate naturally emerges from the supersymmetric seesaw mechanism, which provides a compelling and minimal explanation of the smallness of the observed neutrino masses. The superfields containing the SM-singlet right-handed neutrinos may be gauge singlets and their bosonic components, the right-handed sneutrinos, may in principle play the role of the inflaton in either chaotic inflation \\cite{Murayama:1992ua} or hybrid inflation \\cite{Antusch: 2005 hp}.  In the framework of SUSY models of inflation it is necessary to take into account the effects arising from the effective SUGRA theory, which turn out to be important even for low field values. Such corrections generically occur in locally supersymmetric (SUGRA) theory (including low energy effective SUGRA theories arising from string theory) and threaten the flat direction for inflation. The dangerous corrections arise from the effective SUGRA potential, where the expansion of the K\\\"ahler potential generally leads to masses of the order of the Hubble scale $H$ for all scalar fields, including the inflaton. Such corrections to the inflaton mass would lead to a slow-roll parameter $\\eta \\sim 1$ which would spoil inflation. This is the so-called $\\eta$-problem of SUGRA inflation \\cite{Copeland:1994vg, Dine:1995uk}.  Several attempts have been made to cure the $\\eta$-problem as follows. One option to treat the problem is to arrange for a small SUGRA induced mass by hand, e.g.\\ by a specific choice of the K\\\"ahler potential or by a general expansion of the K\\\"ahler potential in inverse powers of the Planck scale and finally tuning the expansion parameters (see e.g.\\ \\cite{Murayama:1993xu,Antusch: 2005 hp} for examples in the context of sneutrino inflation). However these choices are not motivated by any symmetry argument. Another way to solve the $\\eta$-problem is to impose some symmetry requirement on the K\\\"ahler potential. For example, shift symmetry has been used in several SUGRA inflation model constructions \\cite{Kawasaki:2000yn, Yamaguchi:2000vm, Brax:2005jv}. Another possibility is to use a Heisenberg symmetry\\footnote{The name Heisenberg symmetry is due to the invariance of the theory under non-compact Heisenberg group transformations. An account is given in e.g.~\\cite{Binetruy: 1987}.} in the K\\\"ahler potential which leads to a flat potential for the inflaton at tree-level~\\cite{Gaillard:1995az}. However, using a Heisenberg symmetry in order to solve the $\\eta$-problem requires the introduction of a modulus field, which must be stabilized during inflation.  The problem of combining moduli stabilization and inflation is also common to compactifications of the higher dimensional string theory. In the low energy effective 4-dimensional SUGRA theory there are several scalar fields, which are for instance related to the geometry of the internal space of the higher dimensions, and which are also called moduli fields. The requirements for the moduli fields are exactly opposite to the ones for the inflaton field. The moduli fields must be stable during inflation, in order not to spoil inflation, and must remain stabilized after inflation. The necessity of giving the moduli fields a large mass and stabilizing them at comparatively large field values is called the ``moduli stabilization problem\" in the literature \\cite{Brustein:1992nk, Dine:2000ds}.  It turns out that in potentially realistic SUGRA models it is hard to stabilize the moduli and solve the $\\eta$-problem of inflation simultaneously. The moduli sector is usually not decoupled from the inflaton sector, which means that any mechanism of moduli stabilization would always have an effect on the inflaton sector. In particular, a small variation of the moduli fields can give a significant contribution to the inflationary potential, often spoiling the conditions for inflation. Discussions of the problems associated with moduli stabilization and solving the $\\eta$-problem, as well as directions for possible solutions, can be found in the literature: For example, in \\cite{Ellis:2006ara} moduli stabilization is discussed in the context of chaotic type potentials with the no-scale form of the K\\\"ahler function. It has been noted that inflation may be achieved if the associated moduli are fixed during inflation, however at the expense of large fine-tuning. Similar conclusions have been drawn recently in \\cite{Davis:2008fv}. The possibility of modular inflation in this context has recently been discussed in~\\cite{Covi:2008cn} and it has been pointed out that subleading corrections to the no-scale K\\\"ahler potential could help to realize consistent inflation models. The effect of couplings between moduli and inflaton sectors in hybrid models has been discussed in \\cite{Brax:2006ay}. There it has been found that when the usual $\\eta$-problem in SUGRA is solved by using a shift symmetry, the moduli dynamics in general contributes a large mass to the inflaton and spoils inflation. A possible resolution to this problem has been proposed in \\cite{Davis:2008sa}, again using a shift symmetry. In the context of hybrid type of inflationary models, the possibility of using a Heisenberg symmetry for solving the $\\eta$-problem and issues connected to the associated moduli problem were discussed in \\cite{Gaillard:1998xx}, however no explicit model has been considered.   In this paper we present a class of SUSY hybrid inflation models in which the $\\eta$-problem is solved by a Heisenberg symmetry of the K\\\"ahler potential. The associated modulus field is stabilized and made heavy with the help of the large vacuum energy during inflation without any fine-tuning.\\footnote{ Here we do not address the problem of stabilizing the modulus after inflation, which we assume to be achieved by a different mechanism.} Because of the Heisenberg symmetry of the K\\\"ahler potential, the tree-level potential of the inflaton is flat and only lifted by radiative corrections, induced by Heisenberg symmetry breaking superpotential couplings. The resulting class of models are well motivated from the point of view of string theory since they include the case of no-scale SUGRA K\\\"ahler potentials which are ubiquitous in string constructions. The models are also well motivated from the point of view of particle physics since they allow the possibility that the inflaton may be identified with the right-handed sneutrino in SUSY see-saw models of neutrino masses, for example as in the model of sneutrino hybrid inflation in \\cite{Antusch: 2005 hp}. We emphasize that the general class of models considered here applies to a wider class of singlet inflaton models with Heisenberg symmetry, where the inflaton field is massless during the inflationary epoch, and subsequently acquires a large mass at the end of inflation. However much of the paper is devoted to a particular example inspired by sneutrino hybrid inflation with no-scale SUGRA, and which we discuss in some detail in order to illustrate the approach. In the considered example we demonstrate explicitly how the modulus gets stabilized by the large vacuum energy density provided during inflation, and find that the spectral index $n_s$ is predicted to be below $1$, but above about $0.98$, while the tensor-to-scalar ratio $r$ is below ${\\cal O}(0.01)$. In extensions of this minimal model we find examples where a spectral index as low as $0.95$ can be realized.  The paper is organized as follows: In section \\ref{The Model} the general framework is outlined and our explicit example scenario is presented. In section \\ref{Diagonal  Basis} we describe the background evolution of the fields for the considered example. Section \\ref{Tree-Level Scalar Potential} contains the analysis of the relevant tree-level potential including the relevant inflaton-dependent masses of the scalar fields. The flatness of the tree-level potential is lifted by the radiative corrections which are calculated in section \\ref{One-Loop Effective Potential}. Moreover, it is devoted to numerical solutions of the inflationary dynamics, including the stabilization of the modulus. The predictions of the model are presented in section \\ref{Observations}. Our Summary and Conclusions are given in section \\ref{Conclusions}. ",
        "conclusions": "In this paper we have proposed a class of models in which the $\\eta$-problem of SUSY hybrid inflation is resolved using a Heisenberg symmetry, where the associated modulus field is stabilized and made heavy with the help of the large vacuum energy during inflation without any fine-tuning. The proposed class of models is well motivated both from string theory considerations, since it includes the commonly encountered case of no-scale SUGRA K\\\"ahler potential, and from the perspective of particle physics since a natural candidate for the inflaton in this class of models is the right-handed sneutrino, i.e.\\ the superpartner of the SM-singlet right-handed neutrino, which is massless during the inflationary epoch and subsequently acquires a large mass at the end of inflation.  In order to illustrate the approach we have developed in some detail a specific example motivated by sneutrino hybrid inflation with no-scale SUGRA. In this example the right-handed sneutrino field $N$ appears in the Heisenberg combination in Eq.~\\eqref{Heisenberg}, the superpotential has the form in Eq.~\\eqref{Superpotential}, and the K\\\"ahler potential has the form in Eq.~\\eqref{Kaehlerpotential} where in practice we have assumed the no-scale form in Eq.~\\eqref{specificf(rho)}. In this model the singlet $S$ is stabilized during inflation due to its non-canonical K\\\"ahler potential, and since the right-handed sneutrino contained in $N$ has its mass protected by the Heisenberg symmetry, it becomes a natural candidate for the inflaton, with the associated modulus field stabilized and made heavy with the help of the large vacuum energy during inflation (where we have not addressed the problem of the stability of the modulus after inflation). Because of the Heisenberg symmetry the tree-level potential of the sneutrino inflaton is flat and only lifted by radiative corrections (induced by Heisenberg symmetry breaking superpotential couplings) which we have studied and found to play a key part in the inflationary dynamics.  We have found that in the considered setup the spectral index $n_s$ typically lies below $1$ with typical values shown in Fig.~\\ref{paramspace} for the case of the model defined in Eqs.~\\eqref{Superpotential} and~\\eqref{Kaehlerpotential} with $f(\\rho)$ given in Eq.~(\\ref{specificf(rho)}). However, the inclusion of additional couplings, for instance of $\\kappa_{SH}\\ne 0$ as in Eq.~\\eqref{Eq:Kaehler_general}, could lower the spectral index at the loop-level. For example, for the parameters $(\\kappa,\\lambda,\\kappa_{SH})=(0.05,0.2,10)$ we have found a spectral index of $n_{\\text{s}}\\simeq 0.953$ and $M$ fixed to $M/\\MP\\simeq 0.0029$ by the COBE normalization. We expect further changes to the predictions for the observables in other extensions or variants of our simple example model.  In conclusion, the class of SUSY hybrid inflation models proposed here not only solves the $\\eta$-problem using a Heisenberg symmetry, and stabilize the associated modulus during inflation, but are also well motivated both from string theory and particle physics, and leads to an acceptable value of the spectral index, while predicting very small tensor modes. The specific example of sneutrino hybrid inflation with no-scale SUGRA is a particularly attractive possibility which deserves further study."
    },
    "0808/0808.2049_arXiv.txt": {
        "abstract": "The internal dynamics of a dark matter structure may have the remarkable property that the local temperature in the structure depends on direction. This is parametrized by the velocity anisotropy $\\beta$ which must be zero for relaxed collisional structures, but has been shown to be non-zero in numerical simulations of dark matter structures. Here we present a method to infer the radial profile of the velocity anisotropy of the dark matter halo in a galaxy cluster from X-ray observables of the intracluster gas. This non-parametric method is based on a universal relation between the dark matter temperature and the gas temperature which is confirmed through numerical simulations. We apply this method to observational data and we find that $\\beta$ is significantly different from zero at intermediate radii. Thus we find a strong indication that dark matter is effectively collisionless on the dynamical time-scale of clusters, which implies an upper limit on the self-interaction cross-section per unit mass $\\sigma/m\\lesssim1\\,$cm$^2$g$^{-1}$. Our results may provide an independent way to determine the stellar mass density in the central regions of a relaxed cluster, as well as a test of whether a cluster is in fact relaxed. ",
        "introduction": "Our understanding of dark matter structures has increased significantly over the recent years. This progress has mainly been driven by numerical simulations which have identified a range of universalities of the dark matter structures. One of the first general properties to be suggested was the radial density profile \\citep{1996ApJ...462..563N,1998ApJ...499L...5M,2004MNRAS.353..624D,2006AJ....132.2685M,2006AJ....132.2701G}, whose radial behaviour was shown to change from a fairly shallow decline in the central region to a much steeper decline in the outer regions. This behaviour has been confirmed observationally for galaxy clusters, both through X-ray observations \\citep{2004ApJ...604..116B,2005A&A...435....1P,2005A&A...441..893A,2006ApJ...640..691V,2006A&A...446..429P}, and also more recently through strong and weak lensing observations \\citep{2004ApJ...604...88S,2005ApJ...619L.143B,2006ApJ...642...39C,2008arXiv0802.4292L}.  A slightly less intuitive quantity to be considered is the dark matter velocity anisotropy defined by \\begin{equation}\\label{eq:be} \\beta \\equiv 1 - \\frac{\\sigma^2_t}{\\sigma^2_r}, \\end{equation} where $\\sigma^2_t$ and $\\sigma^2_r$ are the 1-dimensional tangential and radial velocity dispersions in a spherical system~\\citep{1987gady.book.....B}. This anisotropy was shown in pure dark matter simulations to increase radially from zero in the central region to roughly 0.5 in the outer region~\\citep{1997ApJ...485L..13C,1996MNRAS.281..716C,2006NewA...11..333H}. For collisional systems, in contrast, the velocity anisotropy is explicitly zero in the equilibrated regions. Therefore, eventually inferring $\\beta$ from observational data is an important test of whether dark matter is in fact collisionless, as assumed in the standard model of structure formation. On this note, it has been shown that the Galactic velocity anisotropy can affect the detection rates of direct dark matter searches \\citep{2008PhRvD..77b3509V}, and it is in principle measurable in a direction-sensitive detector \\citep{2007JCAP...06..016H}.  The most massive bound structures in the Universe are clusters of galaxies, which consist of an extended dark matter halo, an X-ray emitting plasma making up the intracluster medium (ICM), and the individual galaxies. While the contribution of galaxies to the total mass is small, approximately 10 $\\%$ of the cluster mass resides in the ICM. The present generation of X-ray satellites, \\emph{XMM-Newton} and \\emph{Chandra}, allows very accurate measurements of azimuthally-averaged radial profiles of density and temperature of the ICM. These are used, under the assumption of hydrostatic equilibrium and spherical symmetry of both gas and total mass distributions, to estimate total, gas, and dark matter mass profiles \\citep{1980ApJ...241..552F}.  Below we infer the radial velocity anisotropy profile of dark matter in 16 galaxy clusters using a generally applicable framework without any parametrized modeling of the clusters. In short, we assume a universal relation between the effective temperature of dark matter and the ICM temperature, which allows us to solve the dynamics of the dark matter halo using the radial gas temperature and density profiles determined from X-ray data. We investigate the shape and validity of this temperature relation in two cosmological simulations of galaxy clusters, based on independent numerical codes. We apply our method to 16 galaxy clusters from two different samples and find a velocity anisotropy significantly different from zero in the outer parts, in qualitative agreement with simulations.  Our approach here is a generalization of the non-parametric analysis in \\citet{2007A&A...476L..37H} where $\\beta$ was inferred neglecting the radial dependence. We also note the parametrized analyses in \\citet{2004ApJ...611..175I} and \\citet{2007MNRAS.380.1521M}, where the total dark matter velocity dispersion was inferred assuming either $\\beta=0$, or the analytical $\\beta$-profiles of \\citet{2000ApJ...539..561C} or \\citet{1996MNRAS.281..716C} (see also \\citet{2008MNRAS.tmp..719W}). In particular, \\citet{2007MNRAS.380.1521M} found that the dark matter temperature and the ICM temperature were essentially the same in strong cooling-core clusters.  The structure of the paper is the following: In the next section, we discuss how we relate the temperature of dark matter to the observable gas temperature. In section 3 we show how we can then solve the dynamics of the dark matter. In section 4 we test the assumed temperature relation and our method on numerical simulations, and in section 5 we apply the method to observational data. Section 6 is the summary and discussion. ",
        "conclusions": "In this paper, we have presented a non-parametric method to infer the velocity anisotropy of dark matter in clusters of galxies from the observable temperature and density of the intracluster medium. We assume that the intracluster medium has the same specific energy as the dark matter, and we investigate the validity of this assumption in two different cosmological simulations of the formation of galaxy clusters. Both confirm the simplest possible form of the relation, namely $T_\\mathrm{DM}\\approx T_\\mathrm{gas}$ in the radial range which is resolved.  We have tested how well our method can reconstruct the actual velocity anisotropy in the simulated clusters, and we have found good agreement between the two, although the reconstruction is sensitive to systematic biases connected with deviations from hydrostatic equilibrium.  We have applied our method to the radial ICM density and temperature profiles of 16 galaxy clusters based on {\\it Chandra} and {\\it XMM-Newton} X-ray data. The shape of the velocity anisotropy profiles is always consistent with that seen in simulations, which tends to zero at the innermost radius where the temperature relation is calibrated. It then increases to about 0.5 at $r_{2500}$ and even larger in the outer regions. The same is true of the fiducial analysis applied to simulated data and is likely caused by a deviation from hydrostatic equilibrium outside $r_{2500}$. We also find a significant anisotropy even if we assume radially varying $\\kappa$-profiles, such as can be expected given the strong gas cooling and AGN heating in the core of many clusters, or if we assume $\\kappa\\ne 1$. The agreement between the observed velocity anisotropy and that predicted in numerical simulations shows that we are beginning to understand also the dynamical aspects of dark matter in halos.  In the innermost radial bins we measure a rather large anisotropy, but this is most likely an overestimation due to the neglect of the stellar mass in the center. This can be used as a means to estimate the stellar mass profile of galaxy clusters if one assumes that the velocity dispersion to be isotropic in the central regions. Similarly, our method may be used as a general test of whether a cluster is relaxed. A reconstructed velocity anisotropy which deviates significantly from the simulated profiles would be a strong hint that the data do not support the assumption of hydrostatic equilibrium.  The inferred velocity anisotropy profiles are significantly different from zero which means that the collective behaviour of dark matter is unlike that of baryonic particles in gases. This shows that dark matter is effectively collisionless on the timescale of $\\tau\\sim10^9$, the dynamical timescale of galaxy clusters. By taking typical values at $\\sim0.3\\,r_{2500}$ and allowing only a few scatterings within the time $\\tau$, this corresponds to an order--of--magnitude upper limit to the scattering cross-section of roughly $\\sigma/m=(\\rho_{\\textrm{DM}}\\tau v)^{-1}\\lesssim1\\,$cm$^2$g$^{-1}$. This limit is similar to what has been found for merging clusters \\citep{2004ApJ...606..819M,2008arXiv0806.2320B}, and within an order of magnitude of the scattering cross-section for self-interacting dark matter proposed in \\citet{Spergel:1999mh}.  We emphasize that improvements to the numerical simulations in the near future will improve our understanding of the $\\kappa$ profile and hopefully track the impact of radiative effects in the center. We also hope that improved understanding of deviations from hydrostatic equilibrium will allow us to estimate how large the suspected bias at large radii is. On the observational side, the main problem at present is the uncertainty in the temperature profile. Improvements can be expected both with regards to the deprojection analysis and the amount of data available. Obviously, there is also the possibility of including a kinematical analysis of the galaxy clusters in our method."
    },
    "0808/0808.3601_arXiv.txt": {
        "abstract": "{}{We investigate the dependence of the luminosity-diameter relation of galaxies on the environment.} {This study is based on a comparison between the 80 galaxies in the Shapley-Ames Catalog that are located within a distance of 10 Mpc, and the luminosity-diameter relation for galaxies in great clusters such as Virgo and Coma.}{ A relatively tight linear correlation is observed between the absolute magnitudes and the logarithms of the linear diameters of galaxies located within 10 Mpc. Surprisingly this observed power-law relationship appears to be almost independent of environment and local mass density as defined by Karachentsev \\& Malakov. However, at a given luminosity, early-type galaxies are (on average) slightly more compact than are objects of a later type. Unexpectedly the present results appear to indicate that the luminosity-diameter relation for the galaxies within 10 Mpc is indistinguishable from what is  observed in the much denser Virgo cluster.} { Galaxies appear to form an almost one-dimensional family in parameter space. It remains a mystery why the luminosity-diameter relation for galaxies is so insensitive to environment.}{} ",
        "introduction": "Giuricin et al. (1988) first drew attention to the surprising fact that the luminosity-diameter relation provides the tightest of all correlations observed among the photometric parameters that can be used to characterize galaxies.  Subsequently this conclusion was strengthened and confirmed by Gavazzi et al. (1996) who found a strong correlation between diameter and blue-light luminosity of galaxies, and an even tighter correlation between H-band luminosity and disk galaxy diameter.  Similarly tight relationships can be obtained by comparing photometric and dynamical parameters describing galaxies (Faber \\& Jackson 1976, Tully \\& Fisher 1977). According to present ideas the formation of galaxies is a very messy and chaotic process involving both multiple mergers and complicated feed-back effects. Furthermore the history of the rate of star formation in individual galaxies can differ in a bewildering variety of ways. From currently fashionable ideas on galaxy formation one would expect environment, initial mass, angula momentum, central concentration, star formation history and merger history to all affect the evolution of disk galaxies (Dalcanton et al. 1997, Ma et al. 1998).  This raises the question why parameters, such as galaxy luminosity and diameter, should still be so closely correlated at the end of widely differing evolutionary tracks in different environments. Possibly some hints about the origin of the observed  power-law relation between the luminosity and the diameter of galaxies might be obtained by searching for variations of the luminosity-diameter relation as a function of galaxy environment, which may itself be characterized in a variety of differing ways. Girardi et al. (1991) have pointed out that ``The tightness of this relationship suggests that possible environmental effects should be detectable from the analysis of its shape in different samples\". Giuricin  et al. (1988) claim to have detected ``appreciable differences in the galaxy-luminosity relationships for different clusters''. On the other hand Girardi et al. (1991) reach the opposite conclusion and find no significant differences between the luminosity-diameter relations in a variety of environments. Similarly Gavazzi et al. (1996) obtain broadly similar luminosity-diameterrelations in differing environments. It is the purpose of the present study to re-investigate this question by comparing the luminosity-diameter relations in  three particularly high quality data sets: (1) The Shapley-Ames galaxies with D $<$ 10 Mpc (2) the disk galaxies in the Virgo cluster, and (3) a number of other more distant Abell clusters. ",
        "conclusions": "Figure 1A shows a plot of galaxy diameter $A_{25}$ in kpc versus galaxy absolute magnitude $M_{B}$ for all of the galaxies listed in Table 1. This figure shows that, over a remarkably large range in luminosity, galaxies exhibit a reasonably tight power-law relation between luminosities and diameter. This result is both unexpected and puzzling. One might have expected the chaotic merger history of galaxies, and their widely differing histories of star formation, to have produced a broad spectrum of size-luminosity relationships. Furthermore the sizes of some galaxies are likely to have been affected by feedback produced by active galactic nuclei during early phases of their evolution. Figure 1 shows that, after reducing the Girardi et al. (1991) observations to the distance scale used in the present paper, the objects in Table 1 are well-represented by the relation\\\\  M$_{B}$ = -12.65 -5.7 log A$_{25}$,  \\hspace*{2.5cm}(1)\\\\  ~~~~~~~~~$\\pm$0.1 ~~~$\\pm$0.2     \\\\   {\\flushleft that Girardi et al. (1991) derived for 177 disk galaxies in the Virgo cluster.  It should be emphasized that the errors quoted in Eqn (1) are probably underestimated because they do not include systematic effects that might result from the bias that is introduced by incompleteness at the faint end of the galaxy sample that these authors used.}  The rms dispersion of the galaxies listed in Table 1 derived from the regression line defined by Eqn. (1) is 0.8 mag. From a similar study of 533 disk galaxies in eight relatively nearby rich clusters and in the inter-cluster ``Great Wall'' Gavazzi et al. (1996) find a slope of 6.6, which is marginally greater than the value of 5.7 $\\pm$ 0.2 that Girardi et al. (1991) found for the Virgo cluster. Gavazzi et al. argue that the power-law relationship between the radii and the masses of galaxies suggest ``that to the first order the process of galaxy evolution can be described with a single parameter; the initial mass\". It may be challenging to reconcile this conclusion with the hierarchical merging scenario of the $\\Lambda$CDM Scenario in which the initial mass concept appears meaningless.  In a different context Woo et al. (2008) have recently found that the Local Group dwarfs basically define a one-parameter ``fundamental line\" primarily driven by steller mass.  In Figure 1 distance errors will cause a galaxy to slide along a line with slope -5.0 (corresponding to constant surface brightness),which is close to the  observed slope of -5.7 in Eqn. (1). As a result even moderately large errors in distance estimates for individual galaxies will not add significantly to the observed dispersion of galaxies around the line defined by Eqn. (1). The most strongly deviant object (marked by a plus sign) is NGC 1569 (see panel \\#336 of Sandage \\& Bedke 1994), which is unusually small for its luminosity. Van den Bergh (1960) classifies it as Ir pec III-IV? It appears likely that this object is too bright for its diameter because it underwent a violent star-burst in the relatively recent past  (Greggio et al. 1998).  It is noted in passing that Grocholski et al. (2008) have recently used HST observations to show that NGC 1569 is more distant than previously believed and it probably a member of the IC 342 group.  From the data in Table 1 it is also seen that NGC 221 (M32) is smaller than would be predicted by Eqn. (1). The reason for this is, no doubt (Faber 1973), that M32 has been stripped of its outer envelope by one or more strong tidal encounters with M31.  The largest galaxy within a sphere of radius 10 Mpc is NGC 5457 (M101). This object is almost twice as large as the two next largest galaxies with D $<$ 10 Mpc: The Andromeda galaxy and NGC 4258 (M106). Peebles (2007) has drawn attention to the fact that M101 is located near the edge of the Local Void, rather than in the dense central region of the Local Super-cluster.  Figure 1B shows that field (F), group (G) and cluster (C) galaxies are all distributed similarly relative to the line defined by Equation 1. This result is slightly puzzling because van den Bergh (2007) found a very strong correlation between the morphological types of galaxies and their assignment to F, G and C environments. One might perhaps have expected cluster members (which are mainly of early morphological types) to be smaller than field galaxies that mostly have late morphological types.  Figure 1C shows that field galaxies, and objects that are bound on timescales $< 1/H_{o}$, are distributed in the same fashion. Taken at face value these results suggest that the position of an object on the luminosity-diameter relation is essentially independent of the environment in which it finds itself. This result is again somewhat counterintuitive because one might have expected compact early-type galaxies to occur predominantly among objects that are bound on a Hubble time, whereas more extended late-type galaxies would perhaps have been expected to predominate among unbound field galaxies.  Finally Figure 1D shows that early-type and late-type galaxies are displaced relative to each other. Not unexpectedly the figure shows that, at a given luminosity, early-type galaxies are more compact than galaxies of late type. In other words, the luminosity-diameter relation for early-type galaxies is slightly displaced towards smaller diameters relative to that for late- type galaxies.  These small deviations show that galaxies are almost, but not exactly, a one parameter family.  In this connection it is of interest to note that Gavazzi et al. (1996) find that the power law relationship between galaxy luminosity and diameter is even tighter in H-band than it is in the I-band.  This no doubt is to the fact that environmental factors, that may influence galaxy formation and the presence of dust will, affect the blue luminosity of a galaxy more that they will its infra-red luminosity.  One of the main result of the present investigation is that the relatively well-studied Virgo cluster exhibits a luminosity- diameter relationship that is indistinguishable from that of the lower density region containing the Shapley-Ames galaxies that are situated within a distance of 10 Mpc. In their study of the Virgo cluster Girardi et al. (1991) divided their Virgo sample into three sub-regions: An inner shell within R $\\leq$ 0.5 Mpc of the center of the Virgo cluster, an intermediate shell with 0.5 $< R \\leq$ 1.0 Mpc, and an outer shell with R $>$ 1.0 Mpc. These authors found that these three subsamples exhibited extremely similar luminosity-diameter relations. More fragmentary data that Girardi et al. collected on the diameters of other (mainly more distant) clusters do not appear to show clear-cut differences between the luminosity-diameter relation for the Virgo cluster and those for other dense clusters.  Girardi et al. (1991) have also compared the luminosity-diameter relations of objects in Tully's (1988){\\it  Nearby Galaxies Catalog}, which lists galaxies out to a redshift of 3000 km $s^{-1}$ corresponding to R $\\simeq$ 42 Mpc. Within this largere volume these authors find no significant differences between the luminosity-diameter relations for galaxies that Tully assigns to (1) the field, (2) groups, and (3) clusters. Girardi et al. therefore conclude that ''These results showed that any possible environmental effect is not strong enough to affect significantly the L-D relation for disk galaxies in the samples used in the present paper.'' The results obtained from Tully's catalog are therefore entirely consistent with those shown in Figure 1B which exhibits a similar independence from environmental effects for galaxies that van den Bergh (2007) assigned to clusters, groups and the field from visual inspection of the prints of the {\\it Palomar Sky Survey}. Furthermore these results are also consistent with the data plotted in Figure 1C, which show no obvious dependence of the luminosity-diameter relation among nearby galaxies on the $\\theta$ index of Karachentsev \\& Malakov (1999).  Finally Gavazzi et al. (1996) find no systematic difference between the luminosity-diameter relations in dense Abell clusters and among the more isolated individual galaxies in the ''Great Wall''. The fact that Gavazzi et al. find a much tighter correlation between H-band luminosity and radius than between B-band luminosity and galxy radius confirms the suspicion that environmental effects are smaller at long wavelengths."
    },
    "0808/0808.1118_arXiv.txt": {
        "abstract": "{Magnetic Fields -- Magnetohydrodynamics (MHD) -- Accretion Discs -- Stellar and Planetary Dynamos -- Planetary Magnetic Fields}  Magnetic fields are known to reside in many astrophysical objects and are now believed to be crucially important for the creation of phenomena on a wide variety of scales. However, the role of the magnetic field in the bodies that we observe has not always been clear. In certain situations, the importance of a magnetic field has been over looked on the grounds that the large-scale magnetic field was believed to be too weak to play and important role in the dynamics.  In this article I discuss some of the recent developments concerning magnetic fields in stars, planets and accretion discs. I choose to emphasise some of the situations where it has been suggested that weak magnetic fields may play a more significant role than previously thought.  At the end of the article I list some of the questions to be answered in the future. ",
        "introduction": "Our knowledge of magnetism, and of magnetic fields, began with the study of lodestones several hundred years before Christ. Stones of this type can become magnetic, and then two such stones will naturally align (see, for example, Parker 1979). Scholars of the time proposed a variety of reasons for this peculiar behaviour and it was even suggested at one point that the lodestone might even have a soul!  Many centuries later, at the start of the $17^{th}$ century, William Gilbert noted that the Earth also behaves like a large lodestone (see for example Childress \\& Gilbert, 1995). Gilbert noted that a lodestone always points to align with what is now referred to as magnetic north/south in the same way that a small stone aligns itself relative to a larger one. By the end of the nineteenth century, it had been concluded that the Earth is not unique in having a magnetic field. In 1908, Hale  determined that the Sun has a magnetic field and that the phenomena that are known as sunspots, which are dark patches on the surface of the Sun (Hale 1908), have a magnetic field that is incredibly strong (of the order of 1000 Gauss). In the last century, studies of different stars, planets and other non-Earth objects, have been made and it is now known that many of them have a detectable magnetic field. A useful summary of typical field strengths for some objects can be found in Zeldovich \\textit{et al}.\\ (1983) and Jones (2007).  Given that there is a measurable magnetic field in many objects in the Universe, it is natural for us to ask what its role is in the dynamics of different objects.  At a more fundamental level we wish to know if the magnetic field that is being measuring today is a primordial field, that has resided in the object since its creation, or if the magnetic field is one that is constantly being generated by some kind of hydromagnetic dynamo mechanism.   In the last few decades, the importance of a magnetic field in many of the astrophysical bodies has become increasingly recognised. This has stimulated considerable research in the area of magnetohydrodynamics (MHD) to understand how a plasma and a magnetic field interact. In some phenomena, such as sunspots, the role of such a  strong magnetic field is fairly easy to grasp at a basic level but only recently has a more comprehensive understanding of these phenomena been obtained (see the article by Bushby in this volume). However, it has proved to be more difficult to understand the important and role of weaker fields, where the magnetic energy is much less than the kinetic energy \\footnote{This is the definition of weak that is frequently used for astrophysical scenarios. However, we note here that is some cases the definition of a weak magnetic field is different and is one for which the gas pressure is much bigger than the magnetic pressure.}. Further, understanding why some planets have magnetic fields that are easy to detect, and others don't means that there is a vast array of issues and interesting questions still to be resolved about magnetic fields and their interactions with electrically conducting fluids in objects throughout the universe.  In this short article, I will discuss some of the recent developments in our knowledge and understanding of astrophysical magnetic fields, and discuss some of the interesting open problems and issues. Much of this revolves around the study of how astrophysical object generate magnetic fields i.e.\\ their dynamo mechanisms. ",
        "conclusions": "Over the last 100 years, since Hale made the first discovery of a magnetic field in a non-Earth object, there has been vast progress in our understanding of the interactions between astrophysical plasmas and a magnetic field. There is now a detailed picture of how solar features arise, what ingredients preclude dynamo action in stars and planets, and how turbulence in a disc gives rise to accretion of matter onto the central object. This said, there are still a plethora of questions that remain for us to answer such as:  \\begin{itemize} \\item Why is it the case that there are extended periods in history, such as the Maunder minimum where there is an absence of the `usual' sunspots that are associated with solar cycle? What is the trigger for the sunspot cycle to appear again several decades later? \\item What is it that makes Venus different, such that no dynamo mechanism operates? \\item How does a more complex internal structure within a star affect the maintenance and transport of the magnetic field? \\item Are there further astrophysical areas where the magnetic fields have so far been perceived to be weak, and so neglected, which now warrant reconsideration in the light of the vital role played by weak magnetic fields? \\item In how many stars is there a flip in the magnetic field, and on what time-scales do these flips occur? It is possible to obtain a comprehensive understanding of how the time-scale for flips relates to the structure of each star? \\end{itemize}  Hopefully in the next 100 years new considerable progress will be made on these topics and other areas with the help of new space-based missions and other technology that will improve our current knowledge from observations, and advanced in computer resources."
    },
    "0808/0808.1097_arXiv.txt": {
        "abstract": "In conventional models of gauge-mediated supersymmetry breaking, the lightest supersymmetric particle (LSP) is invariably the gravitino. However, if the supersymmetry breaking sector is strongly coupled, conformal sequestering may raise the mass of the gravitino relative to the remaining soft supersymmetry-breaking masses. In this letter, we demonstrate that such conformal dynamics in gauge-mediated theories may give rise to satisfactory neutralino dark matter while simultaneously solving the flavor and $\\mu / B \\mu$ problems. ",
        "introduction": "In supersymmetric model building, there is often a tension between the experimental constraints of flavor physics and dark matter.  A cold relic with the observed relic abundance of $\\Omega_{DM} \\simeq 0.25$ \\cite{Dunkley:2008ie, Tegmark:2006az} naturally suggests a cross section consistent with a weakly interacting particle of mass $m \\simeq 100-1000$ GeV.  For such a particle to be a good dark matter candidate, it should be stable.  It is often assumed that this can be accomplished in supersymmetry (SUSY) by imposing R-parity.  However, it is often the case that solving the flavor problem in the MSSM implies a gravitino LSP. The gravitino mass arises from gravity-mediated effects, which in general are flavor violating. Experimental constraints on flavor violation require that gravity-mediated contributions to the soft SUSY-breaking masses are small; if additional physics is responsible for generating the remaining soft masses, one might expect the gravitino to be much lighter than the other superpartners.  From this point of view, the most natural means of communicating SUSY-breaking to the MSSM would seem to be gauge mediation \\cite{Dine:1981za,Dimopoulos:1981au, Nappi:1982hm,AlvarezGaume:1981wy,Dine:1995ag}.  In gauge mediation, the MSSM soft masses are generated through gauge interactions and are therefore flavor blind.  One can then easily arrange for the scale of supersymmetry breaking to be low enough to render gravity-mediated contributions negligible. Consequently, the gravitino is generally quite light, and thus it is often stated that a gravitino LSP is one of the firm predictions of gauge mediation \\cite{Giudice:1998bp, Meade:2008wd}. Unfortunately, gravitino dark matter is somewhat untenable, both in terms of cosmology and direct detection.  There has been much ongoing work to engineer satisfactory dark matter candidates within gauge mediation \\cite{Dimopoulos:1996gy, Nomura:2001ub, Fujii:2002fv, Ibe:2006rc,Feng:2008zza, Feng:2008ya,Dudas:2008eq}.  With many direct detection searches looking for a massive particle recoil \\cite{Akerib:2005kh, Angle:2007uj}, it is unfortunate that most of these models do not predict this classic observable signature.  In models with strongly coupled hidden sectors, the na\\\"{i}ve scale of the gravity mediated contribution to soft masses is known to be inaccurate.  Indeed, in models of anomaly mediation, it is necessary to suppress gravity-mediated contributions when SUSY breaking occurs at a high scale \\cite{Randall:1998uk, Giudice:1998xp}.  Such suppression is generally accomplished by ``sequestering.'' In models with extra dimensions, sequestering is obtained by separating the SUSY breaking from the Standard Model by a large distance.  There is also a purely four-dimensional picture known as conformal sequestering \\cite{Luty:2001jh,Luty:2001zv}, in which gravity-mediated operators are suppressed by large anomalous dimensions.  For the most part, sequestering has been considered strictly in the context anomaly mediation.  However, it was recently observed \\cite{Roy:2007nz,Murayama:2007ge} that by incorporating sequestering in models with gauge mediation, the $\\mu / B\\mu$ problem may be solved.  In this letter, we will discuss the implications of sequestered gauge mediation for dark matter.  In particular, we will show that the combination of sequestering and gauge mediation can simultaneously solve the flavor problem and the $\\mu / B\\mu$ problem while furnishing a neutralino dark matter candidate with the right relic abundance.  The organization of this paper is as follows: In Sec. \\ref{sec:seq} we review sequestering and its application to gauge mediated supersymmetry breaking, while in Sec. \\ref{sec:theories} we consider a general class of conformally-sequestered theories suitable to solving $\\mu/ B \\mu$ in which the gravitino is no longer the LSP. Such theories lead to a characteristic spectrum of soft masses at the scale of conformal symmetry-breaking.  Finally, in Sec. \\ref{sec:dm}, we discuss the resultant weak-scale sparticle spectrum and explore the range of parameters for viable neutralino dark matter. ",
        "conclusions": "In this letter, we have demonstrated that neutralino dark matter is possible in gauge mediation.  Specifically, the $\\mu$ problem and flavor problem may be simultaneously solved in a simple model of gauge mediation with conformal sequestering, in which a bino-like neutralino is the lightest supersymmetric particle.  Moreover, the relic abundance of this neutralino dark matter may be consistent with cosmological observations.  In the coming years, the LHC will give us a better picture of physics beyond the Standard Model.  Likewise, dark matter direct detection experiments will increasingly probe some of the most interesting ranges of parameter space. Here we provide an example of a very general class of theories producing satisfactory signatures in both theaters: dark matter within range of direct detection experiments and a characteristic sparticle spectrum that is both consistent with current collider bounds and within reach of future experiments.  While this work was being completed, \\cite{Shirai:2008qt} appeared which uses a similar mechanism to obtain neutralino dark matter in gauge mediation.  Their model employs both a separate solution to the $\\mu$ problem and sequestering above the messenger scale to solve the flavor problem."
    },
    "0808/0808.1574_arXiv.txt": {
        "abstract": "Optical spectra of the bright Type II-L supernova SN 1979C obtained in April 2008 with the 6.5~m MMT telescope are compared with archival late-time spectra to follow the evolution of its optical emission over the age range of 11 to 29 years.  We estimate an H$\\alpha$ flux decrease of around 35\\% from 1993 to 2008 but noticeable increases in the strength of blueshifted emission of forbidden oxygen lines. While the maximum expansion of the broad $\\sim6700$ km s$^{-1}$ H$\\alpha$ emission appears largely unchanged from 1993, we find a significant narrowing of the double-peaked emission profiles in the [\\ion{O}{1}] $\\lambda\\lambda$6300, 6364 and [\\ion{O}{2}] $\\lambda\\lambda$7319, 7330 lines.  A comparison of late-time optical spectra of a few other Type II supernovae which, like SN~1979C, exhibit bright late-time X-ray, optical, and radio emissions, suggests that blueshifted double-peaked oxygen emission profiles may be a common phenomenon. Finally, detection of a faint, broad emission bump centered around 5800 \\AA\\ suggests the presence of WC type Wolf-Rayet stars in the supernova's host star cluster. ",
        "introduction": "SN 1979C was one of the brightest Type II supernovae (SNe) ever observed in the optical and radio.  It was discovered in M100 on 1979 April 19 by G.\\ Johnson \\citep{Mattei79} (NGC 4321; d = $15.2$ Mpc, \\citealt{Freedman01}), a galaxy which has had an unusually high number (five) of supernovae \\citep{Panagia80}. SN 1979C's blue absolute magnitude was $\\approx -20$ mag, about $2-3$ mag brighter than typical supernovae of its kind in the optical, and it reached a 6 cm flux density of $\\sim 8$ mJy, implying a luminosity more than 200 times that of Cas~A in the radio \\citep{Y89,Tammann90,Gaskell92,Weiler86,Weiler89}.  Optical spectra near maximum light showed a featureless continuum that developed emission lines with expansion velocities  $\\ga 8000$ \\kms\\ typical of Type II SNe  \\citep{Branch81,Barbon82}.  Dominating the spectrum after one month was strong and broad  (v~$\\approx 9000-10,800$ \\kms) H$\\alpha$ emission with little or no P Cyg absorption. The relatively slow decline and evolution of the light curve classified the supernova photometrically as a Type II-L \\citep{Panagia80}.  SN 1979C was optically recovered in 1990 over a decade after outburst (\\citealt{FM93}; hereafter FM93). Low-dispersion optical spectra showed broad but blueshifted, double-peaked and asymmetric emission lines of [\\ion{O}{1}] $\\lambda\\lambda$6300, 6364 and [\\ion{O}{2}] $\\lambda\\lambda$7319, 7330. Faint 6000 \\kms\\ broad H$\\alpha$ emission with an asymmetric profile stronger toward the blue was also seen.  Follow-up ground-based optical spectra taken in 1993 with the Multiple Mirror Telescope (MMT) and near-UV spectra ($2200-4500$ \\AA) obtained in 1997 with the Faint Object Spectrograph (FOS) aboard the {\\sl Hubble Space Telescope} ({\\sl HST}) showed these same features as well as blueshifted, double-peaked emission lines of C~II] $\\lambda\\lambda$2324, 2325, [\\ion{O}{2}] $\\lambda$2470, \\ion{Mg}{2} $\\lambda\\lambda$2796, 2803 and [\\ion{O}{3}] $\\lambda$4363 (\\citealt{Fesen99}; hereafter F99).  The blueward peak was much stronger than the less blueshifted peak in the near-UV lines, which was interpreted as due to extinction within the expanding ejecta.  Radio flux density from SN 1979C has been extensively monitored over the last three decades, with observations beginning eight days after maximum optical light.  These data show an initial decline in flux density followed by a flattening or possible brightening at an age of $\\approx10$ yr \\citep{Weiler91,Weiler01,Montes00}. The spectral index at late times has been relatively constant, although recent data suggest a probable flattening \\citep{Montes00,Bartel08}.  Potential quasi-periodic variations in its radio emission may be due to modulations of the progenitor star's circumstellar medium (CSM) by a binary companion \\citep{Weiler92,Schwartz96}.  Very Long Baseline Interferometry (VLBI) observations suggest a near-free expansion and shell structure \\citep{Bartel03,Bartel08}.  Here we present a recent optical spectrum of SN 1979C and compare it against earlier spectra in order to follow the evolution of its late-time emission. Similarities between SN~1979C and a number of other Type II supernovae observed at late times ($> 5$ yr) are also briefly discussed.  \\begin{figure*} \\centering \\includegraphics[width=0.9\\linewidth]{f1.eps} \\caption{MMT spectrum of SN 1979C taken in April 2008.  A blue continuum has been removed from the observed spectrum. } \\end{figure*} ",
        "conclusions": "We present an optical spectrum of SN 1979C taken almost three decades after outburst. This spectrum is compared against archival spectra to follow the evolution of SN~1979C's late-time optical emission.  Overall, H$\\alpha$ emission appears to have decreased by some 35\\% while the [\\ion{O}{1}], [\\ion{O}{2}], and [\\ion{O}{3}] lines show noticeable increases in blueshifted emission from levels observed in 1993.  The separation between emission peaks in the line profiles of [\\ion{O}{1}] $\\lambda\\lambda$6300, 6364 and [\\ion{O}{2}] $\\lambda\\lambda$7319, 7330 has decreased over the past 18 years, with the more blueshifted peak's velocity changing from $-5700$ \\kms \\ to $-4300$ \\kms.  We note that SN~1979C's late-time oxygen line emissions appear similar to other Type II SNe observed $5 - 30$ yr past outburst in that they show largely blueshifted emission, often with a blueshifted, double-peaked line profile. The commonality of these late-time emissions may signal a shared geometry in the oxygen ejecta and/or CSM environment for X-ray, optical, and radio bright core-collapse supernovae."
    },
    "0808/0808.3571_arXiv.txt": {
        "abstract": "In order to understand the influence of magnetic fields on the propagation properties of waves, as derived from different local helioseismology techniques, forward modeling of waves is required. Such calculations need a model in magnetohydrostatic equilibrium as initial atmosphere to propagate oscillations through it. We provide a method to construct such a model in equilibrium for a wide range of parameters to be used for simulations of artificial helioseismologic data. The method combine the advantages of self-similar solutions and current-distributed models. A set of models is developed by numerical integration of magnetohydrostatic equations from the sub-photospheric to chromospheric layers. ",
        "introduction": "In the recent years, local helioseismology has provided new insights into the sub - photospheric structure of quiet and active regions of the Sun \\citep{Duvall+etal1993, Kosovichev1999, Kosovichev2002, Kosovichev+etal2000, Zhao+Kosovichev2003, Braun+Lindsey2000}. However, the influence of magnetic fields on data interpretation has not been fully explored. Theoretical efforts have been made by \\citet{Crouch+Cally2003}, \\citet{Cally2005, Cally2006}, \\citet{Schunker+Cally2006}, \\citet{Cally+Goossens2007}, \\citet{Schunker+etal2008} in order to include mode conversion and to model ray path of waves in magnetized structures by means of analytical theory. These studies confirm the potential importance on helioseismic measurements of the so called surface effects caused by the presence of a magnetic field. A more complete understanding of the problem will, probably, be better reached via direct forward modeling of helioseismological data, since magnetic fields of arbitrary configuration may be used. There are several recent works reported in this direction as, \\eg, \\citet{Gizon+etal2006, Khomenko+Collados2006, Parchevsky+Kosovichev2007, Shelyag+etal2007, Hanasoge2008, Cameron+etal2008}. In all the works cited above \\citep[except][]{Shelyag+etal2007}, the authors apply a similar strategy. In particular, they assume the existence of an equilibrium atmosphere containing a magneto-static structure whose properties may resemble to a larger or lesser extent those of a sunspot or a magnetic flux tube. A small-amplitude perturbation is then applied to the system in order to study wave propagation, wave mode transformations, amplitude and phase behaviour of waves in a complex magnetic field topology, \\etc  The behaviour of waves observed in active regions is very sensitive to their magnetic field configuration. Photosphere and low chromosphere are regions where a small change of parameters such as the size of the magnetic structure, or its temperature, density or magnetic field strength and inclination, may produce significant changes in the resulted wave field. Simulations can be of invaluable help to explore, within a full parameter space, different magneto-static structures in order to understand the effects produced by the magnetic field on the measurable variables used in local helioseismology.  The latter task requires a robust procedure to construct magneto-static structures of desired properties. In this paper, we propose a strategy with that aim  and apply it to obtain thick structures, as prototypes for solar spots and pores. As a minimum requirement, the model should fulfill the following properties: (i) in the photosphere the model should, on average, reproduce the properties of a typical sunspot; (ii) at the border, the model should smoothly merge into a quiet-Sun non-magnetic model atmosphere; (iii) there should be a possibility to choose the profile of thermodynamic parameters at the sunspot axis; Wilson depression should be taken into account; (iv) magnetic field strength, inclination and the radius of the structure should be adjustable; (v) the model should be easily extensible into an arbitrary depth below the photosphere.  There is a vast amount of works on magneto-static models reported in the literature. Leaving apart small-scale flux tube models, those for thick structures can be divided into those possessing a current sheet \\citep[\\eg][]{Pizzo1990}, with a sharp magnetic-non magnetic interface, and those with distributed currents \\citep[\\eg][]{Pizzo1986}, showing a smooth transition. Without discussing advantages and disadvantages of the both, we will proceed here with current-distributed models.  Present current-distributed models apply two different philosophies. In the first category of models, the magnetic structure is prescribed and the distribution of thermodynamic variables is looked for to be in agreement with this structure. This is the class of self-similar models, proposed by \\citet{Schluter+Temesvary1958} and then extended by, \\eg, \\citet{Low1975, Low1980}. In the second category, the pressure distribution is prescribed as the boundary condition at the axis of the magnetic structure and in the far-away non-magnetic atmosphere. Both, pressure and magnetic field, are iteratively changed in the remaining points to reach an equilibrium situation \\citep{Pizzo1986}.  From the point of view of the requirements set above, both classes of models have advantages and disadvantages. The approach by \\citet{Pizzo1986} is more fruitful in the photosphere, since the pressure distributions of the field-free and magnetized atmospheres can be taken from observations and are relatively well known. However, for deep sub-photospheric layers, the models that can be taken as boundary conditions are scarce. More precisely, the quiet-Sun non-magnetic pressure stratification can be taken from helioseismological data, for example, from the standard solar model of \\citet{Christensen-Dalsgaard+etal1996}. As for the sunspot axis, no precise data are available \\citep[see, however][]{Zhao+Kosovichev+Duvall2001, Kosovichev2002, Couvidat+Birch+Kosovichev2006}. The model of \\citet{Pizzo1986} turns out to be very sensitive to the pressure deficit inside sunspots and the method, in general, has poor convergence if the simulation box is too deep and is very sensitive to the guess of the pressure distribution at the sunspot axis. This makes the method unsuitable for the purpose of our work.  On the other side, the procedure proposed by \\citet{Low1980} works better in deep layers, where the gas pressure dominates over the magnetic pressure. In the photosphere, where the plasma becomes magnetically-dominated, negative pressures are frequently obtained from the method of \\citet{Low1980}. It is complicated to guess the parameters of the magnetic field configuration in order to avoid this problem. At the same time, if one wishes to extend the models into the photosphere and higher layers, the magnetic field strength is limited to rather low flux tube-like values, not appropriate for sunspots \\citep[see][]{Hanasoge2008, Cameron+etal2008}.  In this paper, we take advantage of both Pizzo-like and Low-like approaches and propose a method to calculate the magneto-static equilibrium of a thick sunspot-like structure with the properties defined above. Below we describe the equations that allow to successfully merge results from both methods and show examples of MHS solutions for a wide range of parameters. The conclusions are given in the last section. ",
        "conclusions": "In this study, we propose a method to construct a magnetostatic structure with properties and size of a typical sunspot, from the deep interior to the solar surface. Previously published methods to construct such a model fail due to a poor knowledge of the thermodynamic and magnetic parameters of sunspots in sub-photospheric layers. We make use of self-similar models in deep layers and show that such models can naturally merge with models where the pressure distribution is prescribed on the axis, as well as the field-free atmosphere, allowing for a more realistic description of the atmospheric layers of sunspots. The procedure shows a rather good convergence and stability. By changing the parameters of the solution, a set of models can be produced with desired properties. We suggest that these models may be used, among others, in artificial helioseismology data simulations. Given a set of models, a parametric study can be done, investigating the influence of the topology and strength of the magnetic field of sunspots on the parameters inferred by the local helioseismology measurements in solar active regions."
    },
    "0808/0808.3092_arXiv.txt": {
        "abstract": "Theory and simulations are used to study collisionless relaxation of a gravitational $N$-body system.  It is shown that when the initial one particle distribution function satisfies the virial condition -- potential energy is minus twice the kinetic energy -- the system quickly relaxes to a metastable state described {\\it quantitatively} by the Lynden-Bell distribution with a cutoff.  If the initial distribution function does not meet the virial requirement, the system undergoes violent oscillations, resulting in a partial evaporation of mass.  The leftover particles phase separate into a core-halo structure.  The theory presented allows us to quantitatively predict the amount and the distribution of mass left in the central core, without any adjustable parameters.  On a longer time scale  $\\tau_G \\sim N$ collisionless relaxation leads to a gravothermal collapse. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.0869_arXiv.txt": {
        "abstract": "Recent works have shown that the \\TI~of km-sized near-Earth asteroids (NEAs) is more than two orders of magnitude higher than that of main belt asteroids (MBAs) with sizes (diameters) between 200 and 1,000 km. This confirms the idea that large MBAs, over hundreds millions of years, have developed a fine and thick thermally insulating regolith layer, responsible for the low values of their \\TI, whereas km-sized asteroids, having collisional lifetimes of only some millions years, have less regolith, and consequently a larger surface \\TI.  Because it is believed that regolith on asteroids forms as a result of impact processes, a better knowledge of asteroid \\TI~and its correlation with size, taxonomic type, and density can be used as an important constraint for modeling of impact processes on asteroids. However, our knowledge of asteroids' \\TIvs~is still based on few data points with NEAs covering the size range 0.1--20 km and MBAs that $>$100 km.  Here, we use IRAS infrared measurements to estimate the \\TIs~of MBAs with diameters $<$100 km and known shapes and spin vector: {\\hl filling an important size gap between the largest MBAs and the km-sized NEAs.} An update to the inverse correlation between \\TI~and diameter is presented. {\\hl For some asteroids thermophysical modelling allowed us to discriminate between the two still possible spin vector solutions derived from optical lightcurve inversion. This is important for (720) Bohlinia: our preferred solution was predicted to be the correct one by Vokrouhlick\\'y et al. (2003, Nature 425, 147) just on theoretical grounds.} ",
        "introduction": "Thermal inertia is a measure of the resistance of a material to temperature change. It is defined by $\\Gamma = \\sqrt{\\rho \\kappa c}$, where $\\kappa$ is the thermal conductivity, $\\rho$ the density and $c$ the specific heat. \\ti~is a key parameter that controls the temperature distribution over the surface of an asteroid. In the limit of zero thermal inertia the surface of an asteroid is in instantaneous equilibrium with the solar radiation and displays a prominent temperature maximum at the sub-solar point. In the realistic case of a rotating asteroid with finite thermal inertia the temperature distribution becomes more smoothed out in longitude with the afternoon hemisphere hotter than the morning one \\citep[see e.g.][and references therein]{Delbo02, Delbo04, Migo07}.  {\\hl Acquisition of temperature data (e.g. from thermal infrared observations at different wavelengths around the body's heat emission peak) over a portion of the diurnal warming/cooling cycle can be used to derive the thermal inertia of planetary surfaces by fitting a temperature curve calculated by means of a thermal model to the observed data. Asteroids surface temperatures depend also on the bodies' shapes, inclination of their spin axis and rotation rates. For those objects for which this information is available the so-called thermophysical models (TMPs) can be used to calculate infrared fluxes as function of the asteroid's albedo, thermal inertia and macroscopic roughness. Those parameters are adjusted until best fit to the data is obtained \\citep[see][and \\S \\ref{S_method} for details]{Migo07}}  Knowledge of the thermal inertia of asteroid surfaces is important for several reasons: \\begin{enumerate} \\item  \\TI~is a sensitive indicator for the presence or absence of {\\hl thermally insulating} loose material on the surface such as regolith or dust \\citep[see e.g.][]{Chris03}. The value of \\ti~depends on {\\hl regolith depth, degree of induration and particle size}, rock abundance, and exposure of solid rocks and boulders within the top few centimeters of the subsurface {\\hl(i.e. a few thermal skin depths)}. Typical values of \\ti~in (S.I. units \\tiu) are 30 for fine dusts, 50 for the lunar regolith, 400 for coarse sands {\\hl (note that a \\TI~of 400 for coarse sand assumes the presence of some atmosphere, even if as thin as the Martian one)}, and 2500 for bare solid rocks \\citep[][see also http://tes.asu.edu/TESworkshop/Mellon.pdf]{Mellon00, Spencer89, 1986Icar...66..117J}. Information about \\TI~is therefore of great importance in the design of instrumentation for lander missions to asteroids such as the Marco Polo of the European Space Agency, because it allows one to have information about the soil and sub--soil temperatures and the make up of asteroid surfaces.  \\item The presence or absence and thickness of the regolith on km--sized bodies can give hints about the internal structure of asteroids: recent work by \\citet{2007P&SS...55...70M} showed that small asteroids (with sizes $\\sim$1 km) can capture collisional debris and build up regolith if their tensile strength is not high;  \\item Thermal inertia affects the strength of the Yarkovsky effect \\citep[see][and references therein]{Bottke06} which is responsible for the gradual drifting of the orbits of km-sized asteroids and is thought to play an important role in the delivery of near-Earth asteroids (NEAs) from the main belt \\citep{Morby03}, and in the dynamical spreading of asteroid families \\citep[see][]{Bottke06}.  \\item Understanding asteroid \\TI~is important to estimate and reduce systematic errors on sizes and albedos of asteroids, when the these are determined by means of simple thermal models such as the Standard Thermal Model \\citep[STM;][]{Lebo89} neglecting the effect of the rotation of these bodies and their thermal inertia \\citep{Delbo02, Delbo04}. \\end{enumerate}    To date the value of the thermal inertia has been derived for seven large main-belt asteroids (MBAs) \\citep{1998A&A...338..340M,2004A&A...418..347M,Migo06} and six NEAs \\citep{2005Icar..179...95H, 2007Icar..188..414H, Migo04, Migo07, 2004A&A...424.1075M}.  Moreover, the mean value of \\ti~was estimated for the NEAs with multiwavelength thermal infrared data, the latter believed to be representative of the \\TI~of NEAs with sizes between 0.8 and 3.4 km \\citep{Delbo07}.  By comparing MBA and NEA \\TIs, an inverse correlation between \\ti~and asteroid diameter $D$  was derived \\citep{Delbo07} of the form: \\begin{equation} \\Gamma = d_0 D^{-\\xi}, \\label{EGamma_D} \\end{equation} {\\hl where $D$ is the diameter of a sphere with a volume equivalent to that of the asteroid shape}. Equation \\ref{EGamma_D} has also important consequences for the Yarkovsky effect, implying that that the orbital semimajor axis drift rate of MBAs due to the Yarkovsky effect is proportional to $\\sim$$D^{\\xi-1}$ \\citep{Delbo07} rather than to $D^{-1}$, the latter being the expected dependence for size independent thermal inertia. Given the small number of determined asteroid \\TIs, \\citet{Delbo07} used a unique value for $\\xi$ and $d_0$ across an interval of 4 orders in magnitude in $D$. Their best-fit values are $\\xi = 0.48 \\pm 0.04$ and $d_0 = 300 \\pm 47$, where $D$ is km and \\ti~in S.I. units (\\tiu).  However, there are several reasons to suspect that surface properties of large asteroids may be different from those of smaller bodies. In this case $\\xi$ might acquire different values in different size ranges. For example, \\citet{Bottke05} showed that asteroids with $D>$100 km and most bodies with $D>$50 km in size are likely to be primordial objects that have not suffered collisional disruption in the past 4 Gyr. These objects have resided in the asteroid belt long enough to build up a fine regolith to cause their low \\ti--values regardless of their size. Moreover, the same work has shown that objects smaller than $\\sim$30 km are statistically ejecta from the catastrophic collisional disruption of larger parent bodies. In the latter case, the more recent the smaller is an object. The surfaces of these asteroids might be systematically fresher with less mature and less thick regolith, implying higher--\\ti~values. At the smaller end of the size distribution, an unknown role might be played by the YORP effect. By increasing the rotation rate of these bodies, regolith might have been ejected from the surfaces, leading to large \\ti--values. Furthermore, our knowledge of \\ti~for asteroids $<$20 km in size is based on NEAs only. While it is believed that NEA surfaces are representative of the small ($D<20$ km) MBA surfaces, this has still to be demonstrated. Some NEAs might have suffered planetary close approach strong enough to alter their surfaces, for instance by stripping off some of the regolith \\citep[see e.g.][]{2006Icar..180..201W}. This is not the case for small MBAs. Furthermore, thermal inertia is a function of temperature \\citep[\\ti $\\propto T^{3/2}$; see e.g.][for some discussion]{Migo07, Delbo07}. This effect may lead to \\ti~offsets between cooler MBAs and hotter NEAs.  In this work we present new determination of MBA \\TI~from thermophysical modeling of data obtained by the Infrared Astronomical Satellite (IRAS). We focus on MBAs with $D<100$ km in order to fill the gap of data between NEAs and the largest MBAs and improve our understating of the relation between \\TI~and asteroid size.  In \\S \\ref{S_method} we present the method used to derive the \\TI~of MBAs from IRAS data and the selection of the targets. In \\S \\ref{S_result} we describe the results obtained for each studied asteroid. Furthermore, in \\S \\ref{S_discussion}, we discuss our novel determination of asteroid \\TIs~in the context of the aforementioned published results. ",
        "conclusions": "We derived the \\TIvs~of 10 main belt asteroids in the size range between 30 and 100 km from thermophysical modeling of IRAS data. Our results indicate that \\TI~increases with decreasing size more rapidly for main belt asteroids with diameters between 30 and 1000 km than for near-Earth asteroids smaller than 30 km. This might reflect the different regolith properties between the largest, likely primordial asteroid and the smaller ones, catastrophically disrupted and rebuilt during the age of the solar system. We also discuss the comparison between diameters from thermophysical modeling of IRAS data and SIMPS diameters for the asteroids included in this study. \\label{S_conclusions}  \\subsection*"
    },
    "0808/0808.2788_arXiv.txt": {
        "abstract": "{} {To show the importance of high-spatial resolution observations of \\HII\\ regions when compared with observations obtained with larger apertures such as ISO, we present mid-infrared spectra of two Magellanic Cloud \\HII\\ regions, N88\\,A and N160\\,A. } {We obtained mid-infrared ($8-13$\\,\\um), long-slit spectra with TIMMI2 on the ESO 3.6 m telescope. These are combined with archival spectra obtained with the Infrared Spectrograph (IRS) onboard the \\spitzer\\ Space Telescope, and are compared with the low-spatial resolution ISO-SWS data. An inventory of the spectra in terms of atomic fine-structure lines and molecular bands is presented.} {Concerning N88\\,A, an isolated \\HII\\ region with no adjacent infrared sources, the observations indicate that the line fluxes observed by ISO-SWS and \\spitzer-IRS come exclusively from the compact \\HII\\ region of about 3\\arcsec\\ in diameter.  This is not the case for N160\\,A, which has a more complex morphology. We have spectroscopically isolated for the first time the individual contributions of the three components of N160\\,A, two high-excitation blobs, A1 and A2, and the young stellar object N160\\,A-IR.  In addition, extended \\SIV\\ emission is observed with TIMMI2 and is most likely associated with the central star cluster located between A1 and A2. We show the value of these high-spatial resolution data in determining source characteristics, such as the degree of ionization of each high-excitation blob or the bolometric luminosity of the YSO. This luminosity ($2\\times10^5\\,L_\\odot$) is about one order of magnitude lower than previously estimated. For each high-excitation blob, we also determine the electron density and the elemental abundances of Ne, S, and Ar.} {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.1738_arXiv.txt": {
        "abstract": "Encircling the Milky Way at low latitudes, the Low Latitude Stream is a large stellar structure, the origin of which is as yet unknown. As part of the SEGUE survey, several photometric scans have been obtained that cross the Galactic plane, spread over a longitude range of 50\\degree~ to 203\\degree. These data allow a systematic study of the structure of the Galaxy at low latitudes, where the Low Latitude Stream resides. We apply colour-magnitude diagram fitting techniques to map the stellar (sub)structure in these regions, enabling the detection of overdensities with respect to smooth models. These detections can be used to distinguish between different models of the Low Latitude Stream, and help to shed light on the nature of the system. ",
        "introduction": "During the past decade, the structure of the Milky Way (MW) has been mapped in unprecedented detail, owing to deep, wide-field photometric surveys, such as the Sloan Digital Sky Survey \\cite[(SDSS, York et al. 2000)]{york00} and 2MASS. Apart from an improved understanding of the overall shape and stellar populations of the MW, these data have unveiled a plethora of substructures on all scales. Two examples of large structures are the stellar overdensity towards Canis Major \\cite[(CMa, Martin et al. 2004)]{martin04}, and the Low Latitude Stream (LLS), or Monoceros stream, a ring-like structure that seems to encircle the MW at low latitudes \\cite[(Newberg et al. 2002, Ibata et al. 2003)]{newberg02,ibata03}.  The origin of these structures is as yet unclear, but several hypotheses have been put forward. CMa might be (the remnant of) an accreted dwarf galaxy \\cite[(e.g. Martin et al. 2004, Mart\\'inez-Delgado et al. 2005, Bellazzini et al. 2006, Butler et al. 2007, de Jong et al. 2007)]{martin04,martinez05,bellazzini06,butler07,dejong07} with the LLS being tidal debris stripped from CMa. \\cite[Martin et al. (2005)]{martin05} and \\cite[Pe\\~narrubia et al. (2005)]{penarrubia05} have shown that dynamical models of such an accretion can indeed reproduce both the CMa overdensity as well as the LLS.  On the other hand, as these structures are located at very low Galactic latitudes, it cannot be ruled out that they are intrinsic to the disk itself \\cite[(e.g. Momany et al. 2006)]{momany06}. Another explanation for the LLS is that its stars originate in the disk, but have been pulled out of the plane by interactions with satellites \\cite[(e.g. Kazantzidis et al. 2007, Younger et al. 2008)]{kazantzidis07,younger08}.  Here we present preliminary results of the application of colour-magnitude diagram (CMD) fitting techniques to photometry at low Galactic latitudes taken as part of the SEGUE survey. These techniques, developed in \\cite[de Jong et al. (2008)]{sdssmatch}, allow us to map the 3-D distribution of stars. Fig. \\ref{fig:dataoverview} demonstrates that the resulting maps of stellar (sub)structure along the SEGUE imaging scans can help to distinguish between different models of the LLS and provide crucial constraints for further modelling endeavours.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=12cm]{lowlatsubstr_f1.ps} \\caption{ {\\it Left:} Overview of imaging scans used for this analysis, indicated as vertical lines, in Galactic coordinates. The top panel shows reddening according to the dust extinction maps from \\cite[Schlegel et al. (1998)]{sfd}, with the gray scale indicating regions with $E(B-V) >$0.1, 0.25, 0.5 and 1.0 mag. The middle and lower panels show the LLS models of \\cite[Pe\\~narrubia et al. (2005)]{penarrubia05} and \\cite[Martin et al. (2005)]{martin05}, respectively. {\\it Right:} CMDs from the scan at $l$=94\\degree. The top CMD, at $b$=3\\degree, is heavily extincted (E(B-V)=3 mag), while the bottom CMD, at $b$=30\\degree, is typical of the data used for our CMD fitting analysis.} \\label{fig:dataoverview} \\end{center} \\end{figure} ",
        "conclusions": "Colour-magnitude diagram fitting can successfully reproduce the distance distribution of stars. This is corroborated by the fact that the disk and halo model that best fits the density distributions in Fig. \\ref{fig:dens_1pop} gives densities and density ratios between the different model components that are in good agreement with previous determinations from star counts \\cite[(e.g. Siegel et al. 2002 and references therein)]{siegel02}.  After subtraction of a smooth Galaxy model, a wealth of substructure becomes visible at low Galactic latitudes (Fig. \\ref{fig:subtracted}). When searching for pieces of the LLS, possible other overdensities need to be taken into account. The scans at $l$=50\\degree~ and 70\\degree cross the Hercules-Aquila cloud \\cite[(Belokurov et al. 2007)]{hercaqui}, a large overdensity extending both above and below the plane to high latitudes at a distances between 10 and 20 kpc. Consequently, the extended overdensities seen in these scans in Fig. \\ref{fig:subtracted} are likely to be related to this structure (it appears slightly more distant in our results because Hercules-Aquila is more metal-poor than the populations used in our fits). Furthermore, the Sagittarius stream lies partly along the scan at $l$=203\\degree, and is probably partly responsible for the overdensities there. However, even with these restrictions, a large number of overdensities is left, of which at least part corresponds to the LLS. For example, the overdensities at D=10 kpc, Z=+5 kpc in the $l$=178\\degree~ and $l$=187\\degree~ scans correspond to the original detection of the Monoceros overdensity by \\cite[Newberg et al. (2002)]{newberg02}.  More accurate analysis of the distances and metallicities of the detected overdensities is needed to piece together an overview of the LLS and possible other substructures. One avenue of solving for the degeneracies between distance, age, and metallicity is by using colour-colour information to obtain metallicity estimates \\cite[(Ivezic et al. 2008)]{ivezic08}. With locations and metallicities for the overdensities in place, a detailed picture of the LLS will emerge that can help to shed light on the nature of this structure.\\\\  Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/."
    },
    "0808/0808.3484_arXiv.txt": {
        "abstract": "During the last stage of collapse of a compact object into the horizon of events, the potential energy of its surface layer decreases to a negative value below all limits. The energy-conservation law requires an appearance of a positive-valued energy to balance the decrease. We derive the internal-state properties of the ideal gas situated in an extremely strong, ultra-relativistic gravitational field and suggest to apply our result to a compact object with the radius which is slightly larger than or equal to the Schwarzschild's gravitational radius. On the surface of the object, we find that the extreme attractivity of the gravity is accompanied with an extremely high internal, heat energy. This internal energy implies a correspondingly high pressure, the gradient of which has such a behavior that it can compete with the gravity. In a more detail, we find the equation of state in the case when the magnitude of the potential-type energy of constituting gas particles is much larger than their rest energy. This equation appears to be identical with the general-relativity condition of the equilibrium between the gravity and pressure gradient. The consequences of the identity are discussed. ",
        "introduction": "\\label{sec1}  After a star with a mass lower than the Oppenheimer-Volkoff limit has spent its nuclear fuel, it becomes either the white dwarf or neutron star. The rest of its thermal energy is irradiated into the neighbouring space and, thus, its temperature permanently decreases, down to a value approaching the absolute zero. Its internal structure can be described with the help of theory of cold degenerated electron or neutron gas.  In this work, we polemize with the current theory of the supplementary case: the end-state of the star with the final mass exceeding the Oppenheimer-Volkoff limit. Namely, there are several serious arguments that such a star cannot end as a cold object. Even if it was cooled enough during the first stage of the collapse, its internal energy and temperature can be expected to increase over all limits, when its radius approaches the horizon of events, which was found at the Schwarzschild gravitational radius, $R_{g}$. More specifically, the increase of internal energy can be expected when the proper radius of the object is reduced below, say, $1.01\\, R_{g}$ or $1.001\\, R_{g}$.  Our considerations are based on the classical quantum statistics and classical general relativity (with no-hair theorems valid), which both are well-known a long time. Only new element concerning the assumptions in this work is a term corresponding with a potential-type energy in the formulas of quantum-statistics equations of state for the extremely hot gas. In fact, our idea is not new, in principle. The potential-energy term is well known in the classical Maxwell-Boltzmann impulse-distribution law of ideal gases having a low internal energy. In the latter, it is identical to the classical potential energy and situated in the Boltzmann part of the distribution law.  In our description of gas properties, we use the formula for the relativistic energy which is suitable in a situation when the rest energy of gas particles can be neglected. In the classical works by Landau \\cite{landau32}, Chandrasekhar \\cite{chandra35}, or Oppenheimer \\& Volkoff \\cite{oppenheimer:volkoff39}, the influence of conservative gravitational field on the state properties of the gas was not taken into account. A step toward an inclusion of this influence occurred in more recent studies of stellar collapse and neutron stars. Prakash et al. \\cite{prakash:etal97} considered the potential-energy term in course to provide the equation of state of a hot gas, which constitutes a neutron star a short time after its formation. The authors, however, considered that approximative form of the relativistic formula for energy, which is applicable at the values much lower than the rest energy. Their work was focused especially on the nuclear reactions and alternative forms of matter constituting the star. They did not deal with any collapse going on down to the gravitational radius. Very recently, Moustakidis \\& Panos \\cite{moustakidis:panos09}, who continued in the work of Prakash et al., provided an equation of state for a hot nuclear matter with the $\\beta$-decay equilibrium. This kind of matter was again assumed to constitute a neutron star in less extreme conditions, allowing neutrons to be unstable, than we are going to consider.\\\\  The following arguments can be presented to support the concept of the occurrence of extremely hot gas in an object with radius approaching $R_{g}$.  (1) Let us consider a radiation coming from an ultra-relativistic compact object to an observer at a relatively large distance. In his own reference frame, the observer detects the radiation which is red-shifted. In other words, its energy is lower in the reference frame of the observer than was at the moment of its emission, in the reference frame of the source. In a reverse situation, when a radiation comes onto the surface of the compact objects from a distant source, its wave-length is blue-shifted. This time, its energy is increasing with time. It is higher in the object-surface reference frame than in the frame of the distant source.  The numerical calculation reveals that the frequency of the radiation coming to the horizon at $R_{g}$ from an outer space increases over all limits implying the corresponding energy to increase over all limits. Such the increase must, however, occur not only in the case of radiation, but in the case of non-zero rest-mass particles, when these are falling down, as well. The requirement of the local conservation of energy in a volume of the compact-object-surface layer means that the increase must be, on the other hand-side, balanced by a decrease of that type energy of the surface layer, which has a character of potential energy. We refer to this energy as the \"potential-type energy\".  Therefore, when the collapsing-object radius approaches the gravitational radius, its internal thermal energy must increase over all limits, because we do not know any efficient mechanism of the extreme-rate cooling. A fraction of internal energy is irradiated in form of photons and carried out by neutrinos only from the object's surface. Even in this case, the efficiency of both cooling processes decreases down to zero when the gravitational radius is approached (see our more specific considerations in Sect.~\\ref{sec6}). Thus, an extreme accumulation of the internal energy seems to be inevitable in the last stage of the collapse to the horizon and can be the cause of a strong pressure with a steep gradient resisting to the gravity. After the object's radius reaches the horizon, the accumulated energy is trapped inside the object forever. (Some exceptions resulting from our work are discussed in Sect.~\\ref{sec7}. In the context of our work, we do not consider fine effects as, e.g., Hawking's evaporation of black holes.)  (2) Tolman \\cite{tolman30b} and Tolman \\& Ehrenfest \\cite{tolman:ehrenfest30} found that inside the sphere filled in with an ideal material fluid, the product of both proper temperature of the fluid, $T$, and square root of component $g_{44}$ of metrique tensor is a finite constant, i.e. \\begin{equation}\\label{equivT} T\\sqrt{g_{44}} = T\\exp(\\nu/2) = T_{c} = const. \\end{equation} ($\\exp(\\nu )$ is an alternative expression of $g_{44}$; see Eq.(\\ref{line_elem}) and the text below this equation). It means that if a volume of the fluid is situated near the surface of the sphere, the radius of which approaches $R_{g}$, then $g_{44} \\rightarrow 0$ and the temperature of the fluid must approach infinity the product of both could be the finite constant.  (3) At the same time, Tolman \\& Ehrenfest \\cite{tolman:ehrenfest30} demonstrated that the quantity $P_{rad}\\, g_{44}^{2}$ is equal to a finite constant inside the sphere filled in with the black-body radiation. $P_{rad}$ is the radiation pressure. When a volume of gaseous material fluid sinks to a more and more curved space-time at the gravitational radius, radiation pressure $P_{rad}$ must be magnified over all limits to keep the constancy of $P_{rad}\\, g_{44}^{2}$. The requirement of this finite constancy also supports the hot concept of final stage of very massive stars.\\\\  Considering all these facts, we search for the \"hot\" solution of the static configuration of a spherically symmetric object with no generated energy. Specifically, we give a brief summary of the theory and results that have been achieved up to date in Sect.~\\ref{sec2}. The introduction of a new modification of concerning fundamental theory with some obvious steps toward the solution of established fundamental integrals is presented in Sect.~\\ref{sec3}. In Sect.~\\ref{sec4}, we derive the solution, which is applicable on an extremely hot gas. In more detail, we derive the state quantities of the ideal gas situated in a region of extremely curved space-time of the stable spherically symmetric object, whereby the curvature is so large that the magnitude of the potential-type energy of a given particle is much greater than its rest energy. In such a gas, a radiation must necessarily occur. So, an influence of the radiation pressure on the object's stability is discussed in Sect.~\\ref{sec5}. In Sect.~\\ref{sec6}, we attempt roughly sketch a behavior of the collapse of object before it disappears below horizon. Finally, we summarize our most important results and discuss some implications toward a potential observational evidence of this kind of objects, in Sect.~\\ref{sec7}. ",
        "conclusions": "\\label{sec7}  In Sects.~\\ref{sec4} and \\ref{sec5}, we found the dependence of state quantities on the auxiliary function $\\nu$, for an extremely energetic environment. This dependence on $\\nu$ appears to be the same for both particles satisfying the Fermi-Dirac statistics and photons satisfying the Bose-Einstein statistics. In the case of the high energies, the rest mass of fermions is negligible, therefore they seem to be as massless particles as photons, and the identical dependence on $\\nu$ is not very surprising.  In the beginning of Sect.~\\ref{sec6}, we demonstrated that the behavior of the gradient of total pressure, given by Eq.(\\ref{final_sum}), is identical to the relativistic equation of state (\\ref{gradPfield}), being a simplification of Eq.(\\ref{basicequ}) in the limit of infinite energy. In this point, we can see a unification of quantum-physics description of state of ideal gas/radiation with the general-relativistic description of space-time filled in with a gas and radiation.  Although we demonstrated that the efficiency of the URMS cooling is low when the condition $\\chi \\gg W_{o}$ is valid, a certain energy still escapes away and this energy loss obviously leads to shrinking of the URMS until its proper physical radius becomes $R_{b} = R_{g}$. For $R_{b} = R_{g}$, no regular energy loss is possible and the equilibrium set by the identity of Eqs.(\\ref{gradPfield}) and (\\ref{final_sum}) is exact. This equilibrium seems to be the law of Nature, supposed by Eddington, which prevents very massive stars from the {\\em reductio ad absurdum}. According to our result, the final stage of an object with a mass exceeding the Oppenheimer-Volkoff limit is the stable object consisting of a matter and radiation and having radius exactly equal to $R_{g}$. In other words, the outer border of the object is situated just at the horizon of events.  The concept of the space-time surrounding the URMS does not differ from the classical concept of the space-time surrounding the non-rotating, Schwarzschild black hole. Also the existence of the horizon of events is predicted within both concepts. In our concept, this horizon cannot be, however, regarded as an absolute, insuperable barrier for any matter or radiation to escape. Just at the horizon, there are permanently situated the gas and radiation having the internal energy approaching infinity. In this context, it is necessary to note that the derived equations are related to a perfectly quiet fluid with no local phenomena. But in the reality, we do not know any perfectly quiet object. Also on the URMS surface, some fluctuations of quantities of state can be expected. Due to these fluctuations, some matter cannot be excluded to occur above the horizon. If this happens, an emission of a radiation can be expected. It must, of course, be extremely red-shifted, but its eventual detection cannot be excluded. In this aspect, our concept of the URMS is different to the concept of classical black holes. (We started to use the term \"URMS\", instead of the traditional \"black hole\", especially because of this reason.)  To explain the basic principles of our new concept of the extremely hot compact object more transparently, we constrained our considerations to the simplest case of spherically symmetric, non-rotating object. Despite this simplicity, we believe that the concept of URMS will be helpful to explain some observational phenomena related to the concerning objects."
    },
    "0808/0808.1953_arXiv.txt": {
        "abstract": "We present a study of the extinction, traced by the Balmer decrement, in HII regions in the dwarf galaxies NGC 1569 and NGC 4214. We find that the large-scale extinction around the most prominent HII regions in both galaxies forms a shell in which locally the intrinsic extinction can adopt relatively high values ($A_V = 0.8 - 0.9$ mag) despite the low metallicity and thus the low overall dust content. The small-scale extinction (spatial resolution $\\sim$0.3'') shows fluctuations that are most likely due to variations in the dust distribution. We compare the distribution of the extinction to that of the dust emission, traced by {\\it Spitzer} emission at  8 and 24\\mi, and to the emission of cold dust at 850\\mi. We find in general a good agreement between all tracers, expect for the 850\\mi\\ emission in NGC 4214 which is more extended than the extinction and the other emissions. Whereas in NGC 1569 the dust emission at all wavelengths is very similar, NGC 4214 shows spatial variations in the  24-to-850\\mi\\ ratio.  We furthermore compared the 24\\mi\\  and the extinction-corrected \\halpha\\ emission from HII regions in a sample of galaxies with a wide range of metallicities and found a good correlation between both emissions, independent of metallicity. We suggest that  this lack of dependence on metallicity might be due to the formation of dust shells with a relatively constant opacity, like the ones observed here,  around ionizing stars. ",
        "introduction": "Interstellar dust can be studied via its emission and also via the extinction that it causes. Each method has its advantages and difficulties. The most common way of obtaining  the extinction  of the light coming from HII regions is based on the comparison of the \\halpha\\ and H$\\beta$ recombination line fluxes. A major advantage of this method is the high spatial resolution achieved, determined by the resolution of the optical images. It is, however, difficult to derive the distribution of the dust from extinction maps because the relative distribution of the dust and the gas plays a mayor role and because maps of the \\halpha/H$\\beta$ ratio are biased towards low-extinction regions.  Alternatively, the dust can be studied via its emission, which depends on the dust amount, the type of grains and the dust temperature determined by the interstellar radiation field (ISRF). Although very different models for the interstellar dust exist (see, e.g., Zubko et al. 2004 and references therein), they generally need to include three types of grains: big grains that are in equilibrium with the ISRF, very small grains (VSGs) that are stochastically heated and are necessary to explain the mid-infrared emission and Polycyclic Aromatic Hydrocarbons (PAHs) to explain the line features in the 8 \\mi\\ range. So far there have been two major problems for studies of the dust emission: The lack of images at a high spatial resolution and the lack of submillimeter data probing the cold dust. Fortunately, the situation is improving. The  {\\it Spitzer} satellite has for the first time provided images of high spatial resolution in the mid-infrared regime where the emission is dominated by PAHs and VSGs. In the near future, the satellites {\\it Herschel} and {\\it Planck} will provide a large database for the submillimeter emission of galaxies.  We present a study of  the dust extinction (based on the Balmer decrement) and emission in two nearby, starbursting, dwarf galaxies,  NGC 1569 and NGC 4214. The goals are  to derive the small and large-scale structure of the dust  extinction and to compare it to the dust emission at various wavelengths, tracing different types of dust grains and temperatures. This study will allow us to derive information about the dust distribution and variations of the dust temperature and of grain types.   NGC 1569 and NGC 4214 have been chosen in order to study of role of the metallicity and active star formation for  the dust. Both galaxies have a low metallicity  of 12+log(O/H) = 8.2 (Kobulnicky \\& Skillman 1997, Armus et al. 1989) and a similar distance (2.2 Mpc to NGC 1569,  Israel 1988, and 2.9 Mpc to NGC 4214, Ma\\'iz-Apell\\'aniz et al. 2002). They both harbour large HII complexes, and host one (NGC 4214) or two (NGC 1569) Super Star Clusters (SSC) in the middle of a large expanding \\halpha\\ shell (Martin 1998). ",
        "conclusions": ""
    },
    "0808/0808.1481_arXiv.txt": {
        "abstract": "{We are carrying out a program of optical spectroscopy of the complete subsample of the 3CR catalog of radio sources at redshift $z <$ 0.3. The sample consists of 113 3CR sources, comprising FR~I, FR~II radio galaxies and Quasars. Complete datasets in other bands are already or will be soon available for the whole sample but the optical spectra are sparse and inhomogeneous in quality. The observations are carried out at the 3.58m Telescopio Nazionale Galileo (TNG, La Palma). More than 100 sources have been already observed. We present here the preliminary results on the analysis of the high and low resolution spectra. We found that sources can be spectroscopically characterized as: High Excitation Galaxies (HEG), Low Excitation Galaxies (LEG) and ``Relic'' AGNs. This classification is supported by the optical - radio correlations in which objects spectroscopically different follow different correlations. We conclude that AGNs with the same radio power can be fueled with different accretion properties. ``Relic'' radio-galaxies are characterized by extreme low excitation spectra that we interpret as nuclei whose activity has recently turned-off. The full spectral catalog will be made available to the scientific community. ",
        "introduction": "Radio Loud AGNs represent about the 10 per cent of active galaxies, nevertheless they are among the most powerful objects in the Universe. Since they are powerful AGNs, they represent an important laboratory to study the formation and evolution of the nuclear activity.  Our aim is the investigation of their nuclear activity via optical and radio observations.  \\begin{figure*}[t!] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{diag.eps}} \\caption{\\footnotesize The diagnostic diagrams are made of pair of emission lines ratios: the solid curve represents the separation between AGNs (above the line) and star-burst galaxies (below the line). The contour lines indicate the distribution of the SDSS galaxies from \\citet{k06}. Blue circles are High Excitation Galaxies (HEG); red crosses are Low Excitation Galaxies (LEG); green diamonds are Relic AGNs; black asterisks are unclassified AGNs. The symbol size is proportional to the $L$([O~III]) luminosity.} \\label{diag} \\end{figure*}   As emission lines presumably originate from gas photoionized by the radiation generated via the accretion process, their absolute luminosities and intensity ratios can provide information on the accretion itself and thus the physical conditions of the environment in which the jets are formed.  We focus on the homogeneous and vastly studied 3CR catalog of radio sources, being unbiased with respect to the nuclear/accretion properties. ",
        "conclusions": ""
    },
    "0808/0808.2816_arXiv.txt": {
        "abstract": "We use deep far-IR, submm, radio and X-ray imaging and mid-IR spectroscopy to explore the nature of a sample of {\\it Spitzer}-selected dust-obscured galaxies (DOGs) in GOODS-N. A sample of 79 galaxies satisfy the criteria $R-[24]>14$ (Vega) down to $S_{24}>100\\,\\mu$Jy (median flux density $S_{24}=180\\,\\mu$Jy). Twelve of these galaxies have IRS spectra available which we use to measure redshifts and classify these objects as being dominated by star formation or active galactic nuclei (AGN) activity in the mid-IR. The IRS spectra and {\\it Spitzer} photometric redshifts confirm that the DOGs lie in a tight redshift distribution around $z\\sim2$. Based on mid-IR colors, 80\\% of DOGs are likely dominated by star formation; the stacked X-ray emission from this sub-sample of DOGs is also consistent with star formation. Since only a small number of DOGs are individually detected at far-IR and submm wavelengths, we use a stacking analysis to determine the average flux from these objects and plot a composite IR (8--1000$\\,\\mu$m) spectral energy distribution (SED). The average luminosity of these star forming DOGs is $L_{\\rm{IR}}\\sim1\\times10^{12}L_{\\odot}$. We compare the average star forming DOG to the average bright ($S_{850}>5\\,$mJy) submillimeter galaxy (SMG); the $S_{24}>100\\,\\mu$Jy DOGs are 3 times more numerous but 8 times less luminous in the IR. The far-IR SED shape of DOGs is similar to that of SMGs (average dust temperature of around 30$\\,$K) but DOGs have a higher mid-IR to far-IR flux ratio. The average star formation-dominated DOG has a star formation rate of $200\\,\\rm{M_{\\odot}\\,yr^{-1}}$ which, given their space density, amounts to a contribution of $0.01\\,\\rm{M_{\\odot}\\,yr^{-1}Mpc^{-3}}$ (or 5--10\\%) to the star formation rate density at $z\\sim2$. We use the composite SED to predict the average flux of DOGs in future {\\it Herschel}/PACS $100\\,\\mu$m and SCUBA-2 $450\\,\\mu$m surveys and show that the majority of them will be detected. ",
        "introduction": "\\label{sec:intro}  Large extragalactic surveys with the {\\it Spitzer Space Telescope} (Werner et al.~2004) have revealed many high redshift objects which are bright in the mid-infrared (mid-IR) and have red mid-IR to optical colors (e.g.~Houck et al.~2005; Yan et al.~2007). The selection in the mid-IR indicates that these objects are dusty; however, without mid-IR spectra it is not clear whether the dust is heated by active galactic nuclei (AGN) or star formation activity or both. Furthermore, extrapolating from mid-IR to total IR luminosity is uncertain without good constraints spanning the far-IR dust peak (e.g.~Papovich et al.~2007; Daddi et al.~2007).  A sample of IR-luminous Dust Obscured Galaxies (DOGs), selected to have very red $R$-[24] color, was recently presented by Dey et al.~(2008, hereafter D08). From spectroscopic observations DOGs were found to have a tight redshift distribution around $z\\sim2$, very similar to that of the submillimeter selected galaxies (SMGs, e.g. Chapman et al.~2005). The space density and clustering of DOGs are also comparable to those of the bright ($S_{850}>6\\,$mJy) SMGs, suggesting that these two populations might be associated (e.g., in an evolutionary sequence; D08, Brodwin et al.~2008).  The origin of the bolometric luminosity in DOGs is uncertain. Using a similar $R$--[24] selection and an additional $R$--$K$ color criterion, Fiore et al. (2008, hereafter F08) selected a sample of faint DOGs ($S_{24}>40\\,\\mu$Jy) in the CDF-S and concluded, based on their stacked X-ray spectrum, that 80\\% are Compton-thick AGN.  In order to test the relationship between SMGs, DOGs, and the F08 Compton-thick AGN, one needs multi-wavelength data including near-IR, far-IR and submillimeter observations of a large sample of DOGs. In this paper we use the deep multi-wavelength data in the GOODS-N field (Giavalisco et al.~2004) to study a sample of faint DOGs in order to put constraints on their infrared luminosities, determine the role of AGN and star formation activity in these systems and compare them with SMGs. Our goal is to improve our understanding of the role of DOGs in massive galaxy evolution.  All magnitudes in this paper use the Vega system unless otherwise noted. We assume a standard cosmology with $H_{0}=72\\,\\rm{km}\\,\\rm{s}^{-1}\\,\\rm{Mpc}^{-1}$, $\\Omega_{\\rm{M}}=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "The average $L_{\\rm{IR}}$ derived for the DOGs using submm and far-IR measurements is a factor of $\\sim4$ times smaller than that calculated in D08 using a conversion from $S_{24}$ observed (aka $L_{8}$ rest $=\\nu L_{\\nu}|_{8\\,\\mu m}$) to $L_{\\rm{IR}}$. Part of this is because our average $L_{8}$ is lower, since we push three times deeper at $24\\,\\mu$m, and part is because of the assumed conversion between mid-IR and total IR luminosity. Based on our best-fit SED, we calculate that $L_{\\rm{IR}}/L_{8}\\simeq7$ (quartile range from Monte Carlo simulations is 5--10) which is within the lower range of conversion factors assumed in D08 ($L_{\\rm{IR}}/L_{8}$=5--15). On the other hand the SMGs from Pope et al.~(2008) with $S_{850}>5\\,$mJy have an average $L_{\\rm{IR}}/L_{8}$ of $\\sim20$. The conversion between 24$\\,\\mu$m flux and $L_{\\rm{IR}}$ is uncertain for high redshift galaxies since it relies on local galaxy templates. Observations of high redshift ULIRGs indicate an evolution in SED shapes from local ULIRGs (e.g.~Pope et al.~2006, 2008; Rigby et al.~2008). The overestimate of $L_{\\rm{IR}}$ (and SFR) using only $S_{24}$ has been noted for bright high redshift ULIRGs (e.g.~Papovich et al.~2007, Daddi et al.~2007). It is clear from our full SED fits for DOGs and SMGs that a uniform conversion cannot be applied to all high redshift ULIRGs and there must be additional parameters, other than mid-IR luminosity, which are needed to determine the total $L_{\\rm{IR}}$. Progress in determining these parameters will be facilitated with future wide-field far-IR and submm surveys with {\\it Herschel Space Observatory} and SCUBA-2, for example.  As discussed in D08, DOGs with $S_{24}>300\\,\\mu$m have similar surface densities to SMGs with $S_{850}>6\\,$mJy (Coppin et al.~2006). The number density of DOGs in GOODS-N down to $S_{24}=100\\,\\mu$Jy is 0.5$\\,$arcmin$^{-2}$, 6 times more than the shallower sample in D08. Given the average SFR from $L_{\\rm{IR}}$ for the SF DOGs ($200\\,\\rm{M_{\\odot}\\,yr^{-1}}$), we calculate a SFRD of $=0.01\\,\\rm{M_{\\odot}yr^{-1}Mpc^{-3}}$ at $z=2$. Depending on the value adopted for the total SFRD at $z=2$ (e.g.~Chary \\& Elbaz 2001; Caputi et al.~2007), SF DOGs contribute 5--10\\% of the total SFRD at $z=2$. This result appears to conflict with D08 who suggest that ($S_{24}>300\\,\\mu$Jy) DOGs contribute 25\\% of the total IR luminosity density at $z=2$. However, if we remove the AGN DOGs from the D08 sample and use the lower conversion factor of $L_{\\rm{IR}}/L_{8}\\simeq7$, this 25\\% becomes 7\\% which is consistent. The AGN DOGs will further contribute to the SFRD, however the small number of AGN DOGs in GOODS-N does not allow us to put constraints on their average $L_{\\rm{IR}}$ and SFR. While the simple DOGs selection can be used in many of the deep {\\it Spitzer} surveys to isolate large samples of high redshift ULIRGs, this selection alone does not resolve the bulk of the SFRD at $z=2$. For comparison, bright ($S_{850}>5\\,$mJy) SMGs contribute $0.02\\,\\rm{M_{\\odot}yr^{-1}Mpc^{-3}}$, with fainter ($S_{850}>2\\,$mJy) SMGs contributing $0.05\\,\\rm{M_{\\odot}yr^{-1}Mpc^{-3}}$ at $z=2$ (Wall et al.~2008). In addition to the DOG samples, several different selection criteria using {\\it Spitzer} data have been presented in the literature to isolate ULIRGs at high redshifts (e.g.~Yan et al.~2007; Farrah et al.~2008), although none of these select a population as numerous as the DOGs. Table \\ref{tab:ids} shows that 70\\% of DOGs are detected in the radio. Since DOGs by definition are faint in the optical, these radio-detected DOGs would be similar to the optically-faint radio galaxies (OFRGs) discussed in Chapman et al.~(2002, 2004).  SMGs are thought to be galaxies in the early stage of a massive merger (e.g., Conselice et al.~2003; Pope et al.~2008; Tacconi et al.~2008), and D08 propose that bright DOGs might be a later stage in the merger. In support of this, we find that $\\sim30\\%$ of SMGs meet the DOGs criteria and all 3 SMGs which have $>50\\%$ AGN contribution in the mid-IR (Pope et al.~2008) are in our sample of AGN DOGs. This implies that DOGs selection preferentially picks up the more AGN dominated SMGs, although these are among the most luminous DOGs. The average SF DOG shows additional mid-IR emission compared to the normalized SMG SED, which may be enhanced PAH emission or hot dust heated by an AGN or star formation. Regardless of the source of the mid-IR excess emission (which accounts for $<10\\%$ of the total IR luminosity), the average $L_{\\rm{IR}}$ and X-ray luminosity of the SF DOGs is several times less than that of most SMGs indicating that the average DOG is not likely to evolve from SMGs. Fig.~\\ref{fig:bzk} shows that most DOGs satisfy the {\\it BzK} selection; while they have ULIRG-like luminosities {\\it BzK} galaxies are thought to be forming stars continuously over longer timescales and do not necessarily require a major merger as as catalyst for star formation (Daddi et al.~2008). In summary, we remind the reader that this analysis is focussed on the average properties of DOGs in GOODS-N; while the average DOG is less luminous than $S_{850}>5\\,$mJy SMGs, some fraction of the DOGs are related to SMGs as shown in the 30\\% of SMGs which meet the DOGs criteria. In order to obtain a submm sample of comparable number density to the $S_{24}>100\\,\\mu$Jy DOGs sample requires a survey down to $S_{850}>3\\,$mJy (Coppin et al.~2006). This will be achieved with future deep SCUBA-2 surveys and allow for a more detailed comparison between DOGs and submm-emitting galaxies."
    },
    "0808/0808.0813_arXiv.txt": {
        "abstract": "{% RoboNet-II uses a global network of robotic telescopes to perform follow-up observations of microlensing events in the Galactic Bulge. The current network consists of three 2m telescopes located in Hawaii and Australia (owned by Las Cumbres Observatory) and the Canary Islands (owned by Liverpool John Moores University).  In future years the network will be expanded by deploying clusters of 1m telescopes in other suitable locations. A principal scientific aim of the RoboNet-II project is the detection of cool extra-solar planets by the method of gravitational microlensing. These detections will provide crucial constraints to models of planetary formation and orbital migration. RoboNet-II acts in coordination with the PLANET microlensing follow-up network and uses an optimization algorithm (``web-PLOP'') to select the targets and a distributed scheduling paradigm (eSTAR) to execute the observations.  Continuous automated assessment of the observations and anomaly detection is provided by the ARTEMiS system. } ",
        "introduction": "When a massive stellar object intercepts the line of sight of an observer and a bright background source, the light rays originating from the source are bent by the gravity of the intervening object. We call this phenomenon gravitational lensing, where the intervening object is termed the {\\it lens} and the background object the {\\it source}. Depending on the nature of the lens object and relative alignment of observer-lens-source, the lensing effect creates multiple arc-like images of the background source around the edge of the gravitational influence of the lens. If the source, lens and observer are perfectly aligned, the multiple images merge and appear as a bright ring around the lens (Chwolson 1924, Einstein 1936). This ring is usually referred to as the {\\it Einstein ring} of the lens and its radius depends on the lensing mass.  Microlensing is a special case of gravitational lensing in which the images of the source are so close to each other that they cannot be independently resolved. In this case, one observes an increase in the brightness of the background source as the lens traverses it and a gradual dimming back to the original source brightness as the lens moves away (Paczynski 1986). This change in brightness plotted versus time is called the {\\it event lightcurve}. Planets orbiting the lens may be detectable by the revealing tell-tale signs they leave on the microlensing event lightcurve (Mao \\& Paczynski 1991, Paczynski 1991).  Nowadays, microlensing is one of the methods routinely used to detect extra-solar planets and is particularly sensitive to low-mass planets (down to the mass of the Earth) orbiting a few AU from their host stars (Beaulieu et al. 2006, Bond 2004, Gould et al 2006, Han 2007, Udalski 2005). In this respect, it is complementary to the ongoing transit and radial velocity searches which are more sensitive to giant planets in close orbits (Butler 2006, Cameron et al. 2007, Fischer et al. 2007, Lister at al. 2007).  The remainder of the article is structured in the following way: Section 2 describes the network of telescopes. Section 3 gives an outline of the robotic control system that is used to operate the telescopes in automatic mode. Section 4 presents how the observations are queued and handled by autonomous agents. Our method of data acquisition and processing is discussed in section 5. Finally, section 6 briefly presents how we fit the microlensing events. We conclude with a summary of the paper in section 7. ",
        "conclusions": "Robonet-II has developed a complete architecure and supporting software to implement that architecture for the automated detection and characterization of exoplanets detected via the microlensing technqiue.  The RoboNet-1.0 pilot programme in previous seasons has returned promising results and contributed to almost all of the microlensing planetary discoveries to date. Current efforts are directed in further automation and restructuring of the scheduling, data acquisition and image processing, improving the alerting system and responses, as well as significant upgrades to the telescope engineering. These involve the deployment of new instruments and electrical safety and performance reliability upgrades.  LCOGT is in the process of expanding the robotic network of telescopes. Current plans are for 18 new 1m and 24 0.4m telescopes which are expected to be fully integrated and in operation by 2011. The microlensing search for planets, and in particular the method pioneered by the RoboNet project which can potentially make full use of the facilities in an automated way, is a science objective that can be efficiently realized with this network."
    },
    "0808/0808.0214_arXiv.txt": {
        "abstract": "We present first results from the Southern Cosmology Survey, a new multiwavelength survey of the southern sky coordinated with the Atacama Cosmology Telescope (ACT), a recently commissioned ground-based mm-band Cosmic Microwave Background experiment.  This article presents a full analysis of archival optical multi-band imaging data covering an 8 square degree region near right ascension 23 hours and declination -55 degrees, obtained by the Blanco 4-m telescope and Mosaic-II camera in late 2005. We describe the pipeline we have developed to process this large data volume, obtain accurate photometric redshifts, and detect optical clusters. Our cluster finding process uses the combination of a matched spatial filter, photometric redshift probability distributions and richness estimation. We present photometric redshifts, richness estimates, luminosities, and masses for 8 new optically-selected clusters with mass greater than $3\\times10^{14}M_{\\sun}$ at redshifts out to 0.7. We also present estimates for the expected Sunyaev-Zel'dovich effect (SZE) signal from these clusters as specific predictions for upcoming observations by ACT, the South Pole Telescope and Atacama Pathfinder Experiment. ",
        "introduction": "The new generation of high-angular resolution Cosmic Microwave Background (CMB) ground-based experiments represented by the the Atacama Cosmology Telescope (ACT) \\citep{Kosowsky06,Fowler07} and the South Pole Telescope (SPT) \\citep{SPTref} are currently targeting their observations in a common area in the southern sky that will ultimately cover several hundreds to thousands of square degrees.  These experiments will provide a blind survey of the oldest light in the Universe at wavelengths of $1-2$~mm and angular scales beyond the resolution limits of the WMAP and Planck satellites. At these arcminute angular scales, temperature fluctuations in the CMB are dominated by secondary effects arising from the formation of large-scale structure in the universe. One of the strongest effects is the imprint left by galaxy clusters though the Sunyaev Zel'dovich effect (SZE) \\citep{SZ80} in which CMB photons suffer inverse Compton scattering by the hot intracluster gas.  ACT and SPT are designed to detect the SZE, through its frequency-dependence: these experiments will measure temperature shifts of the CMB radiation corresponding to a decrement below and an increment above the ``null'' frequency around 220~GHz.  Much can be learned about the Universe from these surveys. First,  accurate systematics-free maps will allow measurement of the primary power spectrum of temperature fluctuations at all scales on which they are the dominant contribution. Second, these data sets will result in a complete census of massive clusters to arbitrarily large distances, limited only by a minimum cluster mass set largely by the instrumental sensitivity and expected to be several $10^{14}\\, M_\\odot$ \\citep{SPTref,Sehgal07}. Thanks to the relatively clean selection function as well as the redshift independence of the SZE, the cluster sample, especially the evolution of the number density of clusters with redshift, will be quite sensitive to the growth of structure in the Universe offering a potentially powerful probe of dark energy \\citep{Carlstrom-02}. Moreover, the SZE data in combination with optical, UV and X-ray observations can teach us a great deal about the detailed physics of cluster atmospheres and galaxy evolution in these dense environments.  Significant observing time and effort has been devoted to the development of techniques and the detection of galaxy clusters using large-area optical catalogs and X-ray observations. Several projects have taken advantage of large-area CCD imaging and have developed automated cluster detection schemes to produce large catalogs of clusters of galaxies \\citep[see][for example]{MaxBCG,Postman96,Postman01,NoSOCSI,NoSOCSII,NoSOCSIII,RCS} which target the relative over-abundance of galaxies over a range of redshifts. Similarly, X-ray surveys such as the {\\it ROSAT} All Sky Survey \\citep{Ebeling98,Bohringer01,Mullis03} produced catalogs with hundreds of galaxy clusters, while pointed X-ray observations have discovered systems up to $z\\simeq1.4$ \\citep{Mullis05}.  With this article we inaugurate the Southern Cosmology Survey (SCS). This project, funded by the National Science Foundation under the Partnership in Research and Education (PIRE) program, is a multiwavelength (radio, mm-band, optical, UV, and X-ray) large area survey specifically coordinated with ACT observations of the southern sky. The goal of the SCS is to maximize the scientific return from the new ground-based CMB experiments and therefore focuses on specific observational studies relevant to this science, such as the selection function of galaxy clusters across wavebands, cluster mass determination, and the establishment a \"gold\" sample of clusters for cosmology and galaxy evolution studies. Here we present results from an $\\simeq8$~deg$^2$ optical imaging survey of the southern sky that overlaps the common SZ survey region. The purpose of this paper is twofold: (1) to present the details of our data reduction pipeline and analysis software and (2) identify new galaxy clusters, constrain their redshifts and masses, and predict their SZ signals. Photometric redshifts come from the 4-band imaging data, while our mass estimates are inferred from the optical luminosity ($L_{200}$) and richness (\\n200) of the clusters, using relations calibrated by the Sloan Digital Sky Survey (SDSS). For the eight massive clusters, out of 38 identified (37 are new sources) in the survey area, we present positions, richness estimates, masses, and predictions for the integrated Compton $y$-distortion of the SZE using empirical power-law relations based on $N$-body simulations. Throughout this paper we assume a flat cosmology with $H_0=100 h$~km~s$^{-1}$~Mpc$^{-1}$, $h=0.7$ and matter density $\\Omega_m=0.3$. ",
        "conclusions": "In this article we have laid out the techniques and methods for our analysis of a large multi-band optical survey with the Blanco telescope and Mosaic-II instrument under the aegis of the Southern Cosmology Survey.  We have obtained sub-arcsecond astrometric precision and sufficient photometric accuracy for the estimation of redshifts to $\\delta z < 0.1$.  Forty-two optical cluster candidates were identified from an area of the sky covering $\\simeq$8~deg$^2$; the richness and integrated galaxy luminosity of the clusters were measured.  Based on correlations between these optical observables and cluster mass as established by SDSS cluster surveys \\citep{Johnston07, Reyes08} we provide mass estimates (in addition to positions and redshifts) for 8 new clusters whose inferred masses lie above $3\\times 10^{14} M_{\\sun}$. These clusters are all likely to be detected by ACT and SPT if these experiments reach their expected sensitivity levels. Although the uncertainties on the estimated mass and inferred SZE signal are large (factors of 2 or so), the accuracy of our cluster positions and redshifts are quite good and typical for 4-band imaging survey data. Moreover, it is worth noting that the $Y-M$ relation varies quite a bit from simulation to simulation (differences of a factor of 2 in $M$ at fixed $Y$), so our predicted masses are well within the uncertainty range of the latest theoretical predictions. In an effort to reduce the mass errors we have begun to estimate weak lensing masses from the current Blanco data \\citep{McInnes09}. We also note that significant additional areas of the BCS are now publicly available and will be presented in future publications.  The strength of the SCS is its multi-wavelength aspect and large sky area coverage.  The 23hr region of the sky analyzed here has now been surveyed in the UV by GALEX and in the X-ray band by {\\it XMM-Newton}; over the next several months as these data become available, we will incorporate them into the SCS.  Correlation analyses of these multi-wavelength data should allow us to reduce the large errors on inferred cluster masses and study the cluster selection biases across wavebands.  Secure confirmation of all candidates, as well as additional mass constraints, rest on follow-up optical spectroscopy, which we are pursuing at the Southern African Large Telescope and elsewhere.  Finally, once ACT and SPT data become available, the rich optical data set analyzed here will be a valuable source for understanding and quantifying the impact of large scale structure on secondary anisotropies in the CMB."
    },
    "0808/0808.1555_arXiv.txt": {
        "abstract": "{We analyzed HST images of 31 nearby (z $\\lesssim$ 0.1) 3CR radio-galaxies.  We compared their UV and optical images to detect evidence of recent star formation. Six objects were excluded because they are highly nucleated or had very low UV count rates.  After subtracting the emission from their nuclei and/or jets, 12 of the remaining 25 objects, presenting an UV/optical colors NUV - r $<$ 5.4, are potential star-forming candidates. Considering the contamination from other AGN-related processes (UV emission lines, nebular continuum, and scattered nuclear light), there are 6 remaining star-forming ``blue'' galaxies.  We then divide the radio galaxies, on the basis of the radio morphology, radio power, and diagnostic optical line ratios, into low and high excitation galaxies, LEG and HEG. While there is no correlation between the FR type (or radio power) and color, the FR type is clearly related to the spectroscopic type. In fact, all HEG (with one possible exception) show morphological evidence of recent star formation in UV compact knots, extended over 5-20 kpc. Conversely, there is only 1 ``blue'' LEG out of 19, including in this class also FR~I galaxies.  The picture that emerges, considering color, UV, optical, and dust morphology, is that only in HEG recent star formation is associated with these relatively powerful AGN, which are most likely triggered by a recent, major, wet merger. Conversely, in LEG galaxies the fraction of actively star-forming objects is not enhanced with respect to quiescent galaxies. The AGN activity in these sources can be probably self-sustained by their hot interstellar medium. ",
        "introduction": "\\label{introduction}  Both observational and theoretical studies have supported the idea of a co-evolution between supermassive black holes (SMBH) and their host galaxies. The observational evidence is derived from several pieces of evidence: the dynamical signature of the widespread presence of SMBH in galaxies (e.g \\citealt{kormendy95}; \\citealt{richstone98}; \\citealt{kormendy01}), the relation between SMBH mass and the spheroid mass (\\citealt{magorrian98}; \\citealt{mclure02}; \\citealt{marconi03}; \\citealt{haring04}), stellar velocity dispersion (\\citealt{ferrarese00}; \\citealt{gebhardt00}; \\citealt{tremaine02}) and concentration index (\\citealt{graham01}; \\citealt{graham07}). These relations were interpreted by e.g. \\citet{hopkins07c} as various projections of a fundamental plane, which included the SMBH mass, analogous to a similar relation for elliptical galaxies (\\citealt{dressler87}; \\citealt{djorgovski87}); this appears to be a sign of a common evolutionary process involving SMBH and galaxies.  The growth of a SMBH can occur via either gas accretion or coalescence with another SMBH during a merger event. Similarly, the growth of a galaxy can be associated with the capture of the stellar populations of another galaxy, due to either a merger or to star formation, which might be self-sustained or triggered by a merger.  In this framework, nuclear activity (i.e. the manifestation of gas accretion onto a SMBH) and star formation are expected to be related, and for both processes mergers are likely to play a crucial role. Indeed, the hierarchical galaxy evolution scenarios support the idea that gas flows associated with galaxy mergers trigger both starburst and AGN activity (\\citealt{kauffmann00}; \\citealt{dimatteo05}). A connection between mergers, galactic starburst, and AGN activity has been well-established and modeled (e.g. \\citealt{barnes91,barnes96}; \\citealt{mihos94,mihos96}; \\citealt{springel05}; Kapferer et al. 2005). \\citet{hopkins06a} presented a simulation in which starburst, AGN activity, and SMBH growth were connected by an evolutionary sequence, due to mergers between gas-rich galaxies. Observationally, studies of ultraluminous infrared galaxies (ULIRGs), distant submillimeter galaxies (SMGs), and quasars demonstrate that they are the remnants of major mergers, in which massive starbursts occur in combination with a dust enshrouded AGN (\\citealt{barnes96}; \\citealt{schweizer98}; \\citealt{jogee06}; \\citealt{sanders88a,sanders88c}; \\citealt{sanders96}; \\citealt{dasyra07}).  Studies of QSO host galaxies also revealed the widespread presence of a young stellar population (e.g. \\citealt{brotherton99}; \\citealt{kauffmann03}; \\citealt{sanchez04}; \\citealt{zakamska06}). In Seyfert galaxies, evidence of nuclear, dusty, compact starbursts was found (\\citealt{heckman97}; \\citealt{gonzalez98}). Apparently, AGN activity and star formation are connected also from a quantitative point of view, since the most luminous quasars have the youngest host stellar populations (\\citealt{jahnke04a}; \\citealt{vandenb06}) and the most significant post-merger tidal features and disturbances (e.g. \\citealt{canalizo01}; \\citealt{hutchings03}; \\citealt{letawe07}).  Most of studies focused on the general AGN population when radio-quiet AGN are not studied in isolation. However, there have been plentiful studies of analysis of radio-loud AGN, which are the focus of this paper.  Observations of radio galaxies in the local Universe (\\citealt{heckman86}; \\citealt{smith89}; \\citealt{colina95}) detected morphological features such as double nuclei, arcs, tails, bridges, shells ripples, and tidal plumes. In a substantial subset of radio galaxies, this suggests that the AGN activity is triggered by the accretion of gas during major mergers and/or tidal interactions. Furthermore, gas kinematics studies (\\citealt{tadhunter89} and \\citealt{baum92}) supported the interpretation that mergers can represent the triggering process of SMBH/galaxies co-evolution. Studies of the spectral energy distribution (SEDs) of radio galaxies discovered that young stellar populations, indicative of a recent starburst, provide a significant contribution to the optical/UV continua, up to 25-40 \\%, of powerful galaxies at low and intermediate redshift (\\citealt{lilly84}; \\citealt{smith89}; \\citealt{tadhunter96}; \\citealt{melnick97}; \\citealt{aretxaga01}; \\citealt{odea01}; \\citealt{tadhunter02}; \\citealt{wills02}; \\citealt{allen02}; \\citealt{wills04}; \\citealt{raimann05}; \\citealt{tadhunter05}; \\citealt{holt07}).  Apparently, AGN activity initiates at a later stage of the merger event than star formation \\citep{emonts06}, which agrees with predications of numerical simulations \\citep{hopkins07}, although the precise length of the delay is not well constrained.  However, there are still several open questions, such as: (i) what is the fraction of nearby radio galaxies showing evidence for young stars?; (ii) is there a relationship between AGN properties and galaxy star-formation characteristics?; (iii) what is the role of mergers in triggering both AGN activity and star formation in radio-galaxies?  We consider these issues by adopting an alternative approach based on the analysis of UV images with the aim of detecting a young stellar population. More specifically, we analyzed HST/STIS observations of a sample of nearby radio galaxies from the 3CR catalogue. The flux level in the UV band ($\\sim$ 2500 \\AA), relative to the optical emission, is sensitive to very low levels of star formation, as demonstrated by GALEX data (the GALEX/NUV band is similar to the filter used to acquire the STIS images used in this study), by detecting fraction of as low as 1-3 \\% of young stars formed within the last billion years (\\citealt{silk77} and \\citealt{somerville00}). The HST UV images were already used for this purpose by \\citet{odea01}, who studied the radio-galaxy 3C~236. The sample considered is sufficiently large (31 objects) for statistical conclusions and is representative of the overall population of nearby radio galaxies.  Furthermore, \\citet{schawinski06} recently analyzed GALEX and SDSS images of a significant number of non-active early-type galaxies. This study provides a well defined control sample that we can use as a benchmark for the star formation properties of quiescent galaxies to be compared with the measurements we will derive for radio-galaxies. We note that \\citet{schawinski06} focused only on quiescent galaxies to avoid contamination by UV light from an active nucleus.  For the sample that we consider here, the light produced by AGN activity (UV nuclei, and jets) can however be masked by taking advantage of the high resolution of the HST images.  The outline of this paper is the following. In Sect. \\ref{observations and data reduction}, we describe the properties of the sample and HST observations, on which this study is based, as well as the reduction method used. In Sect. \\ref{results}, we present our results by classifying the galaxies in terms of their NUV-r color profiles and UV morphology.  In Sect. \\ref{discussion}, we discuss the link between AGN and star formation properties, while in Sect. \\ref{summary} we summarize our findings.  The cosmological parameters used in this paper are H$_{0}$ = 75 \\kms\\ Mpc$^{-1}$ and q$_{0}$ = 0.5. ",
        "conclusions": "\\label{summary}  We have analyzed images of 3CR radio galaxies for which both optical and UV HST images are available. The sample is composed of 31 radio-galaxies, all but one with z $<$ 0.1, representing $\\sim$ 60 \\% of all 3CR sources below this redshift limit. We have excluded six objects, four are very highly nucleated and two with data of too low signal-to-noise ratio, and have 25 remaining objects. To perform a rigorous comparison with previous studies based on GALEX (NUV band) and SDSS (r band) observations, we derived a cross-calibration between GALEX, SDSS, and HST colors. We conservatively estimated that the error associated with this procedure was $\\lesssim$ 0.2 mag.  On the basis of the integrated colors and UV morphology, the galaxies of the sample can be divided into 3 main categories: (1) the quiescent red galaxies (14 objects) have ``red'' colors, with only diffuse UV emission, tracing the optical light; (2) the ``blue'' UV-clumpy galaxies (7 objects) show a clumpy UV morphology (extended over 5-20 kpc); (3) the UV-disky galaxies (4 objects) have UV emission that is cospatial with their circum-nuclear (on a scale of 0.5-1 kpc) dusty disks.  To recognize the presence of recent star formation, it is necessary to account for other sources of UV emission. We estimated the level of UV contamination from emission lines and nebular continuum. In UV-disky galaxies, we found that UV emission lines contributed $\\sim$ 50 - 80 \\% of the total UV flux within the region coincident with the dusty disks, causing the observed blue NUV-r color. For the 7 UV-clumpy galaxies, the emission line or nebular continuum contamination were both negligible. There was a spatial association between UV light and radio-axis only in one galaxy, which is a possible sign of contamination from scattered nuclear light.  Summarizing, we confirm the presence of a young stellar population in at least 6 of the 7 UV-clumpy galaxies, in full agreement (on an object-by-object comparison) with results obtained for ground-based optical spectroscopy.  The NUV-r color is not simply related to the radio power since, while the bluest galaxies have all relatively large values of P$_{178 \\rm MHz}$, at the high end of radio luminosities we found many of the reddest galaxies. Concerning the FR type, we found no blue galaxy among the FR~I sources, but among the FR~II (or transiction FRI/II) galaxies there was an almost equal number of ``blue'' and ``red'' galaxies.  The clearer association between galaxy color and AGN type is related to the optical spectroscopic classification into HEG and LEG: all six HEG are blue, while there is only 1 blue LEG out of 19 including all FR~I into this class.  The fraction of star forming HEG is far in excess of that found by \\citet{schawinski06} for their sample of quiescent galaxies, while in LEG there appear to be even less star-forming galaxies than in quiescent galaxies. These results can be summarized as follows: \\begin{enumerate} \\item   {While all HEG have a blue UV-r color, the opposite is true for LEG, with only one exception;}  \\item {UV, optical, and dust morphology of HEG are highly chaotic and correspond to unsettled morphologies. In LEG, the UV emission is typically diffuse and traces the optical light (leaving aside the UV-disky galaxies contaminated by line emission), while dust (when present) is mostly arranged in small circum-nuclear disk structures;}  \\item   {HEG have a higher dust content;}  \\item   {HEG have higher nuclear luminosity.} \\end{enumerate}  All of these findings can be explained if for HEG we are seeing the effects of a recent, major, ``wet'' merger. The fresh input of gas and dust causes both the higher star formation rate and stronger nuclear activity.  The most favorable situation for the occurrence of a ``wet'' merger is in groups of galaxies. This can explain the different environment between HEG, usually located in groups, and LEG, more often found in clusters of galaxies.  Conversely, in LEG we did not find evidence for recent star formation (with only one exception). A quantitative comparison with the star-formation fraction of quiescent galaxies is not straightforward due to the relatively small size of our sample and the cross-calibration uncertainties. However, we note that the substantial tail of blue objects found in inactive early-type galaxies is apparently not present for our objects, and this strongly disfavours the possibility of enhanced star formation in LEG hosts.  It appears that there is no clear connection between the recent star formation and the presence of low excitation AGN.  This is in agreement with the idea that in LEG the hot interstellar medium is able to substain the AGN activity via quasi-spherical accretion. This process does not require an external cold gas supply that might become detectable in a star formation event.  The results presented here demonstrate the potential of these studies for investigating the triggering mechanism of nuclear activity and star formation in radio galaxies, a method that can be adopted also for other classes of AGN. However, this approach can not provide a quantitative estimate of the star formation history, an essential ingredient in the study of the coupling between the growth of galaxies and SMBH. This crucial information can only be derived by complementing the existing UV images with further data, e.g. optical and UV spectra, that can be used to constraints the age of the recently formed stellar population."
    },
    "0808/0808.3766_arXiv.txt": {
        "abstract": "We carry out a comprehensive theoretical examination of the relationship between the spatial distribution of optical transients and the properties of their progenitor stars. By constructing analytic models of star-forming galaxies and the evolution of stellar populations within them, we are able to place constraints on candidate progenitors for core-collapse supernovae (SNe), long-duration gamma ray bursts, and supernovae Ia.  In particular we first construct models of spiral galaxies  that reproduce observations of core-collapse SNe, and we use these models to constrain the minimum mass for SNe Ic progenitors to $\\approx 25\\msol$.  Secondly, we lay out the parameters of a dwarf irregular galaxy model, which we use to show that the progenitors of long-duration gamma-ray bursts are likely to have masses above $\\approx 43\\msol.$ Finally, we introduce a new method for constraining the time scale associated with SNe Ia and apply it to our spiral galaxy models to show how observations can better be analyzed to discriminate between the leading progenitor models for these objects. ",
        "introduction": "Transients are any events that exhibit a brightening, dimming, or otherwise noticeable change within a finite and usually short lifetime. They are one of the most important observable objects in astronomy as they represent a major change of state, which often exerts an enormous impact on its surroundings (\\eg McKee \\& Ostriker 1977; Dekel \\& Silk 1986; Gehrz \\etal 1998; Matteucci \\& Recchi 2001).  Furthermore, these events inform a great deal on the physics and evolution of the progenitor stars that cause them.  Yet, despite their enormous importance, the progenitors of many transient types remain poorly constrained.  One would like to be able to make predictions about these events and watch them unfold in real time, but transients, as their name suggests, are fleeting and fickle - many observed only accidentally and often only partially.  In particular, in the case of supernovae (SNe) and gamma ray bursts (GRBs), observers must rely either on extremely demanding  direct methods or  approximate indirect methods to constrain their progenitor systems.  In the direct approach, one seeks to identify the progenitor star responsible for a particular transient on a case by case basis.  This, in turn, requires one of two things: (i) the occurrence of a rare, nearby event, as in the case of SN1987A (Gilmozzi \\etal 1987; Kirshner \\etal 1987) in which the surprising properties of the progenitor (\\eg Arnett \\etal 1989) coupled with neutrino detections (Bionta \\etal 1987; Hirata \\etal 1987; Mayle \\etal 1987) led to important advances in our understanding of stellar physics; or (ii) a laborious search through archival images to extract the one in $\\approx 10^{12}$  stars per year that results in a supernova (\\eg Mannucci \\etal 2005).  Even with this  second, ``brute-force\" approach, a good deal of serendipity is still necessary.  Thus, until recently, only one object (1993J; Aldering \\etal 1994), had been identified from this approach; the situation improving to a handful of objects with the advent of painstaking searches though pre-explosion Hubble Space Telescope (HST) images (Barth \\etal 1996; Van Dyk \\etal 2002; Smartt \\etal 2004; Maund \\etal 2005; Hendry \\etal 2006; Li \\etal 2006; Gal-Yam \\etal 2007).  A second approach relies on indirect constraints and uses the properties of transient host galaxies, or the locations of transient within these  hosts, to derive progenitor properties.  Studies of this type include measurements of the rates of supernovae (\\eg Pain \\etal 1996; Cappellaro \\etal 1999; Tonry \\etal  2003; Blanc  \\etal 2004; Dahl\\' en \\etal 2004; Sullivan \\etal 2006; Barris \\& Tonry 2006; Neill \\etal 2006; Botticella \\etal 2008;  Dahl\\' en \\etal 2008) and gamma ray bursts (\\eg Daigne \\etal 2006; Charay \\etal 2007; Liang \\etal 2007; Guetta \\& Della Valle 2007; Kistler \\etal 2008),  constraints from the metal content of galaxy clusters (\\eg Renzini 1999; Loewenstein 2001; Scannapieco \\etal 2003; Tozzi \\etal 2003; Baumgartner \\etal 2005) and high-redshift quasars (Barth \\etal 2003; Dietrich \\etal 2003), studies of the distribution of SNe relative to spiral   arms (Maza \\& van den Bergh 1976; Della Valle \\& Livio 1994; Bartunov \\etal 1994;  McMillan \\& Ciardullo 1996; Petrosian \\etal 2005), and studies that relate the presence and properties of transients to the environmental properties of their hosts (\\eg Oemler \\& Tinsley 1979;  Wang \\etal 1997; Hamuy \\etal 1996; Fruchter \\etal 1999; Howell 2001; Le Floc'h 2003;  Christensen \\etal 2004; van den Bergh \\etal 2005; Ostlin \\etal 2008).  Recently, Fruchter \\etal (2006, hereafter F06) developed a new such indirect observational method, building on the framework first laid out by Baade (1944) and elaborated on for GRBs by Bloom \\etal (2002), and they successfully applied it to differentiate between  the progenitors of long-duration gamma-ray bursts and supernovae.  The method involves observing the spatial locations of GRBs and type-II SNe in their host galaxies and calculating the fraction of the total host galaxy light contained in pixels fainter than these locations.  After several such observations, a pattern emerges, and F06 demonstrated that long-duration GRBs have a higher propensity to cluster in the brighter regions of a galaxy than type II SNe. Kelly \\etal (2007) expanded upon this analysis and applied it to type Ia and Ic observations, to find that type-Ic SNe have similar environments to GRBs.  Theoretical work has followed to use the F06 method to put constraints on which stars might contribute to GRB observations (Larsson \\etal 2007). In this paper, we show how the F06 method, combined with simple analytic models can be used to derive strong constraints on the progenitors of a wide range of transients.  In particular, we simulate model galaxies that act as hosts for three varieties of transients:  core-collapse SNe (type II's and Ic's), type Ia SNe, and long-duration gamma-ray bursts. Using star-formation rates and stellar evolution profiles, we construct simulated galaxies with all the features of a dynamical system of stellar births and deaths and demonstrate that the transients are primarily determined by their progenitor life-times, and thus progenitors can be constrained by the locations of the transients. By simulating the locations of stellar deaths for massive stars that result in core-collapse SNe and long-duration GRBs, we are able to place constraints on the progenitor masses.  And by simulating the location of low mass stars that result  in white dwarfs and type Ia SNe, we are able to constrain the time scale for these transient events, as well as suggest ways that the F06 method  can be modified to provide further constraints.  The structure of this paper is as follows. In \\S 2, we lay out our approach for simulating transients in model spiral galaxies, carry out tests on the resultant data, and apply it to derive mass constraints for core-collapse SN of various types. In \\S 3, we lay out the parameters and approach for a simulated dIrr galaxy, establish and test the criteria it must follow, and apply the method to derive mass constraints for long-duration gamma-ray bursts. In \\S 4, we return to our spiral galaxy model, employ analytical models for the distributions of type Ia SNe in order to constrain the associated dynamical time-scale, and suggest a modified method to improve these constraints.  Our conclusions are summarized in \\S 5. ",
        "conclusions": "The F06 method is a valuable tool in the search for transient progenitors. By relying on simple observations, the method remains agnostic to the details of the transient events, such as spectroscopic signature or peak intensity, and instead places concise and testable constraints on their spatial distributions.  In this paper, we have demonstrated first, that models of spiral galaxies illustrate the dynamics that lead to the observed spatial distributions of core-collapse SNe. In a spiral galaxy, the rotating star-formation wave is the primary dynamical component for setting the spatial distributions of SNe with varying progenitor masses. The distances that stars are able to travel from this wave during their lifetime determines how F06 profiles for certain transient events will present themselves. Since there is a direct relationship between stellar lifetime and mass, our analysis has constrained type Ic SNe only to stars above $\\approx 25\\msol$ and has supported estimates for the minimum mass of type II SNe of $\\approx 8\\msol$.  In the dwarf irregular model, the role of the star-forming wave is played by the star-formation profile which, in conjunction with the cluster diffusion, places young stars in the brightest portions of the galaxy. By employing GALAXEV results in combination with observed dependencies on SFR and the IMF, our dwarf irregular model fulfills our initial criteria for being a representative counterpart of real galaxies as well as adhering to the observed integrated light profiles of dwarf irregulars. In combining the F06 method with the results of this model, we are able to place constraints on the progenitor masses of GRBs to stars with masses above $\\approx 43\\msol$. Our dIrr models also confirm a minimum progenitor mass for type II SNe of $\\approx 7.7\\msol$.  In future work, we will improve our simulation by adding a reasonable amount of patchy absorption in the dwarf and spiral galaxies. With too much dust, the GRB results would not conform to observational constraints, and so, these models might also serve to constrain the effects of dust in GRB hosts.  For type Ia SNe, the delay time associated with the dynamical component of the two-channel Ia intensity profile is the primary factor in determining their spatial distribution. However, the effect of changing the delay time for the dynamical component is not evident in a standard F06 plot. For this reason, we have suggested a modification to this analysis that normalizes each pixel by overall radial profile of the host.  Applying this method to our simulated galaxies  demonstrates a clear dependence of the resultant profile on the delay time of the dynamical component. Future application of this method to existing observations of type Ia SNe will serve to place concise constraints on the delay time, supporting or refuting candidate progenitor models for SNe Ia.  For each of the objects we have considered in this paper, we have demonstrated that galactic dynamics play an important role in the distributions of stars at the moment of their deaths.  Furthermore, while not considered here, this techniques is likely to be applicable to other types of objects such as classical novae, Wolf-Rayet stars, and planetary nebulae.  From the observed distributions of transients, we are able to constrain the masses and properties of the progenitors from which they are formed. As transients flicker in and out of existence, the physics of stars is written across the night sky."
    },
    "0808/0808.3285_arXiv.txt": {
        "abstract": "Using archival imaging from the Wide Field Planetary Camera 2 aboard the Hubble Space Telescope, we investigate the stellar populations of the Local Group dwarf spheroidal Andromeda V -  a companion satellite galaxy of M31.  The color-magnitude diagram (CMD) extends from above the first ascent red giant branch (RGB) tip to approximately one magnitude below the horizontal branch (HB). The steep well-defined RGB is indicative of a metal-poor system while the HB is populated predominantly redward of the RR Lyrae instability strip. Utilizing Galactic globular cluster fiducial sequences as a reference, we calculate a mean metallicity of $[Fe/H] = -2.20 \\pm 0.15$ and a distance of $(m-M)_0 = 24.57 \\pm 0.04$ after adopting a reddening of $E(B-V) = 0.16$. This metal abundance places And V squarely in the absolute magnitude - metallicity diagram for dwarf spheroidal galaxies. In addition, if we attribute the entire error-corrected color spread of the RGB stars to an abundance spread, we estimate a range of $\\sim$0.5 dex in the metallicities of And V stars. Our analysis of the variable star population of And V reveals the presence of 28 potential variables.  Of these, at least 10 are almost certainly RR Lyrae stars based on their time sequence photometry. \\footnotetext[1]{Based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute} ",
        "introduction": "Dwarf spheroidal and dwarf irregular galaxies are thought to be instrumental in the process that forms larger galaxies (Font et al. 2006, and references therein). As such, their importance is sometimes considered only within this context - that of much more massive systems. However, it is important to keep in mind that dwarf galaxies (DGs) are useful probes of galaxy formation and evolution in and of themselves. In this regard, the relation between the absolute magnitude of a DG and its mean metallicity has provided a number of useful insights. First studied decades ago (Tinsley 1978; Mould, Kristian, \\& Da Costa 1983), this relation shows that more luminous DGs possess a more metal-rich mean abundance. This in turn suggests that self-enrichment by heavy elements is more likely in a system with a larger gravitational potential which can retain supernova ejecta (e.g. Davidge et al. 2002).  Within the context of this mass-metallicity correlation, Andromeda V, a dwarf spheroidal companion galaxy to M31, is somewhat of an anomaly. In their discovery paper, Armandroff et al. (1998) used Milky Way globular cluster fiducials combined with V and I imaging to measure a mean metallicity for And V of $[Fe/H] = -1.5$.  At its measured absolute magnitude of $M_V \\sim -9$, we expect And V to have $[Fe/H] \\sim -2$, similar to the Milky Way dSphs Ursa Minor and Draco, so the Armandroff et al. (1998) value made And V a significant outlier among the Local Group dwarf spheroidal galaxies in the relation between absolute magnitude and metallicity. This led Caldwell (1999) to hypothesize that perhaps And V has a deeper than normal potential well, which could be verified through measurements of its stellar velocity dispersion or mass-to-light ratio.  Motivated by the discrepancy between the expected and observed properties of And V, Davidge et al. (2002) re-examined the question of And V's chemical composition. They used the Gemini Multi-Object Spectrograph (GMOS) in imaging mode on the Gemini North telescope to construct an optical color-magnitude diagram (CMD) in the g', r', and i' filters. The slope of the And V red giant branch (RGB) was then used to calculate the mean metallicity of [Fe/H] = $-2.2 \\pm 0.1$ placing And V squarely on the $M_V - [Fe/H]$ relation.  Davidge et al. (2002) therefore asserted that there was nothing unusual about the metallicity of And V and that the absolute integrated magnitude and metallicity do indeed follow the relation for dwarf spheroidal galaxies. However, this result was again thrown into doubt with the work of McConnachie et al. (2005).  Using Johnson V and Gunn i photometry from the Isaac Newton Telescope Wide Field Camera they calculated $[Fe/H] = -1.6$ from the mean color of the And V RGB.  This once again returned And V  to the status of an outlier in the $M_V - [Fe/H]$ relation.  In light of the uncertainty about the mean abundance of stars in And V, and the overall importance of characterizing the properties of stars in dwarf spheroidal galaxies, we have undertaken a photometric study of And V using archival imaging from the Hubble Space Telescope (HST) Wide FIeld Planetary Camera 2 (WFPC2). In section 2 we detail our observations and reductions, while section 3 presents the color-magnitude diagram (CMD) of And V along with a discussion of its properties; our conclusions are summarized in Sec. 4. ",
        "conclusions": "This work presents HST WFPC2 F450W and F555W observations of Andromeda V, a dwarf spheroidal galaxy satellite of M31.  Comparing the CMD of And V to Milky Way Globular cluster fiducials of known metallicity, we find a mean metal abundance for And V of $[Fe/H] = -2.20 \\pm 0.15$ with a spread of $\\sim$0.5 dex, and a distance modulus of $(m-M)_0 = 24.57 \\pm 0.04$.  This result puts Andromeda V squarely on the absolute magnitude - metallicity relationship for local group dwarf spheroidals.  As suggested by Davidage et al. (2004) this tightens the relationship between absolute magnitude and metallicity, suggesting that there is little scatter in the relationship between mass-to-light ratio and absolute magnitude for dwarf spheroidals.  In addition we find RR Lyraes in And V providing evidence for an old population, but are unable to accurately measure their periods."
    },
    "0808/0808.1249_arXiv.txt": {
        "abstract": "The low-mass protostellar region L1527 is unusual because it contains observable abundances of unsaturated carbon-chain molecules including C$_{\\rm n}$H radicals, H$_{2}$C$_{\\rm n}$ carbenes, cyanopolyynes, and the negative ions C$_{4}$H$^{-}$ and C$_{6}$H$^{-}$, all of which are more associated with cold cores than with  protostellar regions.  Sakai et al. suggested that these molecules are formed in L1527 from the chemical precursor methane, which evaporates from the grains during the heat-up of the region.  With the gas-phase  osu.03.2008 network extended to include negative ions of the families  C$_{\\rm n}^{-}$, and C$_{\\rm n}$H$^{-}$, as well as the newly detected C$_{3}$N$^{-}$, we modeled the chemistry that occurs following methane evaporation at $T \\approx$ 25-30 K.  We are able to reproduce most of the observed molecular abundances in L1527 at a time of $\\approx~5 \\times10^{3}$ yr.   At later times, the overall abundance of anions become greater than that of electrons, which has an impact on many organic species and ions. The anion-to-neutral ratio in our calculation is in good agreement with observation for C$_{6}$H$^{-}$ but exceeds the observed ratio by more than three orders of magnitude for C$_{4}$H$^{-}$. In order to explain this difference, further investigation is needed on the rate coefficients for electron attachment and other reactions regarding anions. ",
        "introduction": "L1527, an envelope of a low-mass star-forming region with IRAS04368+2557 at its center,  has the physical features of a class 0/I protostar, as discussed by \\citet{andre00}. The infalling collapse of circumstellar gas was observed by \\citet{ohashi97}, and the gas outflow was observed by \\citet{tamura96}.  A number of interstellar molecules were discovered in L1527 as part of a molecular survey of dense cores by \\citet{jorgensen04}.  Although one might expect the chemistry to be that of a hot corino \\citep{cecc07}, the temperature of the envelope is $\\approx$ 30 K, well below that of hot corinos (100 K).  More recently, \\citet{sakai08a} detected unsaturated hydrocarbons and cyanopolyynes in the central core region of L1527. These species are unusual in corinos, and are more normally associated with cold ($T \\approx 10$~K) dark clouds such as TMC-1 \\citep{suzuki92,smith04}. \\citet{sakai08a} explained the production of the carbon-chain species in terms of the evaporation of methane during the warm-up to 30 K followed by a rapid gas-phase synthetic chemistry, a process that they termed ``warm carbon-chain chemistry.'' This chemistry is related to a grain mantle desorption model for TMC-1 of \\citet{markwick00}. In a recent paper, \\citet{hassel08} used the osu gas-grain code to study the warm-up chemistry of L1527 in the vicinity of 30 K and found that the carbon-chain abundances could be reasonably reproduced in this ``lukewarm corino'' by a chemistry in which gas-phase methane is the precursor.  In addition to these unusual species, the two negative ion species C$_{4}$H$^{-}$ and C$_{6}$H$^{-}$ were also detected in L1527 \\citep{sakai07,sakai08b,agundez08}.  Anions in the interstellar medium were first discussed by \\citet{dalgarno73}. \\citet{herbst81} proposed that neutral radicals with large enough electron affinities can undergo efficient radiative attachment to form anions in cold dense clouds if they possess more than $\\approx 4~$ atoms.  \\citet{lepp88}  suggested that negatively charged polycyclic aromatic hydrocarbons (PAHs) could also form by radiative attachment.  \\citet{wakelam08} modeled the gas-phase chemistry in dense cold cores with neutral and negatively charged PAHs, and found that the inclusion of PAH's allows good agreement with observations for small species with the use of so-called ``high-metal'' elemental abundances.  Confirmation of anions in interstellar and circumstellar sources awaited the measurement of their rotational spectra in the laboratory by Thaddeus and co-workers \\citep{mccarthy06,gupta07}. In addition to the detections in L1527,  C$_{6}$H$^{-}$ and C$_{8}$H$^{-}$ \\citep{mccarthy06,brunken07} were detected in TMC-1, while C$_{4}$H$^{-}$ \\citep{cernicharo07}, C$_{6}$H$^{-}$ \\citep{mccarthy06,kasai07}, and C$_{8}$H$^{-}$\\citep{remijan07} were observed in IRC+10216. In response to these detections, the anions C$_{n}^{-}$, n=5-10, and C$_{n}$H$^{-}$, n=4-10, were added to the RATE06 (udfa.net) and IRC+10216 chemical networks \\citep{millar00} by \\citet{millar07} with estimated or calculated rate coefficients  \\citep{terz00} for radiative attachment and photodetachment processes as well as measured rate coefficients for associative detachment reactions.  The augmented models were then used to study anionic abundances in TMC-1, IRC+10216, and photon-dominated regions such as the Horsehead Nebula.  \\cite{millar07} found that, as predicted earlier by \\citet{herbst81},  hydrocarbon radical anions can reach abundances exceeding 1\\% of their neutral precursors, depending on their size. Another type of anion, C$_{3}$N$^{-}$, the radiative attachment of which was first studied by \\citet{petrie97}, was also detected recently in IRC+10216 by \\citet{thaddeus08}.  Its radiative attachment rate coefficient has been calculated anew by \\citet{ho08}, with a result similar to that of \\citet{petrie97} used here.  In this paper, we re-study the gas-phase chemistry of L1527 with particular interest in the formation and depletion of anions and how their chemistry affects the overall  chemistry of the lukewarm corino.  The osu.03.2008 gas-phase network (see http://www.physics.ohio-state.edu/\\verb+~+eric/) has been augmented to include formation and depletion reactions for the carbon cluster and  hydrocarbon radical anions as well as for C$_{3}$N$^{-}$.   The remainder of this paper is organized as follows.  In Section~\\ref{code}, we discuss the gas-phase code used and the choice of initial abundances.  In Section~\\ref{res}, the results are discussed and compared with observation for both anions and other species and with the earlier gas-grain results for non-anionic species.  In Section~\\ref{var}, we consider how our results respond to variations in initial abundances and to the exclusion of anionic chemistry.  Finally, we end with a summary. ",
        "conclusions": "L1527 is an unusual source because the unsaturated carbon-chain molecules detected there are more normally associated with cold dense cores at lower temperatures.  \\citet{sakai08a} proposed that the carbon-chain chemistry starts soon after rising temperatures allow the volatile species methane to evaporate into the gas.  Our group has studied this chemistry via two types of models.  In a previous model, \\citet{hassel08} used the osu gas-grain code to study the chemistry in a cold dense core followed by warm-up to temperatures associated with L1527 ($T \\approx 30$ K) and found a rich chemistry leading to carbon-chain species, as hypothesized by \\citet{sakai08a}.  Comparison with observation showed that although the non-carbon-chain species in L1527 could be fit well either as remnants of a previous cold phase or current denizens of the warm-up phase, the carbon-chain species are fit significantly better with the warm model.  Although the gas-grain approach was quite useful, this method does not contain any anionic chemistry, and the chemistry is fixed by what happens in the earlier cold epoch and the rising temperatures.  We have followed this prior study with the isothermal ($T = 30$~K)   gas-phase study reported here because it allows us to include the gas-phase anionic chemistry without worrying about what happens when anions collide with grains and because it allows us to vary some of the initial abundances, which are mainly based on the earlier model.  The abundance of methane, for example, has been studied in the ice phase of cold sources \\citep{oberg08} prior to evaporation and  found to depend somewhat on the source investigated.  Including the chemistry of carbon-chain anions in our gas-phase network  has led to a large increase in the number of reactions, mainly because of the many cation-anion neutralization reactions needed.    With the chosen initial abundances, the best agreement with all 23 observed molecules in L1527 occurs at a time of  $4.8 \\times 10^{3}$ yr, when the abundances of most carbon-chain species are increasing.   This agreement corresponds to an average factor of three discrepancy between observational and calculated values.  Variation in the initial abundances of either atomic oxgyen or methane changes the time of best agreement mainly because these species regulate both the formation and depletion rates of carbon-chain species.  For the observed anions, C$_{4}$H$^{-}$ and C$_{6}$H$^{-}$, we fit both the fractional abundance and the anion-to-neutral ratio reasonably well for the larger anion, but dramatically overproduce these quantities for the smaller anion.  Finally, it is to be noted that the steady-state abundances in the gas-phase model are reached by 10$^{6}$ yr, so that the chemical processes of interest in this paper are completed well before the gas-grain models show declines in the gas-phase abundances \\citep{hassel08}.  A comparison of the chemistry occurring with and without negative ions shows that anions have a major effect on the chemistry of other species, especially at later times than that for the optimal fit, when the overall anionic abundance exceeds that of electrons.  This effect can either lower abundances of a given class of molecules or raise them, depending upon whether the formation of anions tends to deplete other species or enhance their production.  There is also a strange effect in which the dependence of the abundance of a molecule in a family such as the cyanopolyynes on size is changed when anions are included in the chemistry.  Typically, in the absence of anions, an increase in size results in a drop in peak abundance.  The inclusion of anions appears to make this dependence significantly smaller and even to vitiate it completely in some instances.  As will be shown in subsequent work, this effect is also  pronounced in the cold core TMC-1 under oxygen-rich conditions.  With relatively large abundances for the largest carbon-containing species in our network, it may well be necessary to add still larger molecules to the model, either by including all of the formation and depletion reactions that govern their chemistry, as in the IRC+10216 network \\citep{millar00}, or by some more approximate approach \\citep{herbst91}.  In addition, despite the size of the current network, the role of negative ion-neutral reactions remains to be fully explored \\citep{cordiner08,cordmill08}."
    },
    "0808/0808.1280_arXiv.txt": {
        "abstract": "We study gravitational instabilities in disks, with special attention to the most massive clumps that form because they are expected to be the progenitors of globular-type clusters.  The maximum unstable mass is set by rotation and depends only on the surface density and orbital frequency of the disk. We propose that the formation of massive clusters is related to this largest scale in galaxies not stabilized by rotation. Using data from the literature, we predict that globular-like clusters can form in nuclear starburst disks and protogalactic disks but not in typical spiral galaxies, in agreement with observations. ",
        "introduction": "The study of instabilities in disks has a long history, following the seminal work of Toomre (1964), with a considerable literature on many aspects of it.  However, relatively little attention has been given to the most massive agglomerations that can form by the fragmentation of galactic gas disks.  These most massive agglomerations are of interest because they may be the precursors of the most massive star clusters known, the globular clusters.  Globular clusters were until recently viewed as exclusively old objects, and cluster formation models were therefore based on ideas about early stages of galaxy formation.  This view changed with the realization that elliptical galaxies often contain two populations of globular clusters that appear to have different origins, suggested to be a `primordial' population and a `merger' population (Ashman \\& Zepf 1992).  The existence of a merger population was confirmed by the Hubble Space Telescope with the discovery of massive young clusters in merging galaxies (Ashman \\& Zepf 1998; Schweizer 1998).  A modern theory of globular cluster formation should therefore address why globular cluster formation is common in some environments such as merging and high-redshift galaxies but not in others such as the present Milky Way.  Another question to be addressed is why some regions of galaxies are more favorable for the formation of massive clusters than others.  For example, the inner part of the Milky Way galaxy contains young clusters with masses up to several times $10^4\\,\\msun$, including some quite close to the central black hole (Krabbe et al. 1995; Figer 2008), that are an order of magnitude more massive than the typical open clusters found elsewhere in the galaxy.  The Giant Molecular Clouds (GMCs) in which open clusters currently form in the outer Milky Way have masses of order $10^6\\,\\msun$, but they produce clusters with masses of only several times $10^3\\,\\msun$.  The formation of globular clusters with masses of $10^5 - 10^6\\,\\msun$ therefore requires either much more massive super-GMCs (Harris \\& Pudritz 1994) or a much higher star formation efficiency.  In reality, a combination of these effects is probably involved.  Several scenarios have been proposed for the formation of massive super-GMCs, involving for example a hot primordial plasma (Fall \\& Rees 1985), collisions between normal GMCs (Harris \\& Pudritz 1994), or confinement by high pressures in mergers (Ashman \\& Zepf 2001).  Relatively little attention has been given to the possible role of disks in the formation of globular clusters, with the exception of Larson (1988, 1996), but this possibility now seems worth further investigation given the evidence for massive cluster formation in rotationally supported disks in merging galaxies (Ashman \\& Zepf 1998).  In this $Letter$ we study the gravitational instability of galactic gas disks, with special attention to the most massive clumps that can form in them.  We also consider which galaxies or environments will produce the largest unstable clumps and are therefore most favorable for massive cluster formation.  We start by reviewing the theory of instabilities in disks in \\S 2, and continue with its application to  the formation of globular-type clusters in \\S 3.  We discuss the most promising environments for globular-cluster formation in \\S 3.1, and in \\S 4 we summarize the main implications of our work. ",
        "conclusions": "In this paper we have studied gravitational instabilities in disks, with special attention to the most massive clumps that form because they are expected to be the progenitors of globular-type clusters.  The maximum unstable mass is set by rotation and depends only on the surface density and orbital frequency of the disk, unlike other mass scales such as the Jeans mass that depend on the complex gas physics. This mass scale is due to a purely local instability and not dependent on large-scale ordered motion. The maximum unstable mass is therefore a well-defined quantity even if the interstellar medium is not well described by a simple equation of state.  The maximum clump mass can be expressed in terms of the total gas mass and the gas fraction in a galaxy, and this formulation makes it clear that environments with a high gas fraction are the most promising places to form massive clusters.  Using data from the literature, we predict that massive globular-like clusters can form in nuclear starburst disks and protogalactic disks but not in typical spiral galaxies, in agreement with observations.  The scenario proposed here relates massive clusters to the largest scale in galaxies not stabilized by rotation, which is the only scale intermediate between stars and galaxies that has a clear physical basis.  There is no well-defined `Jeans mass' on these intermediate scales, and the next smaller scale that has any clear physical basis is the thermal Jeans scale in molecular clouds, which is related to the masses of individual stars.  The next larger physical scale is that of the galaxy itself, so we predict three physically well-motivated scales corresponding to stars, massive stellar  clusters, and galaxies.  Another application of studies of the stability of nuclear gas disks, not discussed here, is to their likely role in feeding supermassive central black holes in galaxies (Escala 2007).  This requires the outward transfer of angular momentum, and if gravity is the most important force involved, the same mass concentrations that form massive clusters may also play an important role in the outward transfer of angular momentum and the growth of a central black hole.  We plan to investigate this problem further in ongoing work."
    },
    "0808/0808.0881_arXiv.txt": {
        "abstract": "\\noindent We examine the extent to which the self-annihilation of supersymmetric neutralino dark matter, as well as light dark matter, influences the rate of heating, ionisation and Lyman-$\\alpha$ pumping of interstellar hydrogen and helium and the extent to which this is manifested in the 21\\,cm global background signal. We fully consider the enhancements to the annihilation rate from DM halos and substructures within them. We find that the influence of such structures can result in significant changes in the differential brightness temperature, $\\delta T_b$. The changes at redsfhits $z<25$ are likely to be undetectable due to the presence of the astrophysical signal; however, in the most favourable cases, deviations in $\\delta T_b$, relative to its value in the absence of self-annihilating DM, of up to $\\simeq20\\,{\\rm mK}$ at $z=30$ can occur. Thus we conclude that, in order to exclude these models, experiments measuring the global 21\\,cm signal, such as EDGES and CORE, will need to reduce the systematics at 50\\,MHz to below 20\\,mK. ",
        "introduction": "\\noindent The standard cosmological model, motivated by measurements of temperature anisotropies in the Cosmic Microwave Background (CMB) \\citep{Spergel:2003cb,Spergel:2006hy,Komatsu:2008hk,Dunkley:2008ie}, the large scale distribution of galaxies \\citep{Cole:2005sx,Tegmark:2006az}, and by evidence of the accelerated expansion of the Universe from supernova observations \\citep{Astier:2005qq,WoodVasey:2007jb}, requires that the Universe possesses a flat spatial geometry with a corresponding critical density, approximately 27 percent of which consists of physical matter. However these observations also indicate that only 4 percent of this matter is baryonic in nature, implying that the remaining 23 percent consists of an elusive, non-baryonic component called dark matter (DM) owing to the severe constraints that current astronomical data sets on its radiative capabilities.  Despite this compelling evidence for the existence of DM, its precise nature is still a topic of debate. Particle physicists have independently supported DM by postulating the existence of a variety of exotic particles with wide-ranging properties that may potentially solve problems in particle physics whilst resulting in a thermal relic particle density that is consistent with current observational constraints.  The most intensely studied DM candidate is the lightest neutralino \\citep*{bertone_review}, a weakly-interacting massive particle (WIMP) motivated by supersymmetric extensions of the Standard Model of particle physics. In many of these extensions the neutralino is the lightest supersymmetric particle (LSP). In theories where the LSP is stable, for example theories where R-Parity is a conserved quantum number \\citep*{Weinberg:1981wj, Hall:1983id,Allanach:1999ic}, the neutralino is thus a highly-motivated DM candidate. Furthermore, an attractive feature of  neutralinos is that a large region of the relevant supersymmetric parameter space can be investigated using CERN's Large Hadron Collider (LHC)\\footnote{www.cern.ch/LHC}.  Whilst neutralino DM is ``cold'', owing to its negligible free-streaming length (i.e. the length scale below which fluctuations in DM density are suppressed), warm DM (WDM) is typically lighter and possesses a much longer free-streaming length. WDM is a viable alternative to cold dark matter (CDM) models which may potentially resolve several shortfalls of the standard CDM model, such as for example the over-prediction of low mass satellites and the existence of cuspy halos \\citep*{hogan2000,Dalcanton:2000hn,avilareese,colin2008}. Among WDM candidates, there are sterile neutrinos \\citep*{Dodelson:1993je,Asaka:2005an,Asaka:2005pn}, majorons \\citep*{Akhmedov:1992hh,Berezinsky:1993fm,Lattanzi:2007ux} and light DM (LDM) particles \\citep{Boehm:2003hm}. What makes LDM interesting for this study is the fact that it can self-annihilate, as opposed to other forms of WDM, and therefore its annihilation rate can be enhanced by overdensities.  In this paper,  we re-examine the influence of neutralino and LDM annihilations on the thermal history of the Universe at times between the epochs of recombination and reionisation, commonly referred to as the ``Dark Ages'', when gas existed in a nearly uniform, dark, neutral state. The investigation of the Dark Ages is one of the frontiers of modern cosmology, and will be carried on by a new generation of radio interferometers such as LOFAR\\footnote{ http://www.lofar.org }, MWA\\footnote{ http://www.haystack.mit.edu/ast/arrays/mwa }, 21\\,CMA\\footnote{ http://web.phys.cmu.edu/past/ }, and SKA\\footnote{ http://www.skatelescope.org }, as well as single antenna experiments such as EDGES\\footnote{ http://www.haystack.mit.edu/ast/arrays/Edges/index.html } and CORE.  These experiments will look for the redshifted 21\\,cm signal associated with the hyperfine triplet-singlet transition of neutral hydrogen. If DM annihilates or decays, the resulting products subsequently collide and heat the surrounding gas, increasing its kinetic temperature and ionisation fraction. This is in turn manifested as distinct features in the 21\\,cm background signal that can be used to constrain the properties of DM \\cite{furlanetto,Shchekinov:2006eb,valdes}.  With few exceptions (e.g.~\\cite{Chuzhoy:2007fg,Yuan:2009xq}), past studies proclaim that the heating effects associated with the annihilation of SUSY WIMP DM are too small to be detected by current radio interferometers. However, these studies overlook the enhancements to the DM annihilation rate in galactic halos and in their substructures \\citep{taylor2003}, which could be large enough to make the DM signature detectable by the next generation radio telescopes. In this paper, we calculate the effect of neutralino and LDM annihilations on the 21\\,cm signal when accounting for the effect of DM clustering.\\\\    \\noindent The rest of this paper is organized as follows. In \\S\\,\\ref{sec:Basic_Physics} we elaborate on the basic properties of neutralinos and LDM. We also discuss the basic physics describing the way in which energy from annihilations is injected into the intergalactic medium (IGM). In \\S\\,\\ref{sec:EG_DM_ann_rate} and \\S\\,\\ref{sec:DM_ann_rate} we calculate the enhancement in the DM annihilation rate caused by the presence of halos and their substructures. In \\S\\,\\ref{sec:Absorbed_Fraction} we estimate how much of the energy produced in a single DM annihilation is actually injected into the IGM. In \\S\\,\\ref{sec:21cm_Background} we discuss the modifications to the differential equations describing the evolution of the ionised fraction and kinetic temperature of the IGM and subsequently use these equations to calculate the modified 21\\,cm background. In \\S\\,\\ref{sec:Results} we calculate the predicted 21\\,cm background for our benchmark neutralino and LDM models and discuss the potential for a detection. Finally, in \\S\\,\\ref{sec:Discussion} we summarise our results and draw our conclusions. ",
        "conclusions": "\\label{sec:Discussion}  \\noindent We have calculated predictions for the effects on the evolution of the cosmological HI 21\\,cm signal during the Dark Ages for various forms of annihilating neutralino and light dark matter. In doing so, we fully accounted for the significant enhancements made to the annihilation rate of these DM particles arising from DM structures. We utilised results from the state-the-art N-body simulations to calculate the evolution of the aforementioned enhancement in the annihilation rate, referred to here as the ``clumping factor'', owing to the distribution of halos and the near self-similar distribution of increasingly smaller substructures predicted to exist within them. We did this for a diverse range of values of astrophysical parameters consistent with the uncertainties in the dynamics of the simulated halos. We performed detailed calculations of the absorbed fraction of the energy injected into the IGM by the annihilation products of our DM candidates. We used the standard equations for the evolution of the kinetic and spin temperatures of the IGM with modifications to account for the additional energy injected into by the IGM from DM annihilations. Finally, we calculated the resulting deviations in the evolution of the differential brightness temperature $\\delta T_b$ relative to a scenario where DM is absent.  In our calculations of $\\delta T_b$ we have neglected the influence of astrophysical processes affecting the 21~cm background. In fact, these effects dominate the 21~cm signal, thus obscuring the cosmological information, once star formation becomes important at redshifts $z \\simeq 25$ \\cite{Pritchard:2008da}. This means that the effect of the presence of annihilating DM at $z\\lesssim 25$ (corresponding to frequencies $\\nu \\gtrsim 55$~Hz) will likely be undetectable due to the uncertainity in the astrophysical modelization.  The global 21\\,cm signal is the target of several experiments, like the ``Cosmological Reionization Experiment'' (CORE) \\cite{CORE} and the ``Experiment for Detecting the Global EOR Signature'' (EDGES) \\cite{edges1, edges2}. Both experiments roughly operate in the frequency range from $\\sim 100$ to $\\sim 200$~MHz, corresponding to the  redshift range $7 \\lesssim z \\lesssim 14$. In fact, the single antenna experiment EDGES has already released preliminary results \\cite{edges1, edges2}. This experiment attempts to separate the redshifted HI 21\\,cm signal from the contribution of the Galactic and Extragalactic foregrounds at the same frequencies by taking advantage of the fact that these foregrounds are anticipated to be smooth power-law spectra. Conversely, the cosmological signal, is expected to have up to three rapid transitions in brightness temperature corresponding to the cooling and heating of the IGM, and most recently from reionisation. However, the rapidly varying (with respect to frequency) systematic and instrumental contributions to the measured power spectrum can easily be mistaken for a cosmological signal. Currently, the r.m.s. level of systematic contributions is approximately $T\\sim75$\\,mK, but further reductions in these systematics are anticipated in the near future \\cite{JuddPC}. Unfortunately, both CORE and EDGES operate in a frequency range where, as explained above, the evolution of the ionized fraction and of the kinetic temperature is dominated by astrophysics, so that presently they cannot be useful in constraining the models considered here. However, the EDGES team plans to expand the frequency coverage of the instrument down to 50\\,MHz or lower \\cite{edges1}, thus opening the possibility of exploring the regime where the 21\\,cm signal is dominated by the cosmological contribution. In order to exclude at least the most optimistic models considered here, the EDGES experiment should be able to reduce the systematics at 50\\,MHz to below $\\sim20$\\,mK, although the contribution of the Galactic synchrotron foreground increases significantly at the lower frequencies.  Another interesting way to potentially put observational constraints on the energy injection from DM annihilation would be to consider the effects on the spatial fluctuations of the brightness temperature, as encoded by their power spectrum, ${\\cal P}_{T_b}$, rather than the average signal as we have done in this paper. This was done, for example, in \\cite{furlanetto} (although the authors neglect the clumpiness of DM). Such fluctuations are somewhat easier to measure than the average signal since they are less contaminated by foregrounds.  Even more interestingly, peculiar velocities give rise to an anisotropy of the 21\\,cm power spectrum that can be used to separate the cosmological signal from the (uncertain) astrophysical contribution, thus allowing, at least in principle, one to detect the effects of DM annihilation in the astrophysics-dominated regime \\cite{Barkana:2004zy,McQuinn:2005hk,Pritchard:2008da}.  The 21\\,cm fluctuations are themselves the target of several experimens, such as LOFAR \\cite{lofar}, MWA \\cite{mwa}, PAPER \\cite{paper}, 21CMA \\cite{21CMA} and SKA \\cite{ska}. In particular, the impending LOFAR epoch of reionization experiment is designed to observe radio fluctuations at frequencies 115-215 MHz, corresponding to the redshifted 21\\,cm signal in the range $6<z<11.5$ \\citep{LOFAR}. Unfortunately, the cosmological signal is contaminated by a plethora of astrophysical and non-astrophysical components - including, Galactic synchrotron emission from diffuse and localized sources \\citep{shaver99}, Galactic free-free emission \\citep{shaver99}, integrated emission from extragalactic sources, such as radio galaxies and clusters \\citep{shaver99, dimatteo2002, dimatteo2004, oh2003, cooray2004}, ionospheric scintillation and instrumental response - the fluctuations in  which can significantly exceed the cosmological signal (see e.g.~\\citep{lofar_foregrounds, lofar_foregrounds2, Liu2009}).   We did not consider the possibility that the DM annihilation cross section is enhanced inside cold substructures due to non-perturbative quantum corrections \\cite{Lattanzi:2008qa}. This so-called Sommerfeld enhancement could increase the annihilation cross section by orders of magnitude with respect to its early-Universe value. The effect of such an enhancement on the heating/ionisation history of the IGM has been studied in \\cite{Cirelli:2009bb}, where it has been shown that it could similarly increase the IGM temperature by orders of magnitude.  An interesting extension of our analysis would be to consider other DM particles, such as ``Exciting DM'' (XDM) (see, e.g.,\\,\\citep{fink07}).  XDM can annihilate to produce two intermediate scalars $\\phi$ that can subsequently decay to standard model particles. If $2m_e<m_\\phi<2m_\\mu$, the $\\phi$ will mainly decay to an $e^+e^-$ pair, with an energy spectrum extending up to the mass of the XDM particle \\citep{cholis08}. Such particles are well motivated DM candidates and have been proposed to explain several other astronomical observations \\cite{fink07, chen09, cholis08}. Like LDM, the direct production of boosted $e^{+}e^{-}$ pairs following self-annihilation gives XDM the potential to produce observable features in the global 21\\,cm signal. In fact, recently it has been determined that particles with collisional long-lived excited states, and inspired by XDM models, may have observable effects on the CMB and the 21\\,cm background signal \\cite{fink08}.  We would also like to acknowledge the recent simulations by the Virgo Consortium as part of its Aquarius project \\cite{Aq1, Aq2}, conducted during the writing of the most recent version of this paper. These simulations investigate the properties of a Galaxy-sized DM halo and its substructure with unprecedented resolution. Whilst the results are largely consistent with those deduced from the Via Lactea simulations, there are significant differences. These include (i) four generations of resolved substructure (rather than the two in Via Lactea II), (ii) a distribution of substructures that is not self-similar to its host halo (rather than a fractal-like distribution observed in Via Lactea), (iii) a subhalo mass function with an index of -1.9 (rather than the -2.0 used in this study, although a brief discussion of the significance of such deviations on the clumping factor are made in \\S\\,\\ref{subsub}). Whilst we acknowledge that these differences may change some of our conclusions regarding the detectability of the 21\\,cm global signature, we do not pursue an investigation of these results here, but intend to incorporate them into a subsequent study \\cite{21cm_Cumberbatch_new}.  Finally, we note that the effect of DM annihilations on the 21\\,cm signal has been further studied in \\cite{Natarajan:2009bm}, including the effects on the brightness temperature fluctuations, that we have not considered here. On the other hand, the authors do not include the effect of substructures in their calculations.    {\\em Acknowledgements:}  DTC is supported by the Science and Technology Facilities Council. Part of this work was conducted while DTC was at CERCA Case Western Reserve University, funded by NASA grant 4200188792 and also supported in part by the US DoE, and also whilst at the University of Oxford, funded by the Science and Technology Facilities Council. Part of this work was conducted while ML was at the University of Oxford supported by the INFN. We would also like to thank Saleem Zaroubi, Rajat Thomas and the other members of the Kapteyn Institute, Groningen for their very helpful discussions regarding LOFAR sensitivities. We would also like to thank Judd Bowman, Aaron Chippendale and Ron Ekers for their helpful comments regarding the CoRE and EDGES experiments."
    },
    "0808/0808.2884_arXiv.txt": {
        "abstract": "In this work we address the problem of simultaneus multi-frequency detection of extragalactic point sources in maps of the Cosmic Microwave Background. We apply a new linear filtering technique, the so called `matched matrix filters', that incorporates full spatial information, including the cross-correlation among channels, without making any a priori assumption about the spectral behaviour of the sources. A substantial reduction of the background is achieved thanks to the optimal combination of filtered maps. We describe in detail the new technique and we apply it to the detection/estimation of radio sources in realistic all-sky \\emph{Planck} simulations at 30, 44, 70 and 100 GHz. Then we compare the results with the mono-frequencial approach based on the standard matched filter, in terms of reliablity, completeness and flux accuracy of the resulting point source catalogs. The new filters outperform the standard matched filters for all these indexes at 30, 44 and 70 GHz, whereas at 100 GHz both kind of filters have a similar performance. We find a noticeable increment of the number of true detections for a fixed reliability level. In particular, for a $95\\%$ reliability we practically double the number of detections at 30, 44 and 70 GHz. ",
        "introduction": "\\label{sec:intro}  From the search of extrasolar planets to the study of active galactic nuclei, one of the most common task in all the branches of Astronomy is the detection of faint pointlike objects. Such objects have angular sizes that are smaller than the angular resolution of the telescopes that are used to observe them, and therefore they are usually referred to as \\emph{point sources}.  A case of particular interest is the detection of extragalactic point sources (EPS) in maps of the Cosmic Microwave Background (CMB). EPS are known to be a relevant source of contamination for CMB studies, specially at small angular scales, where they hamper the estimation of CMB angular power spectrum both in temperature \\citep{tof98,zotti99,hob99,zotti05} and in polarization \\citep{tucci04,tucci05}. Therefore, for the sake of CMB analysis it is necessary to detect and remove as many extragalactic point souces as possible.  Moreover, in the frequency range spanned by CMB experiments the properties of EPS are poorly studied. Only very recently the \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP) satellite~\\citep{wmap0} has permitted the obtention of the first all-sky complete point source catalogs above $\\sim0.8$--$1$ Jy in the 23--94 GHz range of frequencies \\citep{wmap0,hinshaw07short,NEWPS07,chen08,wright08short}. The upcoming \\emph{Planck} mission~\\citep{planck_tauber05} will allow us to extend these catalogs down to lower flux limits and up to 857 GHz. There is interesting physics to be probed in this frequency range. The new EPS catalogs provided by next generation CMB experiments will not only allow us to follow the behaviour of source counts from existing catalogs to microwave frequencies, but also to study source variability and to discover rare objects such as inverted spectrum radio sources, extreme gigahertz peaked spectrum (GPS) sources and high-redshift dusty galaxies (see for example the \\emph{Planck Bluebook}~\\citep{bluebook} for a brief, yet comprehensive review of the rich phenomenology of EPS at microwave frequencies). Thus, the task of detecting point sources is important not only from the point of view of CMB science but also from the point of view of extragalactic Astronomy as well.  Let us consider a single image taken at a given wavelength. Then the problem consists on how to detect a number of objects, all of them with a common waveform that is generally considered to be well known (basically, the shape of point sources is that of the beam) but with unknown positions and intensities, that are embedded in additive noise (not necessarily white). In the field of microwave Astronomy, wavelet techniques~\\citep{vielva01,vielva03,MHW206,wsphere,NEWPS07}, matched filters~\\citep[MF,][]{tegmark98,barr03,can06} and other related linear filtering techniques~\\citep{sanz01,naselsky02,herr02b,herr02c,can04a,can05a,can05b} have proved to be useful. The common feature of all these techniques is that they rely on the prior knowledge that the sources have a distinctive spatial behaviour (i.e.~a known spatial profile, plus the fact that they appear as compact objects as opposed to `diffuse' random fields) that helps to distinguish them from the noise. Detection can be further improved by including prior information about the sources, i.e. some knowledge about their intensity distribution, in the frame of a Bayesian formalism~\\citep{hob03,psnakesI}.  Most of the current and planned CMB experiments are able to observe the sky at several wavelengths simultaneously. Multi-wavelength information makes it possible to separate different astrophysical components (as for example CMB from Galactic synchrotron emission) that have different spectral behaviour. Although multiwavelength component separation techniques have been very succesful in separating diffuse components~\\citep[for a recent comparative review of several methods applied to sky simulations very similar to the ones we will use in this paper, see][]{challenge08}, the detection of EPS has been usually attempted on a channel by channel basis. The reason for this is that EPS form a very heterogeneous population, constituted by a large number of objects with very different physical properties, and therefore it is impossible to define a common spectral behaviour for all of them. The whole situation remains somewhat unsatisfactory: on the one hand, the channel-by-channel approach based on the spatial behaviour works fine, but a valuable fraction of the information that multi-wavelenght experiments can offer is wasted this way. On the other hand, standard component separation techniques based on spectral diversity have problems when dealing with the very heterogeneous EPS components.  Thus, multi-wavelenght detection of EPS in CMB images remains a largely unexplored field. In recent years, some attempts have been done in this direction. For example, \\citet{naselsky02b} combined simulated multi-wavelenght maps ir order to increase the average signal to noise ratio of point sources. In a similar way,~\\citet{chen08} and~\\citet{wright08short} use combinations of the \\emph{WMAP} W and V bands in order to produce a CMB-free map in which to better detect the elusive radio galaxies. Note that, in any case, combined 'clean' maps are suitable for detecting more sources but not for performing accurate photometry, unless the spectral index of all the sources is known in advance.  An intermediate approach is to design filters that are able to find compact sources thanks to their distinctive spatial behaviour while at the same time do incorporate some multiwavelength information, without pretending to achieve a full component separation and without assuming a specific spectral behaviour for the sources. Very recently, the authors have proposed a new technique, based on the so-called \\emph{`matched matrix filters'} \\citep[MTXF,][]{herranz08a}, that goes in this direction. The basic underlying ideas of the new method are: \\begin{itemize} \\item When a source is found in one channel, it will be also present in the same position in all the other channels. \\item The spatial profile of the sources may differ from channel to channel, but it is a priori known. \\item The second order statistics of the background in which the sources are embedded is well known or it can be directly estimated from the data by assuming that point sources are sparse. This knowledge about the second order statistics (namely, the background's power spectrum for each channel and the its correlations among the different channels) will be used to increase the signal to noise ratio of the sources. \\item We want to perform accurate photometry of the sources at every one of the frequencies covered by the experiment, independently of what is the spectral behaviour of any source in particular. \\end{itemize} In~\\citet{herranz08a} the authors presented the new methodology and demonstrated its potential utility with a few toy simulations. In this paper, we will study its applicability to real CMB experiments by applying it to realistic simulations of the whole sky as will be observed by the upcoming \\emph{Planck} mission. We will focus on the particular case of the detection of radio sources in the four lower frequency \\emph{Planck} channels (33--100 GHz), comparing the performance of the new filters with the performance of the well-stablished standard matched filters. In section~\\ref{sec:matrixf} we will summarize the foundations of the theoretical formulation of matched matrix filters. In section~\\ref{sec:toplanck} we will describe the \\emph{Planck} simulations that we use to test the method and we will outline the main features of the code we have developed for its implementation. The results of the exercise will be commented in section~\\ref{sec:results}. Finally, in section~\\ref{sec:conclusions} we will draw some conclusions. ",
        "conclusions": "\\label{sec:conclusions}  Although there is a relatively large number of works in the literature devoted to the detection of point sources in CMB images, there is practically no one that addresses the problem from the multi-frequency point of view. The reason is that each individual extragalactic point source has its own (a priori unknown) spectral behaviour and this makes it very hard to accomodate them in classic component separation schemes.  In this work we apply a novel linear filtering technique, the \\emph{`matched matrix filters'}, introduced by~\\citet{herranz08a}. Matched matrix filters incorporate full spatial information, including the cross-correlation among channels, without making any a priori assumption about the spectral behaviour of the sources. The basic underlying idea is that in all the considered channels the sources do appear in the same unknown positions (but with different, unknown intensities) and have known spatial profiles given by the experiment point spread function, while some components of the background (i.e. CMB and Galactic emission) are correlated among channels. Then a clever linear filtering/combination of images can lead to a substantial reduction of the background, and therefore to lower source detection thresholds.  The resulting expresion of the filters takes form of a matrix given in terms of the cross-power spectra and the profile of the sources in the different maps.  We describe in detail the formalism of the matched matrix filters, looking in detail into some particular cases of interest, and we apply them to the detection of radio sources in realistic all-sky \\emph{Planck} simulations at 30, 44, 70 and 100 GHz. We use state-of-the-art simulations developed by the \\emph{Planck} WG2 Collaboration that include all the astrophysical components and instrumental specifications of the \\emph{Planck} mission. In order to tackle the strong non stationarity of the sky emission, we divide the sky into 371 $14.656\\times 14.656$ overlapping square degrees flat patches where we apply the two filtering techniques. For each patch and frequency we obtain a sub-catalog of detections. Then all the sub-catalogs corresponding to the same frequency are combined into an all-sky catalog. In order to compare with a well stablished mono-frequencial approach, we repeate the same process using the standard matched filter.  We compare both methods in terms of reliablity, completeness, receiver operating characteristics and flux accuracy. We find that for the three lower frequency channels (30, 44 and 70 GHz) the new matched matrix filters clearly outperform the standard matched filters for all these quality indicators. The matched matrix filters decrease significantly the noise level, what translates into a lower detection threshold and a reduced number of false detections. The flux estimation is consequently improved, with a lower dispersion around the true input value and a lower flux at which the well-know selection Eddington bias occur. The improvement is particularly evident at 44 GHz. At 100 GHz, however, the performance of the two filters is very similar. We indicate some possible reasons for this behavior, based on general analytical considerations about the structure of the filters.  One of the most interesting ways to compare two catalogs obtained with different methods is to set a fixed level of reliability, cut the catalogs at the corresponding points and compare how many detections there are in each of them. We find a noticeable increment of the number of true detections for a fixed reliability level obtained with the matched matrix filters with respect to the standard matched filters. In particular, for a $95\\%$ reliability we practically double the number of detections at 30, 44 and 70 GHz: the ratio between the number of detections obtainted with the MTXF and the MF for these channels are 1.8, 2.2 and 1.7, respectively.  We would like to stress once more the importance of including multi-frequency information in this approach. The new matched multi-filters can be used to increase very significatively the number of extragalactic point source detections in upcoming CMB experiments such as \\emph{Planck} as well as in current experiments such as WMAP. A work on the application of this technique to the 5 year WMAP data is in preparation.  Although here we have tested the MTXF to the particular case of the detection of radio sources in the low frequency channels of \\emph{Planck}, the same technique can be easily applied to other frequency bands, or to other fields of image analysis where pointlike objects appear in different frames (images)."
    },
    "0808/0808.1834_arXiv.txt": {
        "abstract": "This paper analyzes astrophysical scenarios that may be detected at the upper end of the energy range of  the  Gamma Ray Large Area Space Telescope (GLAST), as a result of cosmic-ray (CR) diffusion in the interstellar medium (ISM). Hadronic processes are considered as the source of $\\gamma$-ray photons from localized molecular enhancements nearby accelerators. Two particular cases are presented: a) the possibility of detecting spectral energy distributions (SEDs) with maxima above 1 GeV, which may be constrained by detection or non-detection at very-high energies (VHE) with observations by ground-based Cerenkov telescopes, and b) the possibility of detecting V-shaped, inverted spectra, due to confusion of a nearby (to the line of sight) arrangement of accelerator/target scenarios  with different characteristic properties. We show that the finding of these signatures (in particular, a peak at the 1--100 GeV energy region) is indicative for an identification of the underlying mechanism producing the $\\gamma$-rays that is realized by nature: which accelerator (age and relative position to the target cloud) and under which diffusion properties CR propagate. ",
        "introduction": "1506. This is the number of cosmic photons with energies above 10 GeV that were detected during the 9 years lifetime of the Energetic Gamma Ray Experiment Telescope (EGRET, Thompson et al. 2005). {  Out of this number of photons, 187 photons were found within 1 degree of sources that are listed in the Third EGRET Catalog (Hartman et al. 1999) and can be plausible related to the more energetic extent of the cataloged EGRET sources. The majority of the remaining photons correspond to diffuse Galactic and extragalactic radiation, albeit this conclusion is based on the similarity of their spatial and energy distributions with the diffuse model. } No significant time clustering nor source shining only at such high energies was detected. The scarcity of these detections represents the last electromagnetic window that remains to be opened between already explored energy ranges. EGRET was hampered in performing detailed studies of the $\\gamma$-ray sky above 10~GeV, due to back-splash of secondary particles produced by high-energy $\\gamma$-rays causing a self-veto in the monolithic anti-coincidence detector used to reject charged particles, and due to a non-calibrated detector response. GLAST, soon to be launched, will not be strongly affected by these effects since the anti-coincidence shield was designed in a segmented fashion (Moiseev et al. 2007). The effective area of the GLAST will be roughly an order of magnitude larger than that of EGRET leading to an increased sensitivity for detecting celestial $\\gamma$-ray photons (see Fig. 1 of Funk et al. 2008). GLAST's default observing plan is a survey mode where the sky is uniformly  observed. The increase of sensitivity and the survey mode open the gate to the possible discovery of new phenomenology. This paper analyzes astrophysical scenarios that can be detected  at the upper end of the energy range of  GLAST, as a result of cosmic-ray (CR) diffusion in the ISM. ",
        "conclusions": "Compton peaks (of which the first example could have been found already, Aharonian et al. 2006) are not the only way to generate a maximum in a SED in the range of 1--100 GeV. A large variety of parameters representing physical conditions in the vicinity of a CR accelerator could produce a rather similar effect. Distinguishing between these cases would require multiwavelength information, search for counterparts, and modeling. If such a maximum is interpreted hadronically, as a result of diffusion of CR in the ISM and their subsequent interaction with a nearby target, the results herein presented constrain, given the energy at which the maximum of the SED is reached, the characteristics of the putative accelerator, helping to the identification process. Indeed, one of the most distinguishing aspects of this study is the realization that these signatures (in particular, a peak at the 1--100 GeV energy region) is indicative for an identification of the underlying mechanism producing the $\\gamma$-rays that is realized by nature: which accelerator (age and relative position to the target cloud) and under which diffusion properties CR propagate, as it is exemplified in Figure 3. In a survey mode such as the one GLAST will perform, it might also be possible to observe rather unexpected, tell-tailing  SEDs, like those V-shaped presented here, if observed with instruments having a limited PSF, predictably leaving many Galactic sources unresolved. Indeed, we finally remark about the V-shaped spectra presented here that we have focused in the situation where {\\it from the same place in the sky} we have a double peak structure. If  instead two nearby sources at different energies are indeed resolvable by GLAST, i.e., displaced already in the GLAST range, and both above the threshold for detection, we would have a clear case of morphology change and spatially dislocation with energy what would certainly make the study also interesting. As noted recently \\footnote{In a poster by Gabici, Aharonian and Casanova in the Gamma-2008 meeting, July 6-11, 2008,  Heidelberg, and to our knowledge yet unpublished.}, a double peak could also be seen -coming from the same place in the sky- if we consider two populations of CRs interacting with the same cloud, for instance, the usual $E^{-2.75}$ Galactic bath of cosmic rays (producing a photon-spectrum peaking in the GeV regime) and the escaped CRs from the nearby source (producing a photon spectrum peaking in the TeV range). This possibility would, however, be applicable only to the case of very young accelerators, e.g., 2000 years or so, with clouds significantly separated from it, and thus it is a less general scenario than what we have explored in this paper as the mechanism for V-spectra production."
    },
    "0808/0808.1719_arXiv.txt": {
        "abstract": "We report the discovery of a probable new globular cluster in the disk of the Milky Way. Visible in 2MASS and the GLIMPSE survey, it has an estimated foreground extinction of $A_V \\sim 24$ mag. The absolute magnitude of the cluster and the luminosity function of the red giant branch are most consistent with that of an old globular cluster with a mass of a few $\\times 10^{5}  M_{\\odot}$ at a distance of 4--8 kpc. ",
        "introduction": "Harris (2001) estimated that there were $\\sim 20$ unknown Galactic globular clusters hidden behind substantial foreground extinction in the disk or behind the bulge. Subsequent near-IR surveys of the disk have borne out the prediction of missing clusters, two of which have been discovered in 2MASS by Hurt \\etal~(2000). Another cluster, GLIMPSE-C01, was found by Kobulnicky \\etal~(2005) using the Spitzer/IRAC GLIMPSE survey (Benjamin \\etal~2003) of the Galactic plane. The importance of the Galactic globular cluster system in understanding the formation, evolution, and destruction of globular clusters motivates continuing efforts to finish the census of clusters.  In this paper we report the discovery of a probable globular cluster at Galactic coordinates $l=14.13$, $b=-0.64$ (J2000 coordinates: $18^{\\textrm{h}} 18^{\\textrm{m}} 30^{\\textrm{s}}$ --$16^{\\circ} 58\\arcmin 36\\arcsec$). This object was identified by Mercer \\etal~(2005) in a search for star clusters in GLIMPSE, although it was not suggested to be a globular cluster. As this cluster is \\#3 in their catalog, we refer to the object as Mercer 3. This is relatively consistent with the rather confused naming conventions for Galactic globular clusters. ",
        "conclusions": "A secure identification of Mercer 3 as an old globular cluster will require moderately deep near-IR photometry. The predicted main sequence turnoff is at $K \\sim 19$; if instead it is a massive intermediate-age cluster, the turnoff will be more than a magnitude fainter, and the shape of the subgiant branch will be significantly different. Due to the large and differential reddening, an accurate estimate of the cluster metallicity will probably require near-IR spectroscopy. This is easily accomplished as the brightest red giants have $K \\sim 11$.  Mercer 3 is the second probable globular cluster discovered using Spitzer and 2MASS; a further two globular clusters were found using 2MASS alone. At the opposite end of parameter space, Koposov \\etal~(2007) discovered two extraordinarily low-mass globular clusters in the Sloan Digital Sky Survey at heliocentric distances of $\\sim 40-50$ kpc. The continuing pace of these discoveries suggests that the Galactic cluster census is far from complete."
    },
    "0808/0808.1696_arXiv.txt": {
        "abstract": "{The Fabry--Perot scanning interferometer mounted on the 6-m telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences is used to study the distribution and kinematics of ionized gas in the peculiar galaxy Arp\\,212 (NGC\\,7625, III\\,Zw\\,102). Two kinematically distinct subsystems---the inner disk and outer emission filaments---are found within the optical radius of the galaxy. The first subsystem, at galactocentric distances  $r<3.5$~kpc, rotates in the plane of the stellar disk. The inner part of the ionized-gas disk ($r<1.5$--2~kpc) exactly coincides with the previously known disk consisting of molecular gas. The second subsystem of ionized gas is located at galactocentric distances 2--6~kpc. This subsystem rotates in a plane tilted by a significant angle to the stellar disk. The angle of orbital inclination in the outer disk increases with galactocentric distance and reaches $50\\degr$ at $r\\approx6$~kpc. The ionized fraction of the gaseous disk does not show up beyond this galactocentric distance, but we believe that the HI disk continues to warp and approaches the plane that is polar with respect to the inner disk of the galaxy. Hence Arp\\,212 can be classified as a galaxy with a polar ring (or a polar disk). The observed kinematics of the ionized and neutral gas can be explained assuming that the distribution of gravitational potential in the galaxy is not spherically symmetric. Most probably, the polar ring have formed via accretion of gas from the dwarf satellite galaxy UGC~12549.  } ",
        "introduction": "Currently there is almost a consensus of opinion that gravitational interactions and mergers are among the most important factors in the evolution of galaxies, because such events change the observed structure, kinematics, star-formation history, etc. Polar-ring galaxies (PRG) are one of the consequences of those interactions. Such objects exhibit rings or disks consisting of gas, dust, and stars, which rotate in the plane that is approximately perpendicular (polar) to the disk of the main (``host'') galaxy. PRG are believed in most cases to owe their formation to galaxy mergers and accretion of the matter of the companion galaxy or gaseous filaments from the intergalactic medium onto the host galaxy (Bournaud \\& Combes \\cite{bc03:Moiseev_n}; Combes (\\cite{comb06:Moiseev_n}), and references therein).  PRG exhibit circular rotation in two mutually perpendicular planes thereby making it possible to study three-dimensional mass distribution in the galaxy and determine the shape of its dark halo: oblateness, prolateness, and deviation from axial symmetry (Combes, \\cite{comb06:Moiseev_n}). To do this, it is necessary to obtain sufficiently detailed data about the inner kinematics of PRG. Whitmore et al. (\\cite{w90:Moiseev_n}) list 157 candidate polar-ring galaxies selected mostly by their peculiar appearance. However, the number of ``real'' PRG, i.e., which exhibit rotation in orthogonal planes, is much smaller. Even in the simplest case, where both the central galaxy and the ring are seen edge on, at least two longslit spectroscopic sections are needed to determine the rotation pattern of both subsystems. A number of problems arise in the cases where the spatial orientations of the galaxy and ring are such that one of the rotation planes is seen at a moderate angle to the line-of-sight. First, in many cases it is impossible to establish the PRG nature of an object solely from its direct images. Second, a substantially greater number of spectroscopic sections with different slit orientations  are needed for a detailed kinematical study. In these cases it is better to determine the two-dimensional velocity field using optical panoramic (3D) spectroscopy or radio interferometry in lines of molecular and atomic gas.  The NGC\\,2655 galaxy is a good illustration. It was suspected to have a polar ring due to the powerful dust lane crossing the disk of the galaxy. A comparison of the velocity fields of gas and stars in the circumnuclear  region confirmed such an interpretation (Sil'chenko \\& Afanasiev, \\cite{silafan04:Moiseev_n}), and recent 21-cm line observations allowed the structure and kinematics of the polar ring to be analyzed (Sparke et al., \\cite{sparke08:Moiseev_n}). Note that for such complex objects even an extensive set of both photometric and kinematical data may be insufficient for a definitive conclusion about the structure of subsystems in the galaxy considered, see, e.g., the case of UGC\\,5600 (Shalyapina et al., \\cite{labuda07:Moiseev_n}).  In this paper we analyze new observational data for the peculiar galaxy Arp\\,212. Although this object was suspected to be a possible candidate PRG, no polar component has been found yet. In Section~\\ref{intro:Moiseev_n} we summarize the data related to this galaxy. In Section~\\ref{obs:Moiseev_n} we describe observations and reduction of the data obtained with the 6-m telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences. In Section~\\ref{tilt:Moiseev_n} we show that the velocity field of ionized gas contains two subsystems with different rotation velocities. In the next Section~\\ref{2D:Moiseev_n} we construct a two-dimensional model of the rotation of gas in the outer regions to explain the kinematics of the polar ring. In Section~\\ref{discuss:Moiseev_n} we discuss the overall structure of the galaxy and finally in Section~\\ref{results:Moiseev_n} formulate the main results and conclusions of the paper. ",
        "conclusions": "\\label{results:Moiseev_n}  We constructed and analyzed the velocity field of ionized gas in the peculiar galaxy Arp\\,212 and found it to have two kinematically distinct subsystems of rotation gas---the inner disk and outer HII regions. The inner disk is located within 3.5 kpc of the center and ionized gas in this disk rotates in the plane coincident with the plane of the stellar disk. Note that the inner part of the ionized-gas disk is exactly coincident with the earlier known molecular-gas disk. The second subsystem of ionized gas is located at a galactocentric distance of \\mbox{$r$=2--6\\,kpc} and consists of isolated HII regions whose orbits are tilted by significant angles to the stellar disk. Our own and published data on the kinematics of molecular, ionized, and neutral gas can be explained in terms of the following model. Most of the HI mass in the outer parts of the galaxy concentrates in a broad ring with radius of about 20\\,kpc. Outer ring regions rotate in the plane that is orthogonal to the plane of the stellar disk. With decreasing galactocentric distance the orbits of gaseous clouds precess and approach the plane of the disk. This precession is due to the nonspherical   (maybe a triaxial) distribution of the gravitational potential in the galaxy. A burst of star formation takes place in the inner regions of the warped polar ring. The angle between the ring and the plane of the galaxy is equal to $50\\degr$ at $r=R_{25}$, it continues to decrease with decreasing  radius, and at $r\\approx$2--3~kpc the gas from the ring falls onto the plane of the galaxy.   The powerful dust ring in the central region of the galaxy is indicative of the fall of gas from the polar ring onto the plane of the stellar disk of  Arp\\,212. The polar ring may have formed as a result of the interaction with the gas-rich dwarf galaxy UGC~12549, however, new observations of the distribution and kinematics of neutral hydrogen are required to either confirm or disprove this scenario."
    },
    "0808/0808.0145_arXiv.txt": {
        "abstract": "Mode conversion in the region where the sound and Alfv\\'en speeds are equal is a complex process, which has been studied both analytically and numerically, and has been seen in observations.  In order to further the understanding of this process we set up a simple, one-dimensional model, and examine wave propagation through this system using a combination of analytical and numerical techniques.  Simulations are carried out in a gravitationally stratified atmosphere with a uniform, vertical magnetic field for both isothermal and non-isothermal cases.  For the non-isothermal case a temperature profile is chosen to mimic the steep temperature gradient encountered at the transition region.  In all simulations, a slow wave is driven on the upper boundary, thus propagating down from low-$\\beta$ to high-$\\beta$ plasma across the mode-conversion region.  In addition, a detailed analytical study is carried out where we predict the amplitude and phase of the transmitted and converted components of the incident wave as it passes through the mode-conversion region.  A comparison of these analytical predictions with the numerical results shows good agreement, giving us confidence in both techniques.  This knowledge may be used to help determine wave types observed and give insight into which modes may be involved in coronal heating. ",
        "introduction": "\\label{intro.sec} Mode conversion is the process under which an incident wave may be wholly or partially converted into another wave mode, with the remainder of the original wave being transmitted through the plasma.  This may occur where the sound and Alfv\\'en speeds are equal, or equivalently, where the plasma $\\beta$ (the ratio of the gas to the magnetic pressure) is approximately unity.  This layer will generally lie low down in the solar atmosphere, in the region of the lower chromosphere.  This problem has been investigated extensively in the past, many using analytical techniques.  \\cite{Zhugzhda1979}, \\cite{Zhugzhda1981} and \\cite{Zhugzhda1982} investigated the conversion of fast magnetoacoustic waves into slow magnetoacoustic waves as they propagate up through the solar atmosphere from high- to low-$\\beta$ plasma.  Their model looked at a uniform vertical magnetic field set in an isothermal atmosphere stratified by gravity.  Using the fact that, for harmonic waves, an exact solution may be found in terms of the known Meijer-G functions, coefficients describing the conversion and transmission were found.  \\cite{Cally2001} expanded on this work by noting that these solutions may also be written in terms of the equivalent Hypergeometric $_2F_3$ functions, and another set of coefficients was discovered.  \\begin{figure}[t] \\centering \\includegraphics[scale=0.5]{model.eps} \\caption{Cartoon of our model atmosphere, with uniform vertical magnetic field.  The $z$-axis points upwards (opposite to gravity) and a slow wave is driven on the upper boundary, travelling down towards the mode-conversion layer at $\\beta\\approx1$.} \\label{model.fig} \\end{figure}  We also investigate the mode conversion problem analytically, although we combine this with numerical simulations.  We look at a uniform background magnetic field, but, in contrast to the aforementioned work, we investigate downward propagation - so the wave is travelling from low- to high-$\\beta$ plasma (as shown in Figure~\\ref{model.fig}).  As the fast wave is evanescent in low-$\\beta$ plasma, we look at the conversion of slow to fast magnetoacoustic waves.  In Section~\\ref{eqns.sec} the basic equations are set out, and we look at the characteristic wave behaviour in high- and low-$\\beta$ plasma.  Two different situations are then examined: an isothermal atmosphere in Section~\\ref{iso.sec}, and a non-isothermal atmosphere in Section~\\ref{noniso.sec}.  Finally the conclusions are detailed in Section~\\ref{conc.sec}. ",
        "conclusions": "\\label{conc.sec}  Using a combination of analytical techniques and numerical computations we have studied the phenomenon of mode conversion in a simple, one-dimensional model.  More specifically, we have studied the downward propagation of linear waves in a gravitationally stratified atmosphere permeated by a uniform, vertical magnetic field.  These waves propagate from low- to high-$\\beta$ plasma, and the mode conversion occurs where the sound and Alfv\\'en speeds are equal where $\\beta\\approx1$.  In Section~\\ref{iso.sec} we concentrated on an isothermal atmosphere using a combination of the WKB method, and a method developed by \\cite[Cairns \\& Lashmore-Davies]{cairns1983} to examine mode conversion both outside and around the conversion region.  This allowed conversion and transmission coefficients to be found.  It was then possible to compare the transmission predicted by the analytical method with the numerical data.  These showed excellent agreement, even as $k_x$ becomes large, violating the assumptions.  This analysis was repeated in Section~\\ref{noniso.sec} for a non-isothermal atmosphere.  We selected a $\\tanh$ temperature profile to mimic the steep temperature gradient of the transition region.  Surprisingly the transmission and conversion coefficients found were identical to those for the isothermal atmosphere.  So the temperature profile has no effect on the amount of transmission and conversion as a wave undergoes mode conversion.  Although the model set-up is not physically realistic, it has given us a great insight into the mode conversion process, and the techniques which may be utilised in its analysis.  We plan to use this knowledge to look at a more realistic model of a two-dimensional coronal null point, extending on work done by \\cite{McLaughlin2006}.  In this situation the plasma $\\beta$ is infinite at the null point, so as the wave propagates towards the null it will indeed be passing from low- to high-$\\beta$ plasma, as we have examined here."
    },
    "0808/0808.1335.txt": {
        "abstract": "We analyze seven different viable $f(R)$-gravities towards the Solar System tests and stochastic gravitational waves background. The aim is to achieve experimental bounds for the theory at local and cosmological scales in order to select models capable of addressing the accelerating cosmological expansion without cosmological constant but evading the weak field constraints. Beside large scale structure and galactic dynamics, these bounds can be considered complimentary in order to select self-consistent theories of gravity working at the infrared limit. It is demonstrated that seven viable $f(R)$-gravities under consideration not only satisfy the local tests, but additionally, pass the above PPN-and stochastic gravitational waves bounds for large classes of parameters. ",
        "introduction": "The currently observed accelerated expansion of the Universe suggests that cosmic flow dynamics is dominated by some unknown form of dark energy characterized by a large negative pressure. This picture comes out when such a new ingredient, beside baryonic and dark matter,  is considered as a source in the r.h.s. of the field equations. Essentially, it should be some form of un-clustered, non-zero vacuum energy which, together with (clustered) dark matter, should drive the global cosmic  dynamics.  Among the proposals to explain the experimental situation,  the ``concordance model'', addressed as $\\Lambda$CDM,  gives a reliable snapshot of the today observed Universe according to the CMBR, LSS and SNeIa data,  but presents dramatic shortcomings as the ``{\\it coincidence and  cosmological constant problems}'' which point out its inadequacy  to fully trace back the cosmological dynamics \\cite{Peebles}.  On the other hand, alternative theories of gravity,  extending in some way General Relativity (GR), allows to pursue a different approach  giving rise to suitable cosmological models where a late-time accelerated expansion can be achieved in several  ways. This viewpoint does not require to find out candidates for dark energy and dark matter at fundamental level (they have not been detected up to now), it takes into account only the ``observed'' ingredients (i.e. gravity, radiation and baryonic matter), but the l.h.s. of the Einstein equations has to be modified.  Despite of this modification, it could be in agreement with the spirit of GR since the only request is that the Hilbert-Einstein action should be generalized asking for a gravitational interaction  acting, in principle,  in different ways at different scales \\cite{CCT}.  The idea that Einstein gravity should be extended or corrected at large scales (infrared limit) or at high energies (ultraviolet limit) is suggested by several theoretical and observational issues. Quantum field theory in curved spacetimes, as well as the low-energy limit of String/M theory, both imply semi-classical effective actions containing  higher-order curvature invariants or scalar-tensor terms. In addition, GR has been definitely tested only at Solar System scales while it may show several shortcomings if checked at higher energies or larger scales. Besides,  the Solar System experiments are, up to now, not so conclusive to state that the only viable theory of gravity  is GR: for example, the limits on PPN parameters should be greatly improved to fully remove degeneracies \\cite{gaia}.  Of course, modifying the gravitational action asks for several fundamental challenges. These models can exhibit instabilities \\cite{instabilities-f(R)} or ghost\\,-\\,like behavior \\cite{ghost-f(R)}, while, on the other hand, they have to be matched with observations and experiments in the appropriate low energy limit.  Despite of all these issues, in the last years, some interesting results have been achieved in the framework of the so called $f(R)$-gravity at cosmological, Galactic and Solar System scales. Here $f(R)$ is a general (analytic)  function of the Ricci scalar $R$ (see Refs.~\\cite{odirev,GRGrev,faraoni} for review).   For example, there exist cosmological solutions that give the accelerated expansion of the universe at late times \\cite{f(R)-noi,f(R)-cosmo,mimic,prado}.  In addition, it has been discovered that some stability conditions can lead to avoid ghost and tachyon solutions. Furthermore there exist viable $f(R)$ models which satisfy both background cosmological constraints and stability conditions \\cite{Li,Hu,Star,Appleby,Tsuji,Nojiri:2007cq,sebast,Nojiri} and results have been achieved in order  to place constraints on $f(R)$ cosmological models by CMBR anisotropies and galaxy power spectrum \\cite{bean-f(R),Faul,huspe}. Moreover, some of such viable models lead to the unification of early-time inflation with late-time acceleration \\cite{Nojiri:2007cq,sebast,Nojiri}.  On the other hand, by considering $f(R)$-gravity in the low energy limit, it is possible to obtain corrected gravitational potentials capable of explaining  the flat rotation curves of spiral galaxies or the  dynamics of galaxy clusters without considering huge amounts of dark matter \\cite{noi-mnras,salucci,sobouti,mendoza,Boe,lobo}.  Furthermore, several authors have dealt with the weak field limit of fourth order gravity, in particular considering the PPN limit \\cite{bertolami,ppn-tot1,ppn-tot,Navarro,Erick,lgcpapers,mpla,CST} and the spherically symmetric solutions \\cite{spher-f(R),multamaki06,kainulainen,noether}.  This great deal of work needs an essential issue to be pursued: we need to compare experiments and probes at local scales (e.g. Solar System) with experiments and probes at large scales (Galaxy, extragalactic scales, cosmology) in order to achieve self-consistent $f(R)$ models. Some work has been done in this direction (see e.g. \\cite{Hu}) but the large part of efforts has been devoted to address single data sets (observations at a given redshift) by a single model which, several time, is not working at other scales than the one considered. In particular, a given $f(R)$ model, evading Solar System tests, should be not simply extrapolated at extragalactic and cosmological scales only requiring accelerated cosmological solutions but it should be confronted with data and probes coming from cosmological observations. Reliable models are then those matching data at very different scales (and redshifts).   In order to constrain further viable $f(R)$-models, one could take into account also  the stochastic background of gravitational waves (GW) which, together with cosmic microwave background radiation (CMBR), would  carry a huge amount of information on the early stages of the Universe evolution.  In fact, if detected, such a background could constitute a further probe for these theories at very high red-shift \\cite{tuning}. On the other hand, a key role for the production and the detection of the relic gravitational radiation background is played by the adopted theory of gravity \\cite{Maggiore,BBFN}. This means that the effective theory of gravity should be probed at zero, intermediate and high redshifts to be consistent at all scales and not simply extrapolated up to the last scattering surface, as in the case of GR.  The aim of this paper is to discuss the  PPN Solar-System constraints and the GW stochastic background considering some recently proposed $f(R)$ gravity models \\cite{Star,Li,Hu,Nojiri:2007cq,sebast,Nojiri} which satisfy both cosmological and stability conditions mentioned above. Using the definition of PPN-parameters $\\gamma$ and $\\beta$ in terms of $f(R)$-models \\cite{mpla} and the definition of scalar GWs \\cite{CCD}, we compare and discuss if it is possible to search for parameter ranges of  $f(R)$-models  working at Solar System and GW stochastic background scale. This phenomenological approach is complementary to the one proposed, e.g. in \\cite{Hu,huspe} where also galactic and cosmological scales have been considered to constraint the models.  The layout of the paper is the following.  In  Sec. II, we review the field equations of $f(R)$ gravity in the metric approach and their scalar-tensor representation, useful to compare the theory with observations. In Sec.III, we review and discuss some viable $f(R)$ models capable of satisfying both local gravity prescriptions  as well as the observed cosmological behavior. In particular, we discuss their stability conditions and the field values which have to achieved to fulfill physical bounds. Sec. IV is devoted to derive the values of model parameters in agreement with the PPN experimental constraints while, in Sec. V, we deal with the constraints coming from the stochastic background of GWs. These latter ones have to be confronted with those coming from PPN parameterization. Discussion and conclusions are drawn in Sec. VI. As a general remark, we find out that bounds coming from the interferometric ground-based (VIRGO, LIGO) and space (LISA) experiments could constitute a further probe for $f(R)$ gravity if matched with bounds at other scales. ",
        "conclusions": "In this paper, we have investigated the possibility that some viable $f(R)$ models could be constrained considering both Solar System experiments and upper bounds on the stochastic background of gravitational radiation. Such bounds   come from interferometric ground-based (VIRGO and LIGO) and space (LISA) experiments. The underlying philosophy is to show that the $f(R)$ approach, in order to describe consistently  the observed universe, should  be tested at very different scales, that is at very different redshifts. In other words, such a proposal could partially contribute to remove the unpleasant degeneracy affecting the wide class of dark energy models, today on the ground.  Beside the request to evade the Solar System tests, new methods have been recently proposed to investigate the evolution and the power spectrum of cosmological perturbations in $f(R)$ models \\cite{huspe}. The investigation of stochastic background, in particular of the scalar component of GWs coming from the $f(R)$ additional degrees of freedom, could acquire, if revealed by the running and forthcoming experiments,  a fundamental importance to discriminate among the various  gravity theories \\cite{tuning}. These data (today only upper bounds coming from simulations) if combined with Solar System tests, CMBR anisotropies, LSS, etc. could greatly help to achieve a self-consistent cosmology bypassing the shortcomings of $\\Lambda$CDM model.  Specifically, we have taken into account some broken power law $f(R)$ models fulfilling the main cosmological requirements which are to match the today observed accelerated expansion and the correct behavior in early epochs. In principle, the adopted parameterization allows to fit data at extragalactic and cosmological scales \\cite{Hu}. Furthermore, such models are constructed to evade the Solar System experimental tests. Beside these broken power laws, we have considered also two models capable of reproducing the effective cosmological constant, the early inflation and the late acceleration epochs \\cite{sebast}. These $f(R)$-functions  are combinations of hyperbolic tangents.  We have discussed the behavior of all the considered models. In particular,  the problem of  stability has been addressed determining suitable and physically consistent ranges of parameters. Then we have taken into account the results of the main  Solar System  current experiments. Such results give upper limits on the PPN parameters which any self-consistent theory of gravity should satisfy at local scales. Starting from these, we have selected the $f(R)$ parameters fulfilling the tests. As a general remark, all the functional forms chosen for $f(R)$ present sets of parameters capable of matching the two main PPN quantities, that is $\\gamma_{exp}$ and $\\beta_{exp}$. This means that, in principle, extensions of GR are not {\\it a priori} excluded as reasonable candidates for gravity theories. To construct such extensions, the reconstruction method developed in \\cite{odintsov} may be applied.  The interesting feature, and the main result of this paper, is that such sets of parameters are not in conflict with bounds coming from the cosmological stochastic background of GWs. In particular, some sets of parameters  reproduce quite well both the PPN upper limits and the constraints on the scalar component amplitude of GWs.  Far to be definitive, these preliminary results indicate that self-consistent models could be achieved comparing experimental data at very different scales without extrapolating results obtained only at a given scale.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0808/0808.0003_arXiv.txt": {
        "abstract": "Observations of redshift-space distortions in spectroscopic galaxy surveys offer an attractive method for observing the build-up of cosmological structure. In this paper we develop and test a new statistic based on anisotropies in the measured galaxy power spectrum, which is independent of galaxy bias and matches the matter power spectrum shape on large scales. The amplitude provides a constraint on the derivative of the linear growth rate through $f\\sigma_8({\\rm mass})$. This demonstrates that spectroscopic galaxy surveys offer many of the same advantages as weak lensing surveys, in that they both use galaxies as test particles to probe all matter in the Universe. They are complementary as redshift-space distortions probe non-relativistic velocities and therefore the temporal metric perturbations, while weak lensing tests the sum of the temporal and spatial metric perturbations. The degree to which our estimator can be pushed into the non-linear regime is considered and we show that a simple Gaussian damping model, similar to that previously used to model the behaviour of the power spectrum on very small scales, can also model the quasi-linear behaviour of our estimator. This enhances the information that can be extracted from surveys for $\\Lambda$CDM models. ",
        "introduction": "It is now widely accepted that the large-scale structure we see in the distribution of galaxies arises through a process of gravitational instability, which amplifies primordial fluctuations laid down in the very early Universe. The rate at which structure grows from these small perturbations offers a key discriminant between cosmological models. For instance, dark energy models in which general relativity is unmodified predict different Large-Scale Structure formation compared with Modified Gravity models with the same background expansion \\citep[e.g.][]{dvali00,carroll04,Bra05,NesPer08}.  Structure growth is driven by the motion of matter (and inhibited by the cosmological expansion). Galaxies are expected to act as test particles within this matter flow, so the motion of galaxies carries an imprint of the rate of growth of large-scale structure. Because of this, many previous analyses have shown that observations of these galaxy peculiar velocities can distinguish between classes of models \\citep[e.g.][]{jain07,song08,song08b}. A key technique to statistically measure the growth of the velocity field uses redshift-space distortions seen in galaxy surveys \\citep{Kai87}. Galaxy maps, produced by estimating distances from redshifts obtained in spectroscopic galaxy surveys, reveal an anisotropic galaxy distribution. The anisotropies arise because galaxy recession velocities, from which distances are inferred, include components from both the Hubble flow and peculiar velocities from the comoving motions of galaxies. These distortions encode information about the build-up of structure.  Many previous surveys have been analysed to measure $\\beta\\approx\\Omega_m^{0.6}/b$, where $b$ is the deterministic, local, linear bias of the galaxies.  The latest generation of large surveys have provided ever tighter constraints. Analyses using the 2-degree Field Galaxy Redshift Survey \\citep[2dFGRS;][]{colless03} have measured redshift-space distortions in both the correlation function \\citep{peacock01,hawkins03} and power spectrum \\citep{percival04}. Using the Sloan Digital Sky Survey \\citep[SDSS;][]{york00}, redshift-space distortions have also been measured in the correlation function \\citep{Zeh05,Oku08,cabre08}, and using an Eigenmode decomposition to separate real and redshift-space effects \\citep{tegmark04,tegmark06}. These studies were recently extended to $z\\simeq1$ \\citep{guzzo08} using the VIMOS-VLT Deep Survey \\citep[VVDS;][]{lefevre05,garilli08} and the 2SLAQ survey \\citep{Ang08}. In addition to measuring $\\beta$ at $z=0.8$, \\citet{lefevre05} emphasised the importance of using large-scale peculiar velocities for constraining models of cosmic acceleration.  On linear scales the theory behind the observed redshift-space distortions is well developed \\citep{Kai87,HamiltonReview}, and previous analyses have made use of this theory to measure $\\beta$. On quasi-linear and non-linear scales we are instead reduced to making approximations, or using fitting formulae based on numerical simulations. The standard approach is to use a `streaming' model, where linear theory is spliced together with an approximation for random motion of particles in collapsed objects (see section~\\ref{sec:zspace_theory}). \\citet{TinWeiZhe06} and \\citet{Tin07}, which are developments of work in \\citet{HatCol99}, discuss a number of possible improvements to the streaming model based on fits to numerical simulation results. These approaches aim to allow us to extract information on $\\beta$ from scales where the power spectrum does not match its linear form. One concern is that the theoretical dependence on $\\beta$ might change in the quasi-linear regime, leading to complicated dependencies and that fits to simulations will only be correct in a subset of cosmologies and galaxy formation models similar to those from which the fits were derived\\footnote{The fits are quite sensitive to the satellite fraction, for example.}.  A better approach would be to analyse the physics behind redshift-space distortions, and to try to use this to devise the best estimator with which to extract cosmological information. Because galaxy velocities only depend on the distribution of matter, we can devise an estimator $\\hat{P}(k)$ that is not affected by galaxy bias: instead it measures the large-scale shape of the matter power spectrum. Because the estimator is based on the velocity power spectrum, the large-scale amplitude is $f^2$ times that of matter fluctuations, where $f$ is the logarithmic derivative of the linear growth rate with respect to the scale factor, $f\\equiv d\\ln D/d\\ln a$. From this estimator, we can measure $f\\sigma_{8,\\,{\\rm mass}}$, which is proportional to $dD/d\\ln a$, and provides a good discriminant between modified gravity and dark energy models \\citep{song08b}. $\\hat{P}(k)$ is expected to match the linear matter power spectrum shape on large scales, so small-scale differences will help us to determine the scales on which the theory is breaking down.  For $\\Lambda$CDM models, we show that a simple Gaussian smoothing of the redshift-space power along the line-of-sight is able to match most of this break-down. This places redshift-space distortion measurements on the same footing as weak lensing measurements in the sense that they both allow us to test the matter distribution directly. They provide complementary information, as non-relativistic velocity measurements only depend on temporal metric perturbations, while weak lensing tests the sum of the temporal and spatial metric perturbations.  The layout of our paper is as follows. We first review the theory behind redshift-space distortions (Section~\\ref{sec:zspace_theory}) in the linear, quasi-linear and non-linear regimes. In Section~\\ref{sec:estimators}, we introduce a new estimator, based on the monopole and quadrupole from a Legendre decomposition of the redshift-space power spectrum, which is designed to recover a power spectrum given by $f^2$ times the matter power spectrum on large scales. Section~\\ref{sec:sims} introduces the N-body simulation used to test this theory and our new estimator. We use spherically averaged power spectra where we include distortions along multiple axes to analyse this simulation, an approach described in Section~\\ref{sec:pk_sph_av}. Our analytic theory is compared with results from the numerical simulation in Section~\\ref{sec:sim_results}. The paper ends with a discussion of our results. ",
        "conclusions": "Redshift-space distortions encode key information about the build-up of cosmological structure. In the linear regime, the theory behind the density enhancement was developed over 20 years ago \\citep{Kai87}. Following this development, most observational studies have focused on measuring $\\beta$. One method to do this is to measure the quadrupole-to-monopole ratio, which depends on $\\beta$ in a known way (Eq.~\\ref{eq:P2overP0}). By reviewing the theory of redshift-space distortions in Section~\\ref{sec:zspace_theory}, we argued that we should be able to obtain a bias independent cosmological constraint on $f\\sigma_8({\\rm mass})$ from the large-scale redshift-space distortions. This statistic, proportional to $dD/d\\ln a$ with a proportionality constant depending on the amplitude of fluctuations at early times, provides an excellent test of DE models \\citep[e.g.][]{song08b}.  In Section~\\ref{sec:qlin} \\&~\\ref{sec:nlin}, we set out to extend the theory into the quasi- and non-linear regimes, to investigate the smallest scales on which we can easily recover cosmological information using this physical process. In the linear regime, the power spectrum can be decomposed into terms dependent on $\\mu^0$, $\\mu^2$ and $\\mu^4$. As we push to quasi-linear scales we should expect this behaviour to become more complicated with extra $\\mu^6$ and higher order terms becoming important. We should also expect the simple relations $bP_{g\\theta}=fP_{gg}$ and $b^2P_{\\theta\\theta}=f^2P_{gg}$ to break down. Previous work has introduced a ``streaming'' model for non-linear redshift-space distortions, including a damping term spliced together with linear theory. In the limit $k\\to0$, the standard Gaussian or Exponential forms given in Eqns.~(\\ref{eq:F_exp}) \\&~(\\ref{eq:F_gauss}) for this damping function both reduce to providing an extra $\\mu^6$ term with amplitude dependent on $\\sigma$. We might therefore expect these terms can also be used to match the expected quasi-linear $\\mu$ dependence of the power spectrum. In this case, the same function might not match all quasi-linear and non-linear scales.  Using a large numerical simulation we have set out to test this theory. To facilitate this, we have introduced a novel approach to the analysis of the redshift-space power spectrum measured from simulations by allowing redshift-space distortions to be included along multiple axes. Power spectra calculated from density fields including redshift-space distortions along multiple axes act as moments of $P_g^s(k,\\mu)$, allowing us to decompose into the expected $\\mu^0$, $\\mu^2$ and $\\mu^4$ components. On large scales we see no evidence for strong velocity bias, and the velocity-velocity power spectrum has an amplitude $f^2$ times that of the mass to within 10\\% (right panel of Fig.~\\ref{fig:pk_split}). The density-velocity power spectrum, from which most information is obtained from surveys is less affected by this velocity bias. Fig.~\\ref{fig:pk} shows that the relations $bP_{g\\theta}=fP_{gg}$ and $b^2P_{\\theta\\theta}=f^2P_{gg}$ break down on scales $k>0.1\\hompc$. The right panel of Fig.~\\ref{fig:pk_axes} shows that $\\mu^6$ and higher order terms need to be included for $k>0.2\\hompc$. The Exponential and Gaussian models enable a better fit for $0.2<k<0.4\\hompc$, but neither is perfect on all of these scales, although the Gaussian model does slightly better than the Exponential model in the cases we illustrate. Alternative models which reduce to $\\mu^6$ behaviour in the limit $k\\to0$ might perform better, but this is not necessary for our purposes: we have shown that the Gaussian model extends the range of the fit.  Based on this analysis, we have devised a new estimator of the matter power spectrum $\\hat{P}$ using the monopole and quadrupole power spectra. This matches the large-scale shape of the matter power spectrum with amplitude multiplied by $f^2$, so we can measure $f\\sigma_8({\\rm mass})$. Including the Gaussian damping model allows us to extend the match with the simulation data for $k\\simlt0.2\\hompc$. The match between the shape of the expected linear matter power spectrum and $\\hat{P}$ should allow a test of the scales over which we can measure $f\\sigma_8({\\rm mass})$.  Rather than use such an estimator, it would also be possible to perform a maximum likelihood fit of Eq.~(\\ref{eq:pgs_comb_lin}) to the measured redshift-space power spectrum, with $f\\sigma_8({\\rm mass})$ and $b\\sigma_8({\\rm mass})$ as the free parameters.  Our analysis has assumed the distant observer limit, a standard approximation used in many analyses. It would be straightforward to extend this work to methods that incorporate a proper radial--angular split allowing for standard survey geometries. An additional extension would be to perform such an analysis in configuration space: again, it would also be possible the extend the ideas behind our estimator to analyse the correlation function."
    },
    "0808/0808.0151_arXiv.txt": {
        "abstract": "{} {The AGILE gamma-ray burst GRB 080514B is the first burst with detected emission above 30 MeV and an optical afterglow. However, no spectroscopic redshift for this burst is known.} {We compiled ground-based photometric optical/NIR and millimeter data from several observatories, including the multi-channel imager GROND, as well as ultraviolet \\swift UVOT and X-ray XRT observations. The spectral energy distribution of the optical/NIR afterglow shows a sharp drop in the \\swift UVOT UV filters that can be utilized for the estimation of a redshift.} {Fitting the SED from the \\swift UVOT $uvw2$ band to the $H$ band, we estimate a photometric redshift of $z=1.8^{+0.4}_{-0.3}$, consistent with the pseudo redshift reported by Pelangeon \\& Atteia (2008) based on the gamma-ray data.} {The afterglow properties of GRB 080514B do not differ from those exhibited by the global sample of long bursts, supporting the view that afterglow properties are basically independent of prompt emission properties.} ",
        "introduction": "Gamma-Ray Bursts (GRBs) are the most luminous explosions in the Universe, with the bulk of the released energy emerging in the 0.1 to 1 MeV range (e.g., Kaneko et al. 2006; Preece et al. 2000). Indeed, most bursts have not been observed at energies much above 1 MeV, where low photon counts and typically small instrumental collecting areas hamper the gathering of data. For example, the \\it Burst And Transient Source Experiment \\rm (BATSE; operating from 25 keV to 2 MeV) aboard the \\it Compton Gamma-Ray Observatory \\rm (CGRO) detected 2704 bursts during a nine-year lifetime (1991 to 2000), while the COMPTEL telescope on CGRO, operating in the  0.8 MeV to 30 MeV range, in the same time period observed only 44 events with high significance (Hoover et al. 2005). The number of GRBs detected at even higher energies with the EGRET TASC and spark chamber instruments is even lower (Dingus 1995; Kaneko et al. 2008).  \\begin{figure} \\includegraphics[width=8.8cm]{fig1.eps} \\caption{\\it Top: \\rm IAC80 $I$-band discovery image of the optical afterglow of GRB 080514B (de Ugarte Postigo et al. 2008). The afterglow is highlighted. \\it Bottom: \\rm Keck $R$-band image obtained 24 days after the trigger. The underlying host galaxy is clearly detected. The zoom-in of the Keck image shows the host galaxy (encircled) and the GRB position (red circle).} \\label{fig:finding} \\end{figure}  To date no burst detected above 30 MeV has an observed afterglow. The discovery of GRB 080514B by the Italian \\agile gamma-ray satellite (Tavani et al. \\cite{Tavani}) on May 14, 2008 at 09:55:56 UT (Rapisarda et al. 2008) was therefore of particular interest. \\agile carries three instruments covering the energy range from 20 keV to 50 GeV and detected GRB 080514B at energies well above 30 MeV (Giuliani et al. \\cite{Giuliania},\\cite{Giulianib}). GRB 080514B was a bright, multi-spiked event with a duration (T90) of 5.6 s in the energy range from 350 to 1000 keV, characterizing it as a long burst.  The burst was also observed by Mars Odyssey, operating as part of the Interplanetary Network (IPN) (Hurley et al. 2006), making it possible to constrain the error box of the burst to about 100 arcmin$^2$ (Rapisarda et al. 2008). This localization led to the discovery of its X-ray afterglow by the \\swift satellite at coordinates R.A., Dec. (J2000) =  $21^{\\rm h} 31^{\\rm m} 22 \\fs62, +00^\\circ 42\\arcmin 30\\farcs3$ with an uncertainty of 1\\farcs6 (radius, 90\\% confidence) at 0.43 days after the trigger (Page et al. 2008). Before the announcement of the X-ray afterglow position, however, the optical afterglow had already been discovered by our group by observing the complete IPN error box (de Ugarte Postigo et al. \\cite{Postigoa}, \\cite{Postigob}; fig.\\ref{fig:finding}). Here we report our optical/near-infrared follow-up observations of the afterglow of GRB 080514B starting 0.43 days after the trigger, extending to late times of 24 days.  \\begin{table} \\caption{Log of observations. The first column gives the mid-time in days after the GRB. Vega magnitudes are not corrected for Galactic extinction. Also given are 3$\\sigma$ upper limits.} \\begin{tabular}{rlccl} \\toprule $t$ (days)& filter & instr./telesc. & exposure (s) & mag      \\\\ \\midrule 0.430 &$UVW1$    & UVOT      & 284      & 20.45  $\\pm$ 0.40 \\\\ 0.432 & $U$\t  & UVOT      & 142      & 19.75  $\\pm$ 0.30 \\\\ 0.434 & $B$\t  & UVOT      & 142      & 21.00  $\\pm$ 0.63 \\\\ 0.438 &$UVW2$    & UVOT      & 568      & 21.47  $\\pm$ 0.56 \\\\ 0.443 & $V$\t  & UVOT      & 142      & 20.88  $\\pm$ 1.50 \\\\ 0.446 &$UVM2$    & UVOT      & 413      & 21.97  $\\pm$ 1.30 \\\\ 0.488 &$UVW1$    & UVOT      & 419      & 20.17  $\\pm$ 0.27 \\\\ 0.492 & $U$\t  & UVOT      & 209      & 19.91  $\\pm$ 0.27 \\\\ 0.494 & $B$\t  & UVOT      & 209      & 20.40  $\\pm$ 0.31 \\\\ 0.501 &$UVW2$    & UVOT      & 838      & 21.78  $\\pm$ 0.57 \\\\ 0.507 & $V$\t  & UVOT      & 209      & 19.74  $\\pm$ 0.42 \\\\ 0.512 &$UVM2$    & UVOT      & 616      & 20.63  $\\pm$ 0.37 \\\\ 0.555 &$UVW1$    & UVOT      & 415      & 20.88  $\\pm$ 0.46 \\\\ 0.559 & $U$\t  & UVOT      & 207      & 20.09  $\\pm$ 0.32 \\\\ 0.561 & $B$\t  & UVOT      & 207      & 21.62  $\\pm$ 0.91 \\\\ 0.567 &$UVM2$    & UVOT      & 791      & 22.09  $\\pm$ 0.76 \\\\ 0.640 & $R_C$    & Watcher   & 120x14   & 19.23  $\\pm$ 0.47 \\\\ 0.660 & $R_C$    & Watcher   & 120x15   & 19.89  $\\pm$ 0.56 \\\\ 0.727 & $I_C$    & IAC 80    & 3x300    & 20.26  $\\pm$ 0.21 \\\\ 0.743 & $I_C$    & IAC 80    & 3x300    & 20.59  $\\pm$ 0.20 \\\\ 0.761 & $I_C$    & IAC 80    & 3x300    & 20.16  $\\pm$ 0.16 \\\\ 0.774 & $I_C$    & IAC 80    & 3x300    & 20.03  $\\pm$ 0.14 \\\\ 0.907 & $g^\\prime$   & GROND/2.2m& 3x1501   & 21.53  $\\pm$ 0.04 \\\\ 0.907 & $r^\\prime$   & GROND/2.2m& 3x1501   & 21.16  $\\pm$ 0.03 \\\\ 0.907 & $i^\\prime$   & GROND/2.2m& 3x1501   & 20.77  $\\pm$ 0.08 \\\\ 0.907 & $z^\\prime$   & GROND/2.2m& 3x1501   & 20.43  $\\pm$ 0.05 \\\\ 0.907 & $J$    & GROND/2.2m& 3x1200   & 19.82  $\\pm$ 0.03 \\\\ 0.907 & $H$    & GROND/2.2m& 3x1200   & 19.10  $\\pm$ 0.04 \\\\ 0.907 & $K$    & GROND/2.2m& 2x1200   & $>$17.5          \\\\ 1.021 & $J$    & NEWFIRM/KPNO&23x30x2 & 19.84  $\\pm$ 0.14 \\\\ 1.038 & $J$    & NEWFIRM/KPNO&15x30x2 & 20.06  $\\pm$ 0.07 \\\\ 1.763 & $R_C$  & NOT       & 1x300    & 22.31  $\\pm$ 0.08 \\\\ 1.782 & $B$    & NOT       & 1x300    & 23.03  $\\pm$ 0.13 \\\\ 1.798 & $I_C$  & NOT       & 1x300    & 22.00  $\\pm$ 0.10 \\\\ 1.899 & $i'$  & GMOS/Gemini & 1x200    & 21.83  $\\pm$ 0.06 \\\\ 1.993 & $g^\\prime$   & GROND/2.2m& 1x1501   & 22.74   $\\pm$0.08  \\\\ 1.993 & $r^\\prime$   & GROND/2.2m& 1x1501   & 22.38  $\\pm$ 0.10 \\\\ 1.993 & $i^\\prime$   & GROND/2.2m& 1x1501   & 21.78  $\\pm$ 0.13 \\\\ 1.993 & $z^\\prime$   & GROND/2.2m& 1x1501   &  $>$21.6    \\\\ 1.993 & $J$    & GROND/2.2m&  1x1200  & $>$20.3 \t \\\\ 1.993 & $H$    & GROND/2.2m&  1x1200  & $>$19.1 \t \\\\ 1.993 & $K$    & GROND/2.2m&  1x1200  & $>$17.7 \t \\\\ 2.023  & $J$    & NEWFIRM/KPNO& 15x30x2& 20.95  $\\pm$ 0.30 \\\\ 2.039  & $H$    & NEWFIRM/KPNO& 15x15x4& $>$20.3           \\\\ 2.536 & White  & UVOT      & 5361     & 22.19  $\\pm$ 0.17 \\\\ 8.965 & $g^\\prime$   & GROND/2.2m& 4x1501   & 24.05  $\\pm$ 0.17 \\\\ 8.965 & $r^\\prime$   & GROND/2.2m& 4x1501   & 24.40  $\\pm$ 0.25 \\\\ 8.965 & $i^\\prime$   & GROND/2.2m& 4x1501   & 23.35  $\\pm$ 0.26 \\\\ 8.965 & $z^\\prime$   & GROND/2.2m& 4x1501   & 23.28  $\\pm$ 0.24 \\\\ 8.965 & $J$    & GROND/2.2m&  4x1200   &  $>$21.9 \\\\ 8.965 & $H$    & GROND/2.2m&  4x1200   &  $>$20.5 \\\\ 8.965 & $K$    & GROND/2.2m&  3x1200   &  $>$18.4 \\\\ 24.13  & $R_C$  & Keck      &    960    & 24.17  $\\pm$ 0.33 \\\\ 24.13  & $g^\\prime$    & Keck &  1080      & 24.73  $\\pm$ 0.34 \\\\ \\bottomrule \\end{tabular} \\label{mags} \\end{table} ",
        "conclusions": "To our knowledge, GRB 080514B is the first gamma-ray burst detected above 30 MeV for which an afterglow has been found in the X-ray band and in the optical/NIR bands. Based on our ground-based follow-up observing campaign, in combination with \\emph{Swift} UVOT and \\emph{Swift} XRT data starting 0.4 days after the burst, we find: (1) The X-ray/optical/NIR light curve after 0.4 days is well described by a single power-law with no sign of a jet break. (2) The SED of the afterglow indicates strong Lyman blanketing at short wavelengths, implying a photometric redshift of $z=1.8^{+0.4}_{-0.3}$. This is the first redshift determination for a GRB with prompt emission detected beyond 30 MeV. We find no evidence for extinction by dust in the GRB host galaxy. (3) A comparison of the observed light curve decay with the spectral energy distribution favours a model in which the afterglow blast wave propagated in a wind medium. (4) The intrinsic properties of the optical afterglow are typical for other long-duration GRBs. (5) The putative GRB host galaxy, identified in our GROND and Keck images, has $R_c=24.2 \\pm 0.3$ and an absolute $R$-band magnitude of $M_R=-20.9 \\pm 0.3$. (6) The optical transient was offset from the center of its host by $2.6 \\pm 1.7$ kpc.  We conclude that according to our data set the afterglow properties as well as the properties of the host galaxy match into what is known about the corresponding properties of the long burst sample. The only property that make this burst remarkable is its detection above 30 MeV."
    },
    "0808/0808.2648_arXiv.txt": {
        "abstract": "Using the {\\it HST} ACS, we have obtained deep optical images reaching well below the oldest main sequence turnoff in fields on the southeast minor-axis of the Andromeda Galaxy, 35~kpc from the nucleus.  These data probe the star formation history in the extended halo of Andromeda -- that region beyond 30~kpc that appears both chemically and morphologically distinct from the metal-rich, highly-disturbed inner spheroid.  The present data, together with our previous data for fields at 11 and 21~kpc, do not show a simple trend toward older ages and lower metallicities, as one might expect for populations further removed from the obvious disturbances of the inner spheroid. Specifically, the mean ages and [Fe/H] values at 11 kpc, 21 kpc, and 35 kpc are 9.7~Gyr and $-0.65$, 11.0~Gyr and $-0.87$, and 10.5~Gyr and $-0.98$, respectively.  In the best-fit model of the 35~kpc population, one third of the stars are younger than 10~Gyr, while only $\\sim$10\\% of the stars are truly ancient and metal-poor.  The extended halo thus exhibits clear evidence of its hierarchical assembly, and the contribution from any classical halo formed via early monolithic collapse must be small. ",
        "introduction": "One of the primary quests of astronomy is measuring the star formation history of giant galaxies.  The most direct tool for such work is a color-magnitude diagram (CMD) reaching low-mass stars below the main sequence turnoff.  The Advanced Camera for Surveys (ACS) on the {\\it Hubble Space Telescope (HST)} enabled the application of this technique in the Andromeda Galaxy (M31), the nearest giant galaxy to our own.  In a series of deep surveys, we have been exploring the star formation histories in M31's diverse structures (Brown et al.\\ 2003, 2006, 2007).  The importance of M31 as a target for such studies is twofold. First, M31 is the only giant spiral galaxy where we can resolve the stellar main sequence with an external vantage point, thus overcoming the large uncertainties due to reddening and distance that can hamper studies of Milky Way field populations, most of which are too distant for accurate parallaxes.  Second, recent studies suggest that M31 is more representative of large spiral galaxies than the Milky Way.  By comparing their locations in the planes spanned by various parameters (rotational velocity, absolute $K$ luminosity, disk angular momentum, [Fe/H] in the outskirts, disk scale length, etc.), Hammer et al.\\ (2007) argue that M31 is more ``typical'' of large spirals than the Milky Way, which is systematically offset by $\\sim$1$\\sigma$ from the dominant trends in these planes (see also Flynn et al.\\ 2006).  They attribute this distinction to an unusually quiescent merger history in the Milky Way.  In contrast, the active merger history in M31 is supported by star count maps (e.g., Ferguson et al.\\ 2002; Ibata et al.\\ 2007) showing a variety of substructure, including a giant stellar stream (GSS) from a cannibalized satellite (Ibata et al.\\ 2001).  Going back to the work of Mould \\& Kristian (1986), the M31 ``halo'' was considered unusual because of its relatively high metallicity ($\\sim$10 times higher than that in the Milky Way halo). Brown et al.\\ (2003, 2006) have since demonstrated that these same metal-rich populations were of intermediate age ($\\sim$2--10~Gyr). There is now strong evidence that the inner spheroid (within 15~kpc) is polluted by the progenitor of the GSS; N-body simulations (Fardal et al.\\ 2007) and kinematical surveys (Gilbert et al.\\ 2007) demonstrate that much of the substructure in the inner spheroid is due to debris dispersed from the GSS's progenitor, while deep CMDs (Brown et al.\\ 2006) show strong similarities between the star formation histories of the GSS and inner spheroid.  Recently, two groups independently discovered a vast extended halo in M31 (Guhathakurta et al.\\ 2005; Irwin et al.\\ 2005), spanning radii of 30--165~kpc and reaching lower metallicities similar to those in the Milky Way halo (Kalirai et al.\\ 2006; Koch et al.\\ 2007). Our purpose here is to measure the star formation history in this extended halo and to compare it to that in the interior fields.  We have used the surface brightness profile of Guhathakurta et al.\\ (2005) as a guide to this exploration (Figure 1). Brown et al.\\ (2003, 2006) explored a field at 11~kpc on the minor axis, a metal-rich region of the spheroid that resembles a bulge, with significant debris from the merger that produced the GSS. Brown et al.\\ (2007) investigated a field at 21~kpc, falling in the transition zone between the inner and outer spheroid.  Here we investigate two fields at 35~kpc, where the extended halo begins to dominate.  \\begin{figure}[t] \\epsscale{1.1} \\plotone{f1.eps} \\epsscale{1.1} \\caption{The minor-axis surface-brightness profile ({\\it crosses}, Pritchet \\& van den Bergh 1994; {\\it error bars}, Guhathakurta et al.\\ 2005), along with one possible decomposition into disk, Sersic bulge, and $r^{-2.3}$ power-law halo (from Guhathakurta et al.\\ 2005).  A clear break in the surface-brightness profile occurs at 20--30 kpc.  Our minor-axis fields ({\\it grey shading}) sample the spheroid (i.e., bulge + halo) on either side of this transition region and within it.} \\label{proffig} \\end{figure} ",
        "conclusions": "We have measured the star formation history in three regions of the M31 spheroid, spanning a large range in distance on the minor axis (11~kpc, 21~kpc, and 35~kpc from the nucleus).  The [Fe/H] decreases monotonically with increasing radius, with mean values of $-0.65$, $-0.87$, and $-0.98$, respectively.  However, the age does not monotonically increase with increasing radius, having mean values of 9.68, 10.98, and 10.45~Gyr, respectively.  In their investigations of a metallicity gradient in the spheroid, Koch et al.\\ (2007) find that $<$[Fe/H]$>$ falls more rapidly with radius than Kalirai et al.\\ (2006).  Although our $<$[Fe/H]$>$ at 35~kpc is in good agreement with the trend of Kalirai et al.\\ (2006) and nearly 0.5 dex higher than that of Koch et al.\\ (2007), the latter study includes 15 stars in the vicinity of our ACS fields, and for these, the $<$[Fe/H]$>$ agrees well with our value.  If the Koch et al.\\ (2007) trend is representative of the general gradient in the spheroid, it would imply the ACS fields are in a local region of systematically higher [Fe/H].  Note that our fit uses the lower RGB, which is less sensitive to [Fe/H] than the upper RGB (used by Kalirai et al.\\ and Koch et al.), but it provides a larger RGB sample and avoids the need to screen foreground dwarf contamination.  Our fits to the populations in the inner spheroid, the transition zone, and the outer spheroid all require the presence of intermediate-age stars (age $<$~10~Gyr).  Such populations are expected if the spheroid formed via hierarchical assembly.  This picture is also supported by recent star-count maps (Ibata et al.\\ 2007) that show several stream-like structures crossing the minor axis beyond a distance of 40~kpc from the nucleus; relative to the bright ridge of stars in their Stream D (the nearest of these streams), our fields fall $\\sim$30$^\\prime$ (7 kpc) toward the interior, and may conceivably include debris from these streams. Indeed, Bullock \\& Johnston (2005) argue that spheroids form inside-out, and that substructure should be abundant in the outer regions of the spheroid.  A classical halo that is ancient and metal poor, as might be expected from an early monolithic collapse, does not comprise a large fraction of the population at 35~kpc. If we arbitrarily define such an ancient metal-poor population to be at [Fe/H]$<-1.5$ and age$\\ge 12$~Gyr (considered representative of the outer regions of the Milky Way halo), the fraction of such stars in our best-fit model is only $\\sim$10\\%.  Recent surveys (e.g., Ibata et al.\\ 2007) imply it will be difficult to find a truly ``clean'' region of the spheroid that is not criss-crossed by the debris of mergers."
    },
    "0808/0808.1227_arXiv.txt": {
        "abstract": "We present a new semi-analytic model that self-consistently traces the growth of supermassive black holes (BH) and their host galaxies within the context of the \\LCDM\\ cosmological framework. In our model, the energy emitted by accreting black holes regulates the growth of the black holes themselves, drives galactic scale winds that can remove cold gas from galaxies, and produces powerful jets that heat the hot gas atmospheres surrounding groups and clusters. We present a comprehensive comparison of our model predictions with observational measurements of key physical properties of low-redshift galaxies, such as cold gas fractions, stellar metallicities and ages, and specific star formation rates.  We find that our new models successfully reproduce the exponential cutoff in the stellar mass function and the stellar and cold gas mass densities at $z\\sim0$, and predict that star formation should be largely, but not entirely, quenched in massive galaxies at the present day. We also find that our model of self-regulated BH growth naturally reproduces the observed relation between BH mass and bulge mass.  We explore the global formation history of galaxies in our models, presenting predictions for the cosmic histories of star formation, stellar mass assembly, cold gas, and metals. We find that models assuming the ``concordance'' \\LCDM\\ cosmology overproduce star formation and stellar mass at high redshift ($z\\gtrsim 2$). A model with less small-scale power predicts less star formation at high redshift, and excellent agreement with the observed stellar mass assembly history, but may have difficulty accounting for the cold gas in quasar absorption systems at high redshift ($z\\sim 3$--4). ",
        "introduction": "\\label{sec:intro}  It is now well-established that the Cold Dark Matter paradigm for structure formation \\citep{bfpr:84}, in its modern, dark energy-dominated (\\LCDM) incarnation, provides a remarkably successful paradigm for interpreting a wide variety of observations, from the cosmic microwave background fluctuations at $z\\sim 1000$ \\citep{spergel:03,spergel:07}, to the large-scale clustering of galaxies at $z\\sim 0$ \\citep{percival:02,tegmark:04,eisenstein:05}. Other successful predictions of the \\LCDM\\ paradigm include the cosmic shear field as measured by weak gravitational lensing \\citep{heymans:05,bacon:05,hoekstra:06}, the small-scale power spectrum as probed by the Lyman-alpha forest \\citep{desjacques:05,jena:05}, and the number densities of \\citep{borgani:01}, and baryon fractions within, galaxy clusters \\citep{allen:04,white:93}.  \\LCDM\\ as a paradigm for understanding and simulating galaxy formation has had more mixed success. In this picture, as originally proposed by \\citet{wr:78} and \\citet{bfpr:84}, galaxies form when gas cools and condenses at the centers of dark-matter dominated potential wells, or ``halos''. More detailed calculations, using semi-analytic and numerical simulations of galaxy formation, have shown that this framework does provide a promising qualitative understanding of many features of galaxies and their evolution \\citep[e.g.,][]{kwg:93,cafnz:94,sp:99,kauffmann:99,cole:00,spf:01,baugh:06}. However, it has been clear for at least a decade now that there is a fundamental tension between certain basic predictions of the \\LCDM-based galaxy formation paradigm and some of the most fundamental observable properties of galaxies. In this paper we focus on two interconnected, but possibly distinct, problems: 1) the ``overcooling'' or ``massive galaxy'' problem and 2) the star formation ``quenching'' problem.  The first problem is manifested by the fact that both semi-analytic and numerical simulations predict that a large fraction of the available baryons in the Universe rapidly cools and condenses, in conflict with observations which indicate that only about $\\sim 5$--10 \\% of the baryons are in the form of cold gas and stars \\citep{bell:03a,fukugita_peebles:04}.  This is due to the fact that gas at the densities and temperatures characteristic of dark matter halos is expected to cool rapidly, and is related to the classical ``cooling flow'' problem \\citep{fabian_nulsen:77,cowie_binney:77,mathews:78}. Direct observations of the X-ray properties of hot gas in clusters similarly imply that this gas should have short cooling times, particularly near the center of the cluster, but the condensations of stars and cold gas at the centers of these clusters are much smaller than would be expected if the hot gas had been cooling so efficiently over the lifetime of the cluster \\citep[for reviews see][]{fabian:94,peterson:06}. Moreover, X-ray spectroscopy shows that very little gas is cooling below a temperature of about one-third of the virial temperature of the cluster \\citep{peterson:03}.  A further difficulty is that there is a fundamental mismatch between the \\emph{shape} of the dark matter halo mass function and that of the observed mass function of cold baryons (cold gas and stars) in galaxies \\citep{sp:99,benson:03}. The galaxy mass function has a sharp exponential cutoff above a mass of about a few times $10^{10} \\msun$, while the halo mass function has a shallower, power-law cutoff at much higher mass ($\\sim {\\rm few} \\times 10^{13}\\msun$). There is also a mismatch at the small-mass end, as the mass function of dark matter halos is much steeper than that of galaxies. If we are to assume that each dark matter halo hosts a galaxy, this then implies that the ratio between the luminosity or stellar mass of a galaxy and the mass of its dark matter halo varies strongly and non-monotonically with halo mass \\citep{kravtsov:04a,conroy:06,wang:06,moster:08}, such that ``galaxy formation'' is much more inefficient in both small mass and large mass halos, with a peak in efficiency close to the mass of our own Galaxy, $\\sim 10^{12} \\msun$. On very small mass scales (below halo velocities of $\\sim 30-50 \\kms$), the collapse and cooling of baryons may be suppressed by the presence of a photoionizing background \\citep{efstathiou:92,thoul:96,quinn:96}. For larger mass halos (up to $V_{\\rm vir} \\simeq 150-200 \\kms$), the standard assumption is that winds driven by massive stars and supernovae are able to heat and expell gas, resulting in low baryon fractions in small mass halos \\citep{wr:78,dekel_silk:86,white_frenk:91}.  However, stellar feedback probably cannot provide a viable solution to the overcooling problem in massive halos \\citep{benson:03}: stars do not produce enough energy to expell gas from these large potential wells; and the massive, early type galaxies in which the energy source is needed have predominantly old stellar populations and little or no ongoing or recent star formation. Other solutions, like thermal conduction, have been explored, but probably do not provide a full solution \\citep{benson:03,voigt:04}.  The second problem is related to the correlation of galaxy structural properties (morphology) and spectrophotometric properties (stellar populations) with stellar mass. The presence of such a correlation has long been known, in the sense that more massive galaxies tend to be predominantly spheroid-dominated, with red colors, old stellar populations, low gas fractions, and little recent star formation, while low-mass galaxies tend to be disk-dominated and gas-rich, with blue colors and ongoing star formation \\citep[e.g.][]{roberts:94}. More recently, with the advent of large galaxy surveys such as the Sloan Digital Sky Survey (SDSS), we have learned that the galaxy color distribution (and that of other related properties) is strongly \\emph{bimodal} \\citep[e.g.][]{baldry:04}, and that the transition in galaxy properties from star-forming disks to ``dead'' spheroids occurs rather sharply, at a characteristic stellar mass of $\\sim 3 \\times 10^{10} \\msun$ \\citep{kauffmann:03a,brinchmann:04}. In contrast, the ``standard'' \\LCDM-based galaxy formation models predict that massive halos have been assembled relatively recently, and should contain an ample supply of new fuel for star formation. These models predict an \\emph{inverted} color-mass and morphology-mass relation (massive galaxies tend to be blue and disk dominated, rather than red and spheroid dominated) and no sharp transition or strong bimodality. Thus, the standard paradigm of galaxy formation does not provide a physical explanation for the ``special'' mass scale (a halo mass of $\\sim 10^{12} \\msun$, or a stellar mass of $\\sim 3\\times 10^{10} \\msun$) which marks both the peak of galaxy formation efficiency and the transition in galaxy properties seen in observations.  Several pieces of observational evidence provide clues to the solution to these problems. It is now widely believed that every spheroid-dominated galaxy hosts a nuclear supermassive black hole (SMBH), and that the mass of the SMBH is tightly correlated with the luminosity, mass, and velocity dispersion of the stellar spheroid \\citep{kormendy_richstone:95,magorrian:98,ferrarese:00,gebhardt:00,marconi:03,haering:04}. These correlations may be seen as a kind of ``fossil'' evidence that black holes were responsible for regulating the growth of galaxies or vice versa. This also implies that the most massive galaxies, where quenching is observed to be the most efficient, host the largest black holes, and therefore the available energy budget is greatest in precisely the systems where it is needed, in contrast to the case of stellar feedback. The integrated energy released over the lifetime of a SMBH ($\\simeq 10^{60}-10^{62}$ erg) is clearly very significant compared with galaxy binding energies \\citep{silk_rees:98}. In view of these facts, it seems almost inconceivable that AGN feedback is \\emph{not} important in shaping galaxy properties.  However, in order to build a complete, self-consistent machinery to describe the formation and growth of black holes within the framework of a cosmological galaxy formation model, and to attempt to treat the impact of the energy feedback from black holes in this context, we need to address several basic questions: 1) When, where, and with what masses do seed black holes form? 2) What triggers black hole accretion, what determines the efficiency of this accretion, and what shuts it off?  3) In what form is the energy produced by the black hole released, and how does this energy couple with the host galaxy and its surroundings? In order to address some of these questions, we first identify two modes of AGN activity which have different observational manifestations, probably correspond to different accretion mechanisms, and have different physical channels of interaction with galaxies.  \\subsection{The Bright Mode of Black Hole Growth} \\label{sec:bright}  Classical luminous quasars and their less powerful cousins, optical or X-ray bright AGN, radiate at a significant fraction of their Eddington limit \\citep[$L \\sim (0.1-1) L_{\\rm Edd}$;][]{vestergaard:04,kollmeier:06}, and are believed to be fed by optically thick, geometrically thin accretion disks \\citep{shakura:73}. We will refer to this mode of accretion as the ``bright mode'' because of its relatively high radiative efficiency (with a fraction $\\etabh \\sim 0.1-0.3$ of the accreted mass converted to radiation). The observed space density of these quasars and AGN is low compared to that of galaxies, implying that if most galaxies indeed host a SMBH, this ``bright mode'' of accretion is only ``on'' a relatively small fraction of the time. Constraints from quasar clustering and variability imply that quasar lifetimes must be $\\lesssim 10^{8.5}$ yr \\citep{martini:01,martini:03}. These short timescales combined with the large observed luminosities immediately imply that fueling these objects requires funneling a quantity of gas comparable to the entire supply of a large galaxy ($\\sim 10^{9}-10^{10} \\msun$) into the central regions on a timescale of order the dynamical time ($\\sim {\\rm few} \\times 10^7-10^8$ yr).  These considerations alone lead one to consider galaxy-galaxy mergers as a promising mechanism for triggering this efficient accretion onto nuclear black holes. The observational association of mergers with enhanced star formation, particularly with the most violent observed episodes of star formation exhibited by Ultra Luminous Infrared Galaxies (ULIRGS), is well-established \\citep{sanders:96,farrah:01,colina:01,barton:00,woods:06,woods:07,barton:07,lin:07,li:07}. Moreover, numerical simulations have shown that tidal torques during galaxy mergers can drive the rapid inflows of gas that are needed to fuel both the intense starbursts and rapid black hole accretion associated with ULIRGS and quasars \\citep{hernquist:89,barnes:92,barnes:96,mihos:94,mihos:96,springel:05a,dimatteo:05}. As well, it seems that if one can probe sufficiently deep to study the SED beneath the glare of the quasar, one always uncovers evidence of young stellar populations indicative of a recent starburst \\citep{brotherton:99,canalizo:01,kauffmann:03b,jahnke:04,sanchez:04,vandenberk:06}. Near-equal mass (major) mergers also have the attractive feature that they scramble stars from circular to random orbits, leading to morphological transformation from disk to spheroid \\citep{toomre:72,barnes:88,barnes:92,hernquist:92,hernquist:93}. If spheroids and black holes both arise from violent mergers, this provides a possible explanation for why black hole properties always seem to be closely associated with the spheroidal components of galaxies.  What impact does the energy associated with this rapid, bright mode growth have on the galaxy and on the growth of the black hole itself? Long thought to be associated only with a small subset of objects (e.g. Broad Absorption Line (BAL) quasars), high-velocity winds have been detected in a variety of different types of quasar systems \\citep{dekool:01,pounds:03,chartas:03,pounds:06}, and are now believed to be quite ubiquitous \\citep{ganguly:08}. However, their impact on the host galaxy remains unclear, as the mass outflow rates of these winds are difficult to constrain \\citep[though see][]{steenbrugge:05,chartas:07,krongold:07}.  Recently, numerical simulations of galaxy mergers including black hole growth found that depositing even a small fraction ($\\sim 5$ \\%) of the energy radiated by the BH into the ISM can not only halt the accretion onto the BH, but can drive large-scale winds \\citep{dimatteo:05}. These winds sweep the galaxy nearly clean of cold gas and halt further star formation, leaving behind a rapidly reddening, spheroidal remnant \\citep{springel:05b}.  To study how the interplay between feedback from supermassive black hole accretion and supernovae, galaxy structure, orbital configuration, and gas dissipation combine to determine the properties of spheroidal galaxies formed through mergers, hundreds of hydrodynamical simulations were performed by Robertson et al. (\\citeyear{robertson:06a,robertson:06b,robertson:06c}) and Cox et al. (\\citeyear{cox:06a,cox:06c}) using the methodology presented by \\citet{dimatteo:05} and \\citet{springel:05a}. Robertson et al. and Cox et al. analyzed the merger remnants to study the redshift evolution of the BH mass-$\\sigma$ relation, the Fundamental Plane, phase-space density, and kinematic properties. This extensive suite has been supplemented by additional simulations of minor mergers from \\citet{cox:08}. Throughout the rest of this paper, when we refer to ``the merger simulations'', we refer to this suite.  Based on their analysis of these simulations, \\citet{hopkins_cosmo1:07} have outlined an evolutionary sequence from galaxy-galaxy merger, to dust-enshrouded starburst and buried AGN, blow-out of the dust and ISM by the quasar- and starburst-driven winds, to classical (unobscured) quasar, post-starburst galaxy, and finally ``dead'' elliptical. \\citet{hopkins_bhfpth:07} find that in the merger simulations, the accretion onto the BH is eventually halted by a pressure-driven outflow. Because the depth of the spheroid's potential well determines the amount of momentum necessary to entrain the infalling gas, \\citet{hopkins_bhfpth:07} find that this leads to a ``Black Hole Fundamental Plane'', a correlation between the final black hole mass and sets of spheroid structural/dynamical properties (mass, size, velocity dispersion) similar to the one seen in observations \\citep{hopkins_bhfpobs:07,marconi:03}.  Furthermore, Hopkins et al. (\\citeyear{hopkins:05a,hopkins:05b,hopkins_qsolifetime:05}) have shown that the self-regulated nature of black hole growth in these simulations leads to a characteristic form for quasar lightcurves. As the galaxies near their final coalescence, the accretion rises to approximately the Eddington rate. After the critical black hole mass is reached and the outflow phase begins, the accretion rate enters a power-law decline phase. Although most of their growth occurs in the near-Eddington phase, quasars spend much of their time in the decline phase, and this implies that many observed low-luminosity quasars are actually relatively massive black holes in the last stages of their slow decline. \\citet{hopkins_rss:06} found that when these lightcurves are convolved with the observed mass function of merging galaxies, the predicted AGN luminosity function is consistent with observations. Moreover, Hopkins et al. (\\citeyear{hopkins_qsolifetime:05,hopkins_optx:05,hopkins_unified:06,hopkins_redgal:06}) have shown that this picture reproduces many quasar and galaxy observables that are difficult to account for with more simplified assumptions about QSO lightcurves, such as differences in the quasar luminosity function in different bands and redshifts, Eddington ratio and column density distributions, the X-ray background spectrum, and relic red, early type galaxy population colors and distributions.  \\subsection{The Radio Mode}  The second mode of AGN activity is much more common, and in general less dramatic. A fairly large fraction of massive galaxies (particularly galaxies near the centers of groups and clusters) are detected at radio wavelengths \\citep{best:05, best:07}. Most of these radio sources do not have emission lines characteristic of classical optical or X-ray bright quasars \\citep{best:05,kauffmann:07}, and their accretion rates are believed to be a small fraction of the Eddington rate \\citep{rafferty:06}. They are extremely radiatively inefficient \\citep{birzan:04}, and thought to be fuelled by optically thin, geometrically thick accretion as expected in ADAF and ADIOS models such as those proposed by \\citet{narayan:94} and \\citet{blandford:99}. Because these objects are generally identified via their radio emission, we refer to this mode of accretion and BH growth as the ``radio mode'' \\citep[following][]{croton:06}.  Although these black holes seem to be inefficient at producing radiation, they can apparently be quite efficient at producing kinetic energy in the form of relativistic jets. Intriguingly, the majority of cooling flow clusters host these active radio galaxies at their centers \\citep{dunn:06,dunn:08}, and X-ray maps reveal that the radio lobes are often spatially coincident with cavities, thought to be bubbles filled with relativistic plasma and inflated by the jets \\citep[][and references therein]{mcnamara:07}. The observations of these bubbles can be used to estimate the work required to inflate them against the pressure of the hot medium \\citep{birzan:04,rafferty:06,allen:06}, and hence obtain lower limits on the jet power.  While the idea that radio jets provide a heat source that could counteract cooling flows has been discussed for many years \\citep[e.g.][]{binney_tabor:95,churazov:02,fabian:03,omma:04a,binney:04}, these observations now make it possible to investigate more quantitatively whether the heating rates are sufficient to offset the cooling rates in groups and clusters. Several studies conclude that in the majority of the systems studied, the AGN heating traced by the power in the X-ray cavities alone is comparable to or in excess of the energy being radiated by the cooling gas \\citep{mcnamara:06,rafferty:06,fabian:06,best:06,mcnamara:07,dunn:08}. Moreover, the net cooling rate is correlated with the observed star formation rate in the central cD galaxy, indicating that there may be a self-regulating cycle of heating and cooling \\citep{rafferty:06}.  Several other physical processes that could suppress cooling in large mass halos have been suggested and explored, such as thermal conduction \\citep{benson:03,voigt:04}, multi-phase cooling \\citep{maller_bullock:04}, or heating by sub-structure or clumpy accretion \\citep{khochfar:07,naab:07,dekel_birnboim:08}. While some or all of these processes may well be important, in this paper we will investigate whether it is plausible that ``radio mode'' heating alone can do the job.  \\subsection{A Unified Model for Black Hole Activity and AGN Feedback}  All of this begs the question: what determines whether a black hole accretes in the ``bright mode'' or ``radio mode'' state? An interesting possible answer comes from an analogy with X-ray binaries \\citep{jester:05,koerding:06}. Observers can watch X-ray binaries in real time as they transition between two states: the ``low/hard'' state, in which a steady radio jet is present and a hard X-ray spectrum is observed, and the ``high/soft'' state, in which the jet dissapears and the X-ray spectrum has a soft, thermal component \\citep{maccarone:03,fender:04}. The transition between the two states is thought to be connected to the accretion rate itself: the ``high/soft'' state is associated with accretion rates of $\\gtrsim (0.01$--0.02) $\\dot{m}_{\\rm Edd}$ and the existence of a classical thin accretion disk, while the ``low/hard'' state is associated with lower accretion rates and radiatively inefficient ADAF/ADIOS accretion \\citep{fender:04}.  Recently, \\citet{sijacki:07} have applied this idea in cosmological hydrodynamic simulations, by assuming that when the accretion rate exceeds a critical value, ``bright mode'' feedback occurs (AGN-driven winds), while when the accretion rate is lower, ``radio mode'' (mechanical bubble feedback) is implemented. The results of their simulations appear promising --- they produced black hole and stellar mass densities in broad agreement with observations. In addition, they found that their implementation of AGN feedback was able to suppress strong cooling flows and produce shallower entropy profiles in clusters, and to quench star formation in massive galaxies. However, the very large dynamic range required to treat the growth of black holes and galaxies in a cosmological context --- from the sub-pc scales of the BH accretion disk to the super-Mpc scales of large scale structure --- means that numerical techniques such as these will likely need to be supplemented by semi-analytic or sub-grid methods for some time to come.  Our approach is in many respects very similar in spirit to that of \\citet{sijacki:07}, although of course we are forced to implement both modes of AGN feedback in an even more schematic manner because we are using a semi-analytic model rather than a numerical simulation. We adopt fairly standard semi-analytic treatments of the growth of dark matter halos via accretion and mergers, radiative cooling of gas, star formation, supernova feedback, and chemical evolution. We then adopt the picture of self-regulated black hole growth and bright mode feedback in mergers discussed in \\S\\ref{sec:bright}, and implement these processes in our model using the results extracted from the merger simulations described above. We assume that the radio mode is fueled instead by hot gas in quasi-hydrostatic halos, and that the accretion rate is described by Bondi accretion from an isothermal cooling flow as proposed by \\citet[][NF00]{nulsen_fabian:00}. We calibrate the heating efficiency of the associated radio jets against direct observations of bubble energetics in clusters.  A number of authors have previously explored the formation of black holes and AGN in the context of CDM-based semi-analytic models of varying complexity \\citep{efstathiou_rees:88,kh:00,wyithe_loeb:02,bromley:04,scannapieco:04,volonteri:03,volonteri:05}, and recently several studies have also investigated the impact of AGN feedback on galaxy formation using such models \\citep{cattaneo:06,croton:06,bower:06,menci:06,schawinski:06,kang:06,monaco:07}. The models that we present here differ from previous studies of which we are aware, in two main respects: 1) we implement detailed modelling of self-regulated black hole growth and bright mode feedback based on an extensive suite of numerical simulations of galaxy mergers and 2) we adopt a simple but physical model for radio mode accretion and heating, and calibrate our model against direct observations of accretion rates and radio jet heating efficiencies. We present a broader and more detailed comparison with observations than previous works, and highlight some remaining problems that have not previously been emphasized. As well, unlike most previous studies, we calibrate our models and make our comparisons in terms of ``physical'' galaxy properties such as stellar mass and star formation rate, which can be estimated from observations, rather than casting our results in terms of observable properties such as luminosities and colors. Our results are therefore less sensitive to the details of dust and stellar population modelling, and easier to interpret in physical terms.  The goals of this paper are to present our new models in detail, and to test and document the extent to which they reproduce basic galaxy observations at $z=0$ and the global cosmic histories of the main baryonic components of the Universe.  The structure of the rest of this paper is as follows. In \\S\\ref{sec:model}, we describe the ingredients of our models and provide a table of all of the model parameters. In \\S\\ref{sec:results}, we present predictions for key properties of galaxies at $z\\sim0$ and for the global history of the main baryonic components of the Universe: star formation, evolved stars, cold gas, metals, and black holes. We conclude in \\S\\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude}  We have presented a new semi-analytic model for the self-consistent evolution of galaxies, black holes and AGN in the framework of the Cold Dark Matter model of structure formation. Our models are built on those described in SP99, SPF2001, and subsequent papers, but we present several important new ingredients here:  \\noindent (i) Improved modelling of tidal stripping and destruction of orbiting sub-halos and of dynamical friction and satellite merger timescales.  \\noindent (ii) Improved modelling of disk sizes, including realistic DM halo profiles and the effect of ``adiabatic contraction''.  \\noindent (iii) A more realistic recipe for star formation in quiescent disks, based on the empirical Kennicutt Law and including a surface density threshold for star formation.  \\noindent (iv) Updated modelling of the efficiency and timescale of starbursts, based on an extensive suite of numerical hydrodynamic simulations of galaxy mergers.  \\noindent (v) Tracking of a ``diffuse stellar halo'' component, which is built up of tidally destroyed satellites and stars scattered in mergers.  \\noindent (vi) A self-regulated model for black hole growth and ``bright mode'' accretion ``lightcurves'' based on numerical merger simulations with AGN feedback.  \\noindent (vii) Galaxy-scale AGN-driven winds, based on numerical merger simulations.  \\noindent (viii) Fueling of black holes with hot gas via Bondi accretion, regulated according to the isothermal cooling flow model of \\citet{nulsen_fabian:00}.  \\noindent (ix) Heating by radio jets, calibrated against observations of X-ray cavities in cooling flow clusters.  We explored the predictions of our new models for a broad range of physical properties of galaxies, groups, and clusters at $z\\sim 0$. One key result is an explanation of the origin of the characteristic shape of the galaxy stellar mass or luminosity function. We cast this in terms of the ``star formation efficiency function'' (i.e., the fraction of baryons in stars as a function of host halo mass), which can be derived empirically by mapping between (sub)-halo mass and stellar mass such that the observed stellar mass function is reproduced \\citep{moster:08,wang:06}. We found that, in our models, the shape of this function arises from the fact that supernova-driven winds can more efficiently heat and expell gas in lower mass halos, while radio-mode heating by AGN is more efficient in higher mass halos, because these halos have larger mass black holes and these black holes can accrete hot gas more efficiently. The peak in this function at $M_{\\rm vir} \\sim 10^{12} \\msun$ occurs because these halos are too massive for supernova-driven winds to escape easily, and do not have massive enough black holes or a large enough virial temperature to fuel efficient radio-mode heating.  We showed that our model also reproduces the stellar mass function as a function of galaxy morphology (at least for galaxies more massive than $\\sim 10^{10} \\msun$), the cold gas fractions of disk galaxies as a function of stellar mass and the cold gas mass function, the stellar mass-metallicity relation for galaxies, and the BH mass vs. spheroid mass relation.  Our model for heating by radio jets is similar in concept, but different in detailed implementation, to other models that have been presented in the literature. We tested the basic assumption of our model for fueling of BH by hot gas, the isothermal cooling flow model proposed by \\citet{nulsen_fabian:00}, using observations of central temperatures and densities in hot X-ray emitting gas in nine systems by \\citet{allen:06}, and found consistency. We further compared the jet power required to solve the overcooling/massive galaxy problem in our models with direct measurements of jet power from the energetics of bubbles detected in X-ray gas around elliptical galaxies \\citep{allen:06,rafferty:06}. We found that our required jet powers lie above the observations of \\citet{allen:06} but agree with the higher values obtained by \\citet{rafferty:06}.  We compared the results of our fiducial model with a very simple implementation of AGN heating, in which we simply switched off cooling in halos more massive than $\\sim 10^{12} \\msun$ (the ``Halo Quenching'' or HQ model). We found that this model produced very similar results to our fiducial model for the $z=0$ stellar mass fraction as a function of halo mass ($f_{\\rm star}(M_{\\rm halo})$), the stellar mass function, the cold gas fractions and cold gas mass function, and the global star formation and stellar mass assembly histories.  We investigated whether the same models which provided a good match to the stellar mass function also reproduced the distribution of (specific) star formation rates as a function of stellar mass. We found that our fiducial model qualitatively reproduces the main features of the observed distribution of SSFR vs. stellar mass: the models produce a star forming sequence, which dominates at low stellar masses ($\\lesssim 2-3 \\times 10^{10}\\msun$) and a quenched population at high stellar masses. The transition mass is naturally predicted by the model, and again corresponds to the halo mass scale where AGN heating becomes effective. This is in marked contrast to models without AGN feedback, in which all massive galaxies are actively star-forming.  However, a more detailed comparison with the observations indicates that the low-mass star-forming sequence in the models occurs at too low a specific star formation rate, and does not have the right slope (our star forming sequence is flat in SSFR vs. stellar mass space, rather than tilted such that low mass galaxies have high SSFR, as the observations indicate). This problem is present in all of our models, and is independent of AGN feedback and robust to the star formation recipe and the values of our free parameters. It may be a symptom of the same malady that is responsible for producing low-mass galaxies with older ages than those estimated from absorption lines in the spectra of nearby galaxies.  In the Halo Quenching model, the quenching is clearly too sharp a function of stellar mass. Essentially all galaxies more massive than $m_{\\rm star} \\sim 10^{11} \\msun$ are completely quenched, while the observations indicate that massive galaxies have a wide range of specific star formation rates, from small to moderate amounts of recent star formation to activity levels that place them on the star-forming sequence.  We investigated the cosmic histories of the major baryonic components of the Universe as predicted in our models: star formation and stellar mass, cold gas, warm/hot gas, and metals. We found that our fiducial Concordance \\LCDM\\ model produces very good agreement with the global star formation rate density at $z\\lesssim 2$, but predicts a higher and flatter SF history at $2 \\lesssim z \\lesssim 6$ than most observations indicate.  We found that our prediction of the global SFR density is apparently not affected by AGN feedback at redshifts above $z\\gtrsim 3$, because most star formation is taking place in small mass halos.  All of the C-\\LCDM\\ models predict a significantly larger amount of mass in long-lived stars in galaxies at redshifts $z \\gtrsim 0.5$, by about a factor of two at $z\\sim1$ and a factor of three at $z\\sim 2$ for the fiducial and HQ models. The fact that we reproduce the stellar mass density at $z\\sim0$ (by construction) then implies that there is not enough evolution in the galaxy population between $z\\sim1$ and $z\\sim0$ in our models with AGN feedback. This problem does not seem to be specific to our models, but is common to all the recently published models in which AGN feedback is implemented in a similar way \\citep{cattaneo:06,croton:06,bower:06}. Indeed, one can see from the predictions of the global SFR history presented in these works (e.g. Fig.~5 of Croton et al. 2006, Fig.~8 of Bower et al. 2006) that these models all produce very similar predictions for the global SF history and for the integrated stellar mass density.  The WMAP3 model, in which structure formation occurs later due to the reduced power on small scales, has a much lower and more steeply declining global star formation rate at $z \\gtrsim 2$ (about an order of magnitude lower than C-\\LCDM\\ at $z=6$). The SFR is lower than the observational compilation of \\citet{hopkins_beacom:06} at these redshifts by about 0.2--0.3 dex, but still well within the observational errors. In contrast to the C-\\LCDM\\ model, the stellar mass density predicted by the WMAP3 model is in excellent agreement with observational estimates from $z\\sim 4$--0. The delayed SF and stellar mass assembly history in WMAP3 relative to C-\\LCDM\\ has also been illustrated by \\citet{wang:07}.  We also compared the predicted global mass density of cold gas in galactic disks $\\Omega_{\\rm cold}$ as a function of redshift in our models with estimates of \\HI\\ gas mass density at $2 \\lesssim z \\lesssim 4.5$ from Quasar Absorption Systems (Damped Lyman-$\\alpha$ systems and Lyman limit systems). We found that our C-\\LCDM\\ models had no difficulty producing enough cold gas in disks up to $z\\sim 4$, where the observations are the most secure, but the WMAP3 model did not fare so well. The predicted values of $\\Omega_{\\rm cold}$ in this model were too low by a factor of 2--10 at $z \\gtrsim 3$.  We analyze the redshift-dependent breakdown of all the baryons in our fiducial model into each of the various components that we track: hot gas in halos, warm/hot diffuse gas in the IGM, cold gas in galactic disks, stars in galaxies, and stars in Diffuse Stellar Halos (DSH). We find that hot gas in halos and warm/hot gas in the IGM dominate at all redshifts, in agreement with the predictions of numerical cosmological simulations \\citep[][and references therein]{bertone:08} and with the observational baryon census at low redshift \\citep{fukugita_peebles:04}. Therefore, in order to produce more cold gas at high redshift, we either require more efficient cooling of hot gas or less efficient reheating and ejection of cold gas by supernova feedback. Since our models are already overproducing stars at high redshift, the latter is probably a more promising solution. Our models predict efficient early metal enrichment, with the average stellar mass-weighted stellar metallicity reaching 25\\% of solar at $z\\sim6$ and 50\\% of solar at $z\\sim 2.5$.  The picture of the cosmic build-up of stars that we see in our models is interesting in the context of a problem pointed out in several recent papers \\citep[e.g.][]{hopkins_beacom:06,fardal:07,wilkins:08,dave:08}: when the best available observational estimate of the dust- and incompleteness corrected SFR density is integrated, accounting for gas recycling under the assumption of a universal stellar IMF, the mass of long-lived stars is overestimated by about a factor of 2--3. These authors suggest that a possible resolution to this problem is a non-universal IMF, which was more top-heavy at high redshift. In the context of the two models we have considered, the C-\\LCDM\\ models predict early structure formation, accompanied by a lot of early star formation. We might be able to reconcile these models with all the data if in fact the IMF was top-heavy at early times, so that most of the star formation that we see does not produce long-lived stars. On the other hand, the WMAP3 model implies later structure formation, and hence less star formation at high redshift. This provides an alternate means of reconciling observations of the SF history and stellar mass assembly history, which requires only that the current observational estimates of the SFR at $z \\gtrsim 2$ are too high by about a factor of two to three. While the WMAP3 picture seems more attractive in many respects, it is a concern that the WMAP3 models do not seem to be able to account for DLAS at $z\\gtrsim 3$, again as a consequence of the reduced small scale power. Better constraints on the amount of cold gas at high redshift from new facilities such as ALMA could help to resolve this question.  In agreement with previous work, we have shown that the inclusion of AGN feedback in semi-analytic models can plausibly solve the over-cooling problem, the massive galaxy problem, and the star formation quenching problem in the local Universe --- a huge step forward. However, we have also shown that several potentially serious discrepancies still remain, and we have argued that these discrepancies are not peculiar to our implementation, but are common to all of the CDM-based semi-analytic models currently on the market. These discrepancies are connected with low-mass galaxies ($m_* \\lesssim {\\rm few} \\times 10^{10} \\msun$), however, in the picture that we are developing, small galaxies grow into massive galaxies, and the growth of galaxies and AGN are intimately interconnected, so these problems on small-scales may indicate or cause more pervasive problems. It is likely that these problems are connected to the modelling of cooling, star formation and/or supernova feedback, or possibly to the CDM power spectrum on small scales. \\emph{Direct} AGN feedback\\footnote{By this we mean feedback by an AGN within the galaxy itself. Indirect feedback from AGN in external galaxies, e.g. via pre-heating, may be a promising solution \\citep[e.g.][]{scannapieco:04}} is probably not the solution.  In this paper, we have focussed on predictions of the physical properties of galaxies at $z\\sim0$ and the global evolution of the major baryonic components of the Universe over time. In a planned series of papers, we will investigate the predictions of the models we have presented here for multi-wavelength, observable properties (e.g. UV through FIR luminosities and colors) of galaxies at low and high redshift, and examine in more detail the distribution functions and scaling relations of intrinsic galaxy properties (e.g. stellar mass, star formation rate, metallicity, etc) at high redshift. In addition, we will explore the predictions of our models for AGN properties as a function of redshift and environment, and the relationship between AGN and their host galaxies."
    },
    "0808/0808.3014_arXiv.txt": {
        "abstract": "A microlensing lensing zone refers to the range of planet-star separations where the probability of detecting a planetary signal is high. Its conventional definition as the range between $\\sim 0.6$ and 1.6 Einstein radii of the primary lens is based on the criterion that a major caustic induced by a planet should be located within the Einstein ring of the primary.  However, current planetary lensing searches focus on high-magnification events to detect perturbations induced by another caustic located always within the Einstein ring very close to the primary lens (stellar caustic) and thus a new definition of a lensing zone is needed. In this paper, we determine this lensing zone.  By applying a criterion that detectable planets should produce signals $\\geq 5\\%$, we find that the new lensing zone varies depending on the planet/star mass ratio unlike the fixed range of the classical zone regardless of the mass ratio.  The lensing zone is equivalent to the classical zone for a planet with a planet/star mass ratio $q\\sim 3\\times 10^{-4}$ and it becomes wider for heavier planets.  For a Jupiter-mass planet, the lensing zone ranges from 0.25 to 3.9 Einstein radii, corresponding to a physical range between $\\sim 0.5$ AU and 7.4 AU for a typical Galactic event. The wider lensing zone of central perturbations for giant planets implies that the microlensing method provides an important tool to detect planetary systems composed of multiple ice-giant planets. ",
        "introduction": "The microlensing signal of a planet is a perturbation to the smooth standard light curve of the primary-induced lensing event occurring on a background source star.  The planetary signal lasts for a short period of time: several days for a gas giant planet and several hours for an Earth-mass planet.  Currently, the observational frequency of lensing surveys is $\\sim 1$ per night and thus it is not enough to detect planetary signals.  To achieve the observational frequency required to detect planetary signals, microlensing planet searches are using a combination of survey and follow-up observations.  Survey observations (e.g., OGLE, {http://www.astrouw.edu.pl/$\\sim$ogle}, \\citet{udalski03}; MOA, http://www.physics.canterbury.ac.nz/moa/, \\citet{bond02a}) aim to maximize the event rate by monitoring a large area of the Galactic bulge field on a roughly nightly basis. They issue alerts of ongoing events in the early stage of lensing magnification by analyzing data in real time.  Follow-up observations (e.g., PLANET, http://planet.iap.fr, \\citet{kubas08}; MicroFUN, http://www.astronomy.ohio-state.edu/\\%7Emicrofun/, \\citet{dong06}) are focused on the alerted events in order to detect short-lived planetary signals.   \\begin{figure*}[t] \\epsscale{0.85} \\plotone{f1.eps} \\caption{\\label{fig:one} Maps of magnification pattern involved with a point source in the central region of planetary systems with various planet-star separations $s$ and planet/star mass ratio $q$.  The coordinates $(\\xi,\\eta)$ represent the axes that are parallel with and normal to the planet-star axis, respectively.  In each map, the star is at the center and the planet is on the left.  The closed figures drawn in black curves represent the caustics.  Greyscale is drawn such that brighter tone represents the region of a higher magnification. }\\end{figure*}   In practice, however, the number of telescopes available for follow-up observations is far less for intensive followups of all alerted events and thus priority is given to those events which will maximize the planetary detection probability.  Currently, the highest priority is given to high-magnification events \\citep{bond02b, abe04, rattenbury02}.  This is based on the fact that in addition to a major caustic located away from the primary lens (planetary caustic), a planet induces an additional tiny caustic in the region close to the primary lens (stellar or central caustic).  Due to its location, the perturbations induced by the stellar caustic always occur near the peak of high-magnification events and this makes follow-up observations right on the period of greatest sensitivity possible. The strategy focusing on high-magnification events was proposed by \\citet{griest98} and the adoption of the strategy has led to the discoveries of six planet candidates (OGLE-2005-BLG-071Lb, OGLE-2005-BLG-169Lb, OGLE-2006-BLG-109Lb,c, MOA-2007-BLG-400b, MOA-2007-BLG-192) among the total eight microlensing planet candidates discovered to date \\citep{bond04, udalski05, beaulieu06, gould06, gaudi08, dong08, bennett08}.   In planetary microlensing, a `lensing zone' refers to the range of planet-star separations where the probability of detecting a planetary signal is high.  It was first discussed by \\citet{gould92}, who found that the planet detection probability is maximized when the planet is near the Einstein radius of the primary star.  Later, the definition of the lensing zone is refined by \\citet{wambsganss97} and \\citet{griest98} as the range between approximately 0.6 and 1.6 Einstein radii of the primary lens.  In defining the lensing zone, the location of a caustic is important because perturbations can be covered by follow-up observations when they occur during the progress of the lensing magnification induced by the primary star.\\footnote{We note that events are occasionally monitored in current follow-up observations even after the source has exited the Einstein ring. An example is the event OGLE-2005-BLG-390 \\citep{beaulieu06} for which the planet was detected at the moment when the source was about to exit the ring. For these events, the lensing zone is more extended and broader than the classical definition of the lensing zone.} When a planet is located within this range, not only the size of the planet-induced caustic is maximized but also the caustic is located inside the Einstein ring of the primary lens.  However, this definition of the lensing zone is based only on the location of the planetary caustic.  Considering that current lensing observations focus on perturbations induced by stellar caustics and these caustics are always located within the Einstein ring regardless of the planet-star separation, a new definition of the lensing zone is needed.  In this paper, we determine the range of a lensing zone for microlensing planets detectable through the channel of high-magnification events.   The paper is organized as follows.  In \\S\\ 2, we briefly describe the properties of planet-induced caustics.  In \\S\\ 3, we investigate the patterns of lensing magnifications in the central region of planetary lens systems with various planet/star mass ratios and planet-star separations.  From the investigation of the magnification pattern, we set the criterion for planet detections and determine the lensing zone based on this criterion. We discuss about the implication of the determined lensing zone.  We summarize the result and conclude in \\S\\ 4. ",
        "conclusions": "While the convention of defining the range of planet separation detectable by the microlensing method is based on the location of the planet-induced caustic, the range for planets detectable through the channel of high-magnification events cannot be defined by the conventional criterion because the caustic is always located close to the primary lens regardless of the planet separation. Instead, the most important restriction to the detection of central perturbation is given by the finite-source effect. We investigated the magnification pattern in the central region of various planetary systems and found that perturbations with $\\geq 5\\%$ can be detected if the caustic is bigger than roughly one fourth of the caustic size.  Based on this finding, we derived an analytic expression for the range of planetary separations detectable through the central perturbations. From the comparison of classical lensing zone, we found that the new lensing zone is equivalent to the classical zone for a planet with a planet/star mass ratio $q\\sim 3\\times 10^{-4}$ and it becomes wider for heavier planets.  The wider lensing zone of the central perturbations for giant planets implies that the microlensing method provides an important tool to detect planetary systems composed of multiple ice-giant planets."
    },
    "0808/0808.4128_arXiv.txt": {
        "abstract": "In view of the imminent start of the LHC experimental programme, we use the available indirect experimental and cosmological information to estimate the likely range of parameters of the constrained minimal supersymmetric extension of the Standard Model (CMSSM), using a Markov-chain Monte Carlo (MCMC) technique to sample the parameter space. The 95\\% confidence-level area in the $(m_0, m_{1/2})$ plane of the CMSSM lies largely within the region that could be explored with 1~fb$^{-1}$ of integrated luminosity at 14~TeV, and much of the 68\\% confidence-level area  lies within the region that could be explored with 50~pb$^{-1}$ of integrated luminosity at 10~TeV. A same-sign dilepton signal could well be visible in most of the 68\\% confidence-level area with 1~fb$^{-1}$ of integrated luminosity at 14~TeV. We discuss the sensitivities of the preferred ranges to variations in the most relevant indirect experimental and cosmological constraints and also to deviations from the universality of the supersymmetry-breaking contributions to the masses of the Higgs bosons.  \\bigskip \\begin{center} CERN-PH-TH/2008-181, DCPT/08/118, FTPI-MINN-08/33, IPPP/08/59, UMN-TH-2715/08 \\end{center} \\vspace{-0.5cm} ",
        "introduction": "\\label{sec:intro}  Supersymmetry (SUSY)~\\cite{Nilles:1983ge,Haber:1984rc,Barbieri:1982eh} is one of the most highly favoured extensions of the Standard Model (SM), and is often considered to be a prime candidate for discovery at the Large Hadron Collider (LHC). With the start of experiments at the LHC now becoming imminent, it is natural and topical to make the best possible assessment of the likelihood that the LHC will indeed discover SUSY, based on the best available experimental, phenomenological and cosmological information. Most of the current constraints on possible physics beyond the SM are negative, in the sense that they reflect the agreement of data with the SM, and set only lower limits on the possible masses of supersymmetric particles~\\cite{pdg}. Examples are direct constraints such as lower limits on specific sparticles, e.g., the chargino, and indirect constraints such as the lower limit on the possible mass of a SM-like Higgs boson~\\cite{Barate:2003sz,Schael:2006cr}.  However, there are two observational constraints that, within the context of SUSY, may be used also to set upper limits on the possible masses of supersymmetric particles, since they correspond to measurements that cannot be explained by the SM alone. These hints for new physics are the anomalous magnetic moment of the muon, \\gmt, which appears to differ by over three standard deviations from the best SM calculation based on low-energy $e^+ e^-$ data~\\cite{Bennett:2006fi,Moroi:1995yh,Czarnecki:2001pv,Davier:2007ua,Miller:2007kk,Jegerlehner:2007xe,Passera:2008jk}, and the density of cold dark matter, $\\Omega_{\\rm CDM}$~\\cite{Dunkley:2008ie}, which has no possible origin within the SM. Each of these discrepancies has many possible interpretations, of which SUSY is just one. Nevertheless, given the strong motivations for SUSY, which include the naturalness of the mass hierarchy and grand unification, as well as the existence of a plausible candidate for the astrophysical cold dark matter, it is natural to ask what \\gmt\\ and $\\Omega_{\\rm CDM}$ may imply for the parameters of supersymmetric models. Any such analysis should also take into account the constraints imposed by precision measurements of electroweak observables (EWPO) and $B$-physics observables (BPO) such as \\bsg, where most observables agree quite well with the SM.  In this paper we revisit the indirect information on supersymmetric model parameters obtainable in the light of these experimental, phenomenological and cosmological constraints, using a Markov-chain Monte Carlo (MCMC) approach, see, e.g., Ref.~\\cite{Allanach:2007qj} and references therein. This is practical only in simplified versions of the minimal supersymmetric extension of the Standard Model (MSSM), in which some universality relations are imposed on the soft SUSY-breaking parameters. Initially, we work in the framework of the constrained MSSM (CMSSM), in which the scalar and gaugino mass parameters, $m_0$, $m_{1/2}$, and the trilinear coupling $A_0$ are each assumed to be equal at the input GUT scale. Furthermore as low-energy parameter we have $\\tan\\beta$, the ratio of the two vacuum expectation values of the two Higgs doublets. At the end we also comment on the possible changes in our results if the common soft SUSY-breaking contribution to the Higgs scalar masses-squared, $m_H^2$, is allowed to differ from those of the squarks and sleptons, the single-parameter non-universal Higgs model or NUHM1.  There have been many previous studies of the CMSSM parameter space~\\cite{Djouadi:2001yk,deBoer:2001nu,deBoer:2003xm,eos2,baer,arndut,eoss,baer2,lahnan,nath,munoz,Belanger:2004ag,Ellis:2003si,Ellis:2004tc,Ellis:2005tu,Ellis:2006ix,Ellis:2007aa,Ellis:2007ss,Baltz:2004aw,Allanach:2005kz,Allanach:2006jc,Allanach:2006cc,Allanach:2007qk,Feroz:2008wr,deAustri:2006pe,Roszkowski:2006mi,Roszkowski:2007fd,Roszkowski:2007va,Heinemeyer:2008fb,Bechtle:2004pc,Lafaye:2007vs,Buchmueller:2007zk,Allanach:2008iq}, including estimates of the sparticle masses, and a number of these have used MCMC techniques~\\cite{Baltz:2004aw,Allanach:2005kz,Allanach:2006jc,Allanach:2006cc,Allanach:2007qk,Feroz:2008wr,deAustri:2006pe,Roszkowski:2006mi,Roszkowski:2007fd,Roszkowski:2007va}. These have been used to extract the preferred values for the CMSSM parameters using low-energy precision data, bounds from astrophysical observables and flavour-related observables. These analyses differ in the precision observables that have been considered, the level of sophistication of the theory predictions that have been used, and the way the statistical analysis has been performed.  Here we use the MCMC technique to sample efficiently the SUSY parameter space, and thereby construct the $\\chi^2$ probability function, $P(\\chi^2,N_{\\rm dof})$. This accounts correctly for the number of degrees of freedom, $N_{\\rm dof}$, and thus represents a quantitative measure for the quality-of-fit.  Hence $P(\\chi^2,N_{\\rm dof})$ can be used to estimate the absolute probability with which the CMSSM describes the experimental data. Our probabilistic treatment is explained in detail in Sec.~\\ref{sec:mpfit}.   Many previous analyses found evidence for a relatively low SUSY mass scale in  the stau-coannihilation region, e.g.,~\\cite{Ellis:2007aa,Buchmueller:2007zk,trotta}, and a mild preference for  $\\tan\\beta \\sim 10$ was found in~\\cite{Buchmueller:2007zk}. A comparison of Bayesian analyses yielding varying results under different assumptions was made in~\\cite{Allanach:2008iq}. Some differences between analyses may also be traceable to the treatments of the \\bsg\\ and \\gmt\\ measurements, which we discuss in some detail below.   Our main objectives in this paper are threefold. One is to discuss explicitly the prospects for discovering sparticles in early LHC running, another is to discuss the robustness of the fit results by analyzing the implications of relaxing the constraints due to \\gmt, \\bsg, $\\Omega_{\\rm CDM}$ and other observables, and the third is to discuss the extension of the CMSSM results to the NUHM1, in which an extra parameter is introduced that allows a common degree of non-universality for the two Higgs multiplets.  We find that the 95\\%~C.L.\\ area in the $(m_{1/2}, m_0)$ plane of the CMSSM lies largely within the region that could be explored with 1~fb$^{-1}$ of integrated LHC luminosity at 14~TeV in a single experiment, and that much of the 68\\%~C.L.\\ area  lies within the region that could be explored with 50~pb$^{-1}$ of integrated luminosity at 10~TeV (the projected initial LHC collision energy). A same-sign dilepton signal could well be visible in the 68\\%~C.L.\\ area with 1~fb$^{-1}$ of integrated luminosity at 14~TeV, and the lightest Higgs boson might also be detectable in squark decays with 2~fb$^{-1}$ of integrated luminosity at 14~TeV. We find that removing the $\\Omega_{\\rm CDM}$ constraint has little effect on the preferred regions of the CMSSM parameter space in the $(m_0, m_{1/2})$, $(\\tan \\beta, m_{1/2})$, and $(\\tan \\beta, m_0)$ planes, apart from expanding the range of $m_0$, particularly for $\\tan \\beta \\sim 10$. On the other hand, rescaling the present error in \\gmt\\ may have quite an important effect: the preferred ranges in $m_{1/2}$ and $m_0$ would expand quite significantly if the error on the present experimental discrepancy with the SM were to be increased. Conversely, if this error could be reduced, e.g., by a more precise measurement of \\gmt\\ and/or a more refined theoretical estimate within the SM, the predictions for sparticle masses could be significantly improved. We also discuss the effects of possible variations in the errors in \\bsg\\ and other observables. Finally, we show that our results would not be greatly changed in the NUHM1: we leave a more complete study of the NUHM1 and the NUHM2 (in which the masses of the two Higgs multiplets are independently non-universal) for future work. ",
        "conclusions": "Making a probabilistic analysis using a MCMC technique, we have presented in this paper the regions preferred in the CMSSM and the NUHM1 parameter spaces at the 68\\% and 95\\% C.L., as well as the spectra at the best-fit points, in the light of the present direct and indirect constraints on the models' parameters. Particularly important roles are played by \\gmt\\ and \\bsg, and we have analyzed the ways in which effects of these constraints vary with the sizes of their theoretical uncertainties and experimental errors. We have quantified how strengthening (or relaxing) either of these constraints would reduce (or expand) considerably the preferred regions in the CMSSM $(m_0, m_{1/2})$ plane. We have also studied the impact of the constraint on the cold dark matter density imposed by WMAP. We find that the results for the best-fit points are remarkably robust with respect to imposing or dropping the constraint on the cold dark matter density. Perhaps surprisingly, we find that this constraint does not restrict significantly most two-dimensional projections of the preferred region in the CMSSM parameter space. Encouragingly, we find that the preferred regions in the NUHM1 are quite similar to those in the CMSSM.  The 95\\% exclusion regions in the $(m_0, m_{1/2})$ plane extend significantly further than the discovery regions shown above. Therefore, if SUSY were to be excluded at the LHC with 1~fb$^{-1}$ (100~pb$^{-1}$) of integrated luminosity at 14~TeV, the 95\\% (68\\%) C.L.\\ regions in both the CMSSM and the NUHM1 would be ruled out. On the other hand, SUSY could be discovered at the 5~$\\sigma$ level at the LHC with 1~fb$^{-1}$ of integrated luminosity at 14~TeV in a single experiment over most of the 95\\% C.L.\\ regions in the $(m_0, m_{1/2})$ planes of the CMSSM and the NUHM1. Only the highest $m_{1/2}$ values would require a larger integrated luminosity, or the combination of data from both ATLAS and CMS. Indeed, SUSY could be discovered over all of the 68\\% C.L.\\ regions in both the CMSSM and the NUHM1 with just 100~pb$^{-1}$ of integrated luminosity at 14~TeV, and even 50~pb$^{-1}$ of (good-quality) data at 10~TeV would offer significant prospects for SUSY detection. The same-sign dilepton search would cover most (all) of the 68\\% C.L.\\ region in the CMSSM (NUHM1). If Nature were to choose the best-fit CMSSM point, a measurement of the same-sign dilepton endpoint would impose a strong constraint on the SUSY spectrum.  One way or the other, there are good prospects that the initial runs of the LHC will determine the fate of many speculations about the relevance of low-energy SUSY to particle physics."
    },
    "0808/0808.2318_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.3in}}  Dark matter may be hidden, with no standard model gauge interactions. At the same time, in WIMPless models with hidden matter masses proportional to hidden gauge couplings squared, the hidden dark matter's thermal relic density may naturally be in the right range, preserving the key quantitative virtue of WIMPs. We consider this possibility in detail. We first determine model-independent constraints on hidden sectors from Big Bang nucleosynthesis and the cosmic microwave background. Contrary to conventional wisdom, large hidden sectors are easily accommodated. A flavour-free version of the standard model is allowed if the hidden sector is just 30\\% colder than the observable sector after reheating. Alternatively, if the hidden sector contains a 1-generation version of the standard model with characteristic mass scale below 1 MeV, even identical reheating temperatures are allowed. We then analyze hidden sector freezeout in detail for a concrete model, solving the Boltzmann equation numerically and understanding the results from both observable and hidden sector points of view. We find that WIMPless dark matter indeed obtains the correct relic density for masses in the range $\\kev \\lesssim m_X \\lesssim \\tev$. The upper bound results from the requirement of perturbativity, and the lower bound assumes that the observable and hidden sectors reheat to the same temperature and is raised to the MeV scale if the hidden sector is 10 times colder. WIMPless dark matter therefore generalizes the WIMP paradigm to the largest mass range possible for viable thermal relics and provides a unified framework for exploring dark matter signals across nine orders of magnitude in dark matter mass. ",
        "introduction": "\\label{sec:intro}  At present, all solid evidence for dark matter is gravitational.  At the same time, the possibility that dark matter has electromagnetic or strong interactions is highly constrained~\\cite{Dimopoulos:1989hk,Starkman:1990nj}.  A straightforward possibility, then, is that dark matter is hidden, consisting of particles that have no standard model (SM) gauge interactions.  Hidden sectors have a long and distinguished history.  For example, the idea that a hidden sector may restore parity on cosmological scales is as old as the idea of parity violation itself~\\cite{Lee:1956qn}.  Such hidden sectors, containing ``mirror matter,'' have been studied by many groups, as have their possible implications for dark matter~\\cite{Blinnikov:1982eh,Hodges:1993yb,% Berezhiani:1995am,Mohapatra:2000qx,Berezhiani:2000gw,Mohapatra:2001sx,% Ignatiev:2003js,Foot:2003iv,Berezhiani:2003wj,Foot:2004wz}.  String theory also motivates hidden sectors with their own gauge interactions~\\cite{Green:1984sg}, as in the case of the heterotic string~\\cite{Gross:1984dd} and intersecting brane models~\\cite{Blumenhagen:2005mu}.  More recently, hidden sectors have been central to several phenomenological developments, including Higgs portals~\\cite{Schabinger:2005ei,Patt:2006fw}, hidden valleys~\\cite{Strassler:2006im}, unparticles~\\cite{Georgi:2007ek}, and quirks~\\cite{Kang:2008ea}, and the possibility that dark matter may reside in a hidden sector has been discussed in several recent works~\\cite{Chen:2006ni,Kikuchi:2007az,MarchRussell:2008yu,% Hooper:2008im,McDonald:2008up,Kim:2008pp,Krolikowski:2008qa,% Gong:2008uz,Feng:2008ya,Foot:2008nw}.  In considering hidden dark matter, one would seemingly be sacrificing critical virtues of more conventional dark matter candidates, namely, the connection of dark matter to the gauge hierarchy problem, and the fact that weakly-interacting massive particles (WIMPs) naturally have the desired thermal relic density.  This is not necessarily true, however.  In the recently proposed framework of WIMPless dark matter~\\cite{Feng:2008ya}, hidden sector dark matter also naturally has the desired thermal relic density.  The essential idea is that, very generally, a dark matter candidate that decoupled from the thermal bath non-relativistically has a thermal relic density \\begin{equation} \\Omega h^2 \\sim \\frac{1}{\\langle \\sigma_A v \\rangle } \\sim \\frac{m^2}{g^4} \\ , \\label{omega} \\end{equation} where $\\langle \\sigma_A v \\rangle$ is the thermally-averaged annihilation cross section times velocity, and $g$ and $m$ are the dark matter particle's gauge coupling and mass.  For WIMPs with weak interaction coupling constant $\\gweak \\simeq 0.65$ and weak-scale mass $\\mweak \\sim 100~\\gev - 1~\\tev$, this relic mass density is naturally near the observed $\\Omega_{\\text{DM}} h^2 \\simeq 0.11$.  In WIMPless models, hidden sector particles $X$ have masses and couplings that may be very different from WIMPs, but which nevertheless satisfy \\begin{equation} \\frac{m_X}{g_X^2} \\sim \\frac{\\mweak}{\\gweak^2} \\ . \\label{scaling} \\end{equation} In the WIMPless examples discussed in Ref.~\\cite{Feng:2008ya}, \\eqref{scaling} follows from the structure of gauge-mediated supersymmetry breaking (GMSB), which generates hidden sector masses that are proportional to couplings squared.  The WIMPless framework therefore naturally generalizes the WIMP paradigm to other cold dark matter candidates without sacrificing the key thermal relic density virtue of WIMPs.  If additional connector particles with both SM and hidden gauge quantum numbers are present, they will mediate SM-hidden interactions.  The WIMPless framework may therefore also predict qualitatively new signals, such as an explanation of the DAMA signal in terms of GeV dark matter~\\cite{Feng:2008dz}.  In this paper, we study hidden thermal relics in general, and WIMPless dark matter in particular.  For the most part, we assume that there are no connector particles, so that the hidden sector interacts with the SM only very weakly.  In \\secref{BBN}, we carry out a model-independent analysis of constraints on hidden sectors from Big Bang nucleosynthesis (BBN), the cosmic microwave background (CMB), and other cosmological data.  As is well known, these exclude the possibility that the hidden sector is an exact copy of the SM~\\cite{Kolb:1985bf}.  This statement assumes that the hidden and observable sectors are at the same temperature at late times, however. In the present context, there is no strong reason to assume that the observable and hidden sectors are reheated to the same temperature~\\cite{Hodges:1993yb,Berezhiani:1995am}, and, even if they are, for a truly decoupled hidden sector, their temperatures will differ by the time of BBN.  Following previous works in the context of mirror matter (see, \\eg, Refs.~\\cite{Hodges:1993yb,Berezhiani:1995am,% Foot:1996hp,Berezhiani:2000gw,Ignatiev:2003js}), we consider the cosmological constraints allowing different temperatures and find that even large hidden sectors are easily accommodated. For example, a hidden version of the SM is allowed even if the hidden sector's reheating temperature is as much as 70\\% of the observable sector's. Alternatively, if the hidden sector contains a 1-generation version of the MSSM with characteristic mass scale below 1 MeV, even identical reheating temperatures are allowed.  We then present a concrete model for the hidden sector in \\secref{concrete} and study the freezeout of dark matter in this hidden sector in \\secref{relic}.  The relations of \\eqsref{omega}{scaling} are, of course, only rough relations capturing the parametric dependences.  In these Sections, we present a detailed analysis of a specific case that highlights several subtleties and generalizes previous freezeout analyses to cases with sectors at different temperatures.  This analysis confirms the validity of conclusions based on \\eqsref{omega}{scaling} and highlights various subtleties of the WIMPless framework.  In particular, we find that WIMPless dark matter may have the desired thermal relic density for the entire range of $\\kev \\alt m_X \\alt \\tev$, where the lower bound is set by the requirement that dark matter freezeout is non-relativistic, and the upper bound follows from the requirement of perturbative gauge couplings. WIMPless dark matter therefore encompasses as large a range of masses as one could expect of dark matter that has the naturally correct thermal relic density, and it provides a unified framework for addressing many diverse dark matter signals and phenomenology.  Finally, in \\secref{adding}, we discuss modifications from the presence of connector particles with both hidden and observable sector gauge interactions.  We find that, under general assumptions, these fields do not upset the conclusions derived earlier.  We present our conclusions in \\secref{conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In the paper, we have studied the possibility that dark matter is hidden, that is, has no SM gauge interactions.  Hidden sectors appear in many frameworks for physics beyond the SM, and hidden dark matter is perfectly viable, given all observations to date.  We are particularly motivated by the recent proposal of WIMPless dark matter. As with WIMPs, WIMPless dark matter has the key virtue that its thermal relic density is in the right range to be cold dark matter. In contrast to WIMPs, however, WIMPless dark matter may have masses and couplings that differ drastically from WIMPs.  Although hidden sectors interact very weakly with the SM, hidden particles contribute to the energy density of the Universe.  This impacts the expansion rate of the Universe, and so bounds from BBN and the CMB constrain hidden sectors.  As is well known, current constraints completely exclude a hidden sector that is an exact copy of the MSSM.  We find, however, that these constraints are rather brittle and easily avoided if the hidden sector is just slightly colder than the observable sector or if the mass scale of hidden superpartners differs from the MSSM.  Several viable examples are discussed in \\secref{BBN}, and \\figsref{gstar}{gstarcmb} give model-independent bounds that can be used to determine the viability of other models.  In \\secref{concrete}, we then defined a concrete model, a 1-generation flavor-free version of the MSSM.  This model has features generic to GMSB models and provides a WIMPless dark matter candidate, the hidden stau $\\stau_R^h$.  The $\\stau_R^h$ relic density depends sensitively on only a small number of parameters.  This relic density was examined in \\secref{relic} by solving the Boltzmann equation numerically, and understanding these results through well-known analytic approximations, generalized to hidden sectors with different temperatures.  These results confirm that in WIMPless models with $m_X \\propto g_X^2$, the relic density is naturally in the right range for dark matter masses $\\kev \\alt m_X \\alt \\tev$, greatly extending the conventional WIMP mass range.  WIMPless dark matter therefore provides a class of dark matter candidates that shares the key relic density virtue of WIMPs.  This generalization may have interesting applications, given the diverse and interesting phenomenology and observational anomalies already present at the keV, MeV, GeV, and TeV mass scales.  We have determined the allowed range of $m_X$ by requiring that this dark matter freezes out with the right relic density while non-relativistic.  Of course, this range may be further constrained by other considerations.  In particular, the lightest WIMPless candidates, with $m_X \\sim \\kev$, may have interesting implications for structure formation.  After freeze out, they can still couple to the hidden sector thermal bath through elastic scattering processes $\\stau_R^h \\gamma^h \\to \\stau_R^h \\gamma^h$ and $\\stau_R^h \\nu^h \\to \\stau_R^h \\nu^h$.  After their kinetic decoupling at temperature $T_{\\text{kd}}$, they then stream freely between over- and under-dense regions.  Free-streaming of thermal relics causes damping in density perturbations~\\cite{Bertschinger:2006nq} at a comoving scale $\\lambda_{\\text{FS}} \\propto m^{-1/2}_X\\, T_{\\text{kd}}^{-1/2}$; in contrast to WIMPs~\\cite{Loeb:2005pm,Profumo:2006bv}, this is the dominant effect for MeV or lighter dark matter~\\cite{Hooper:2007tu}. If this scale falls in the range $(1-80)\\, h^{-1}$ Mpc, the linear matter power spectrum from high-resolution Lyman-$\\alpha$ forest data can be used to place a lower bound on the mass of the dark matter particle. Current lower mass limits on warm dark matter of $0.55 - {\\cal O}(1)~\\kev$ from Lyman-$\\alpha$ data~\\cite{Narayanan:2000tp,Viel:2005qj,Abazajian:2005xn,Viel:2007mv} therefore imply lower mass bounds in the WIMPless scenario, which depend on the kinetic decoupling temperature $T_{\\text{kd}}$ for different $\\xiRH$.  In addition, another kind of dark matter lower mass limit follows from the observation that the phase-space density of the dark matter particles in galaxy halos cannot exceed the maximum value of the phase-space density when the dark matter particles were in kinetic equilibrium~\\cite{Tremaine:1979we}. With knowledge of the WIMPless dark matter's chemical potential, one can then apply these generalized Tremaine-Gunn bounds from Ref.~\\cite{Madsen:1991mz} to our scenario. Detailed considerations of these effects are beyond the scope of this study, but they may modify the lower bound on $m_X$ or, alternatively, provide interesting astrophysical signals for the hidden dark matter scenarios discussed here."
    },
    "0808/0808.0637_arXiv.txt": {
        "abstract": "The radio emission in radio loud quasars originates in a jet carrying relativistic electrons. In radio quiet quasars (RQQs) the relative radio emission is $\\sim 10^3$ times weaker, and its origin is not established yet. We show here that there is a strong correlation between the radio luminosity ($L_R$) and X-ray luminosity ($L_X$) with $L_R\\sim 10^{-5}L_X$, for the radio quiet Palomar-Green (PG) quasar sample. The sample is optically selected, with nearly complete radio and X-ray detections, and thus this correlation cannot be due to direct selection biases. The PG quasars lie on an extension of a similar correlation noted by Panessa et al., for a small sample of nearby low luminosity type 1 AGN. A remarkably similar correlation, known as the G{\\\"u}del-Benz relation, where $L_R/L_X\\sim 10^{-5}$, holds for coronally active stars. The G{\\\"u}del-Benz relation, together with correlated stellar X-ray and radio variability, implies that the coronae are magnetically heated. We therefore raise the possibility that AGN coronae are also magnetically heated, and that the radio emission in RQQ also originates in coronal activity. If correct, then RQQ should generally display compact flat cores at a few GHz due to synchrotron self-absorption, while at a few hundred GHz we should be able to see directly the X-ray emitting corona, and relatively rapid and large amplitude variability, correlated with the X-ray variability, is likely to be seen. We also discuss possible evidence that the radio and X-ray emission in ultra luminous X-ray sources and Galactic black holes may be of coronal origin as well. ",
        "introduction": "AGN emit continuum radiation from the radio to the hard X-rays, and in some objects beyond. The overall spectral energy distribution (SED) of AGN has a characterstic shape, with a relatively small dispersion (e.g. Sanders et al. 1989, Elvis et al. 1994), except in the radio band, where the relative strength of the radio emission spans a range of $>10^6$, from the most radio loud AGN to the most radio quiet AGN, where the radio emission is undetectable. Furthermore, there is evidence that the distributions of absolute radio power (Miller et al. 1990) and relative radio power, commonly measured using $R\\equiv f_{\\rm 6~cm}/f_{\\rm 4400~\\AA}$, are bimodal (Kellerman et al. 1989), with radio loud (RL) AGN having $R\\sim 10^2-10^5$ and radio quiet (RQ) AGN having $R\\la 0.1-10$, though the strength of the bi-modality may be weaker than initially estimated (White et al. 2007). The radio emission of RL AGN originates in a fast jet carrying relativistic electrons, as clearly established through high resolution radio imaging (e.g. Begelman et al. 1984, and citations thereafter), however the origin of the radio emission in RQ AGN is not established yet.  Early radio imaging of RQ AGN confined the radio emission to arcsec scale, which allowed the option of supernovae related radio emission from compact starburst activity (Wilson \\& Willis 1980; Ulvestad \\& Wilson 1984; Sopp \\& Alexander 1991; Condon et al. 1991; Terlevich et al. 1992). However, sub arcsecond imaging of RQ AGN (Kellerman et al. 1994; Kukula et al. 1995, 1998; Leipski et al. 2006), followed by higher resolution VLBI imaging (Blundell et al. 1996, Blundell \\& Beasley 1998; Ulvestad \\& Ho 2001; Nagar et al. 2002; Anderson et al. 2004; Ulvestad et al. 2005), found significant compact radio emission on mas, i.e. pc scale. Most recently, radio variability monitoring of RQ AGN confine the radio emission to $\\la 0.1$~pc in luminous AGN (Barvainis et al. 2005), and $\\la 10^{-4} - 10^{-2}$~pc in low luminosity AGN (Anderson \\& Ulvestad 2005), which clearly indicates a compact non thermal source. Unlike high resolution imaging of RL AGN, which often resolves a significant fraction of the total radio flux already on arcsec (i.e. kpc) scale, in RQ AGN the emission is mostly unresolved at arcsec scale, and often remains largely unresolved down to mas. {\\em What produces this radio emission?}. A plausible explanation is a scaled down version of the RL AGN mechanism, i.e. a low power jet (Miller et al. 1993; Falcke et al. 1996), which dissipates before leaving the core, explaining the general lack of resolved emission. However, the bimodal distribution in jet power remains to be understood.  Further hints on the origin of the radio emission in RQ AGN can be provided by correlating it with other emission properties. In this paper we focus on the correlation of the radio emission with the X-ray emission (herafter the R-X relation). Earlier studies of this relation (Brinkmann et al. 2000; Salvato et al. 2004; Wang et al. 2006), indicated a roughly linear relation for luminous RQ AGN. These studies were based on a correlation of the ROSAT all sky survey, and large area radio surveys (VLA FIRST and NVSS). Most objects in flux limited surveys tend to cluster close to the flux limit, and thus a linear relation would appear if the range in luminosities is significantly larger than the range in fluxes, as the luminosity in both bands is then necessarily correlated with the redshift. More recently, Panessa et al. (2007) found a linear R-X relation in an optically selected sample of nearby low luminosity AGN (see also Terashima \\& Wilson 2003). This sample is still somewhat affected by selection effects and incompleteness, and a significant fraction of the objects there may be affected by X-ray absorption, which cannot be accurately corrected.  To avoid a correlation induced by selection effects, one needs a well defined and complete sample of AGN, which is selected independently of the radio and X-ray emission properties, and yet includes complete radio and X-ray detections. The best sample for this purpose is PG Bright Quasar Survey sample (Schmidt \\& Green 1983). This sample is selected based on flux ($B<16.16$), morphology (point like on the Palomar Schmidt Survey plates), and color ($U-B<-0.44$). It includes the 114 brightest AGN in $1/4$ of the sky (at high celestial and Galactic latitudes). The completeness of the sample has been debated, and the recent thorough study of Jester et al. (2005) finds a somewhat different effective cut ($U-B<-0.7$), but with no systematic incompleteness (at $z<0.5$). The PG quasars are representative of $B$ band selected quasars, and are generally similar to bright quasars selected in other wavelengths. The bright magnitude limit facilitates studies of this sample in many other bands, including the X-ray and radio bands. An unpublished study (Ian George, private communication), based on {\\em ASCA} observations of 26 PG quasars (George et al. 2000), indicated a strong R-X relation. Motivated by this result, and the presence of a similarly strong R-X relation in coronally active stars, we study here the R-X relation in a complete sample of 87 $z\\le 0.5$ PG quasars (Boroson \\& Green 1992). We find that both coronally active stars and active galaxies appear to follow the same relation, as further discussed below. In section 2 we present the R-X relation for the PG quasars, and expand the luminosity range to include a smaller sample of low luminosity AGN, ultraluminous X-ray sources, Galactic Black Holes (GBHs), and coronally active stars. In section 3 we discuss various implications for the origin of the radio emission in RQ AGN, and possible future tests of the suggested origin.  \\begin{figure*} \\includegraphics[width=165mm,angle=0]{fig1.ps} \\caption{The distribution of the PG quasars in the $L_{\\rm R}$ vs. $L_{\\rm X}$ plane. RL AGN are designated by stars, and their $L_{\\rm R}/L_{\\rm X}$ is about a factor of $10^3$ higher than for RQ AGN. Circles mark objects with \\C\\ absorption EW$>1$\\AA\\ (some of which are RL), where left attached arrows mark upper limits to $L_{\\rm X}$. These UV absorbed AGN tend to be underluminous by a factor  of 10-30 in $L_{\\rm X}$ for a given $L_{\\rm R}$ (and also for a given optical luminosity), most likely due to absorption. Squares are RQ AGN without \\C\\ absorption, where the filled/empty squares mark objects with radio detections/upper limits. Note the small scatter in the $L_{\\rm R}$ vs. $L_{\\rm X}$ correlation in {\\em optically} selected AGN, once RL and UV absorbed AGN are excluded.} \\end{figure*} ",
        "conclusions": "We find that RQ AGN lie on the G{\\\"u}del-Benz relation, $L_{\\rm R}/L_{\\rm X}=10^{-5}$, found for coronally active stars. Since the X-ray emission in AGN most likely originates from a hot corona, and the corona is likely to be magnetically heated, it is natural to associate the radio emission in RQ AGN with a coronal origin as well.  The ``coronal paradigm\" for the radio emission implies a very compact source, which is synchrotron self-absorbed at the GHz range. This implies that a compact flat spectrum source should generally be present in RQ AGN. Self-absorption will become negligible at mm wavelengths, and this emission should originate from the same volume producing the X-ray emission. The mm emission is likely to be strongly variable, and it should be correlated with the X-ray variability, possibly as seen in the Neupert effect of stellar coronal flares.  Despite the factor of $\\sim 10^{10}-10^{15}$ difference in $L_{\\rm R}$ and $L_{\\rm X}$ between AGN and coronally active stars, the estimated values of $B$ in the active regions of the two systems are comparable ($\\sim 10^3-10^4$~Gauss), which may imply similar underlying microphysics in both coronae. The physical mechanism underlying the $L_{\\rm R}/L_{\\rm X}=10^{-5}$ ratio in these coronae remains a puzzle.  Applying the stellar analogy to the extended radio emission, we suggest that it originates from coronal mass ejections. This emission mechanism differrs from the alternative jet interpretation in that the outflow is relatively slow, and is not well confined, as suggested in observations of some nearby RQ Seyfert galaxies.  There are two additional populations of active radio and X-ray emitting objects, intermediate between stars and AGN. ULXs, where the four objects with radio and X-ray data also fall close to the G{\\\"u}del-Benz relation, and overlap with the position of NGC~4395, the lowest luminosity type 1 AGN. This is consistent with ULXs being scaled down AGN, rather than strongly beamed GBHs. The other population is GBHs, where $L_{\\rm R}$ is a factor of 10-100 lower than expected from the G{\\\"u}del-Benz relation. This offset may result from a dependence of $L_{\\rm R}/L_{\\rm X}$ on the local disk temperature. We show that the Merloni et al. ``fundamental plane\" relation, which incorporates both GBHs and AGN, can be recast as a dependence of $L_{\\rm R}/L_{\\rm X}$ on the surface accretion disk effective temperature. This may indicate that the radio emission in GBHs has a significant coronal component as well, in particular given the observed variability pattern, where $L_{\\rm o}\\propto -dL_{\\rm X}/dt$, which is reminiscent of the Neupert effect in stellar coronae.  We thank M. G{\\\"u}del for providing data in electronic form and A. Merloni for the fundamental-plane relation in his radio-quiet sub-sample. We thank R. Antonucci and S. Jester for valuable comments on the manuscript. We thank the referee for many useful comments that helped us improve the manuscript. This research was supported in part by the Asher Space Research Institute at the Technion."
    },
    "0808/0808.2542_arXiv.txt": {
        "abstract": "We present a CCD survey of variable stars in the Draco dwarf spheroidal galaxy.  This survey, which has the largest areal coverage since the original variable star survey by Baade \\& Swope, includes photometry for 270 RR Lyrae stars, 9 anomalous Cepheids, 2 eclipsing binaries, and 12 slow, irregular red variables, as well as 30 background QSOs.  Twenty-six probable double-mode RR Lyrae stars were identified. Observed parameters, including mean $V$ and $I$ magnitudes, $V$ amplitudes, and periods, have been derived. Photometric metallicities of the ab-type RR Lyrae stars were calculated according to the method of Jurcsik \\& Kovacs, yielding a mean metallicity of $\\langle [Fe/H] \\rangle = -2.19 \\pm 0.03$.  The well known Oosterhoff intermediate nature of the RR Lyrae stars in Draco is reconfirmed, although the double-mode RR Lyrae stars with one exception have properties similar to those found in Oosterhoff type II globular clusters.  The period-luminosity relation of the anomalous Cepheids is rediscussed with the addition of the new Draco anomalous Cepheids. ",
        "introduction": "The Draco dwarf spheroidal (dSph) galaxy ($\\alpha_{2000.0} = 17^{h} 20^{m} 12.39^{s} $, $\\delta_{2000.0}= +57^{\\circ} 54' 55.3''$), a satellite of the Milky Way Galaxy, was first extensively studied by \\citet{Baade:1961} (hereafter known as B\\&S). They reported discovering over 260 variable stars and obtained photometry for 138 variables in the central region of Draco, 133 of which were of RR Lyrae (RRL) type.  Several subsequent studies have investigated aspects of the variable star population in Draco. \\citet{Zinn:1976a} reported new observations of the anomalous Cepheids in Draco. \\citet{Nemec:1985a} reanalyzed the B\\&S photometry and produced updated periods for the B\\&S variables.  Both \\citet{Nemec:1985a} and \\citet{Goranskij:1982} reported new double-mode RRL in Draco.  Recently \\citet{Bonanos:2004} provided a photometric study of Draco which produced light curves for 146 RRL stars, four anomalous Cepheids, an SX Phe star, and a field eclipsing binary.  In this work, we use CCD observations to update the census of variable stars in Draco.  We cover an area slightly larger than the full B\\&S survey, and we discover new variables with smaller amplitudes than those found by B\\&S.  We provide photometric data, periods, and light curves for over 320 variable stars.  This paper is organized in the following manner: Section 2 describes our data acquisition and data reduction processes.  Section 3 covers our analysis techniques. Periods, light curves, and classifications of the variable stars are presented in Section 4.  A re-discussion of the Oosterhoff classification of the Draco dwarf spheroidal is presented in Section 5. Conclusions are summarized in section 6. ",
        "conclusions": "We have presented the latest census of variable stars of the Draco dwarf spheroidal galaxy.  We have found 81 new RRab stars, 8 new RRc stars, and 16 probable new RRd stars, thus bringing to 214 RRab, 30 RRc and 26 RRd the total number of RRL stars of the different types known in Draco. We have increased the number of anomalous Cepheids to nine from five.  Draco cannot be clearly classified as an Oo I or an Oo II type system.  Based upon the mean period of its RRab stars and their location in the period-amplitude diagram, Draco is Oosterhoff intermediate.  Objects exactly like Draco are thus not likely to be important building blocks in forming the Galactic halo. We note, however, that the properties of the RRd stars in Draco are, with a single exception, similar to those of RRd stars in Oo II clusters.  The anomalous Cepheids in the Draco dSph galaxy show a possible dual P-L relation stemming from the pulsational modes of the stars.  However, with so few stars populating the first-overtone relation, we cannot say with certainty that two P-L relation alternative is the only one capable of describing the relationships of luminosity and period for Draco AC stars. In addition to the pulsating variable stars, we find two field eclipsing binary stars, 30 background QSOs, and 12 long period variable stars."
    },
    "0808/0808.3772_arXiv.txt": {
        "abstract": "The Milky Way has at least twenty-three known satellite galaxies that shine with luminosities ranging from about a thousand to a billion times that of the Sun. Half of these galaxies were discovered ~\\cite{Willman2005,Belokurov2007} in the past few years in the Sloan Digital Sky Survey, and they are among the least luminous galaxies in the known Universe. A determination of the mass of these galaxies provides a test of galaxy formation at the smallest scales~\\cite{Mateo1998,Gilmore2007} and probes the nature of the dark matter that dominates the mass density of the Universe ~\\cite{Spergel2007}. Here we use new measurements of the velocities of the stars in these galaxies~\\cite{SimonandGeha,Walker2007} to show that they are consistent with them having a common mass of about $10^7$ M$_\\odot$ within their central 300 parsecs. This result demonstrates that the faintest of the Milky Way satellites are the most dark matter-dominated galaxies known, and could be a hint of a new scale in galaxy formation or a characteristic scale for the clustering of dark matter. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.1893_arXiv.txt": {
        "abstract": "Type 2 quasars are luminous Active Galactic Nuclei whose central engines are seen through large amounts of gas and dust. We present \\spi\\ spectra of twelve type 2 quasars selected on the basis of their optical emission line properties. Within this sample, we find a surprising diversity of spectra, from those that are featureless to those showing strong PAH emission, deep silicate absorption at 10\\micron, hydrocarbon absorption, high-ionization emission lines and H$_2$ rotational emission lines. About half of the objects in the sample are likely Compton-thick, including the two with the deepest Si absorption. The median star-formation luminosity of the objects in our sample measured from the strength of the PAH features is $5\\times 10^{11}L_{\\odot}$, much higher than for field galaxies or for any other AGN sample, but similar to other samples of type 2 quasars. This suggests an evolutionary link between obscured quasars and peak star formation activity in the host galaxy. Despite the high level of star formation, the bolometric output is dominated by the quasar in all cases. For a given strength of 10\\micron\\ Si absorption, ULIRGs are significantly colder than are type 2 quasars (their $F_{\\nu}[14.5\\micron]/F_{\\nu}[27.5\\micron]$ ratio is 0.5 dex lower), perhaps reflecting different obscuration geometries in these sources. We find that the appearance of the 10\\micron\\ feature (i.e., whether it shows in emission or in absorption) is well-correlated with the optical classification in type 1 and type 2 quasars, contrary to some models of clumpy obscuration. Furthermore, this correlation is significantly stronger in quasars ($L_{\\rm bol}\\ga 10^{45}$ erg/s) than it is in Seyfert galaxies ($L_{\\rm bol}\\ll 10^{45}$ erg/s). ",
        "introduction": "Type 2 quasars are powerful ($L_{\\rm bol}\\ga 10^{45}$ erg/s) active galactic nuclei (AGNs) whose central regions are shielded from the observer by large amounts of gas and dust. Until recently, the existence of dust-enshrouded quasars was in question, but now these objects are routinely discovered at X-ray (e.g., \\citealt{szok04, netz06}), infrared \\citep{lacy04, ster05, mart06} and optical \\citep{zaka03, reye08} wavelengths. A significant controversy remains regarding their space density and contribution to the AGN population luminosity budget. Some analyses suggest that there are at least as many obscured as unobscured quasars even at the highest luminosity end \\citep{reye08, mart06, poll08}, whereas others argue for a strong decline of obscured quasar fraction with luminosity \\citep{ueda03, maio07, trei08}.  If an AGN is entirely embedded in an optically thick dusty cloud, virtually all of its original radiation (emitted primarily at optical to soft X-ray wavelengths) is absorbed by dust and is then re-emitted thermally in the IR. In this case, there is no optical signature of the AGN, but its presence may be revealed when hard X-rays and a large IR luminosity are detected. In this way, some ultra-luminous infrared galaxies (ULIRGs) have been shown to contain embedded AGNs \\citep{risa00, fran03, teng05}. The dominant energy source in ULIRGs remains a matter of controversy. Some authors argue that they are powered by star formation, with the AGN contributing only a small fraction \\citep{fran03}, while others suggest that buried AGNs are energetically dominant in many ULIRGs \\citep{iman07}.  An alternative possibility is that the nucleus is only partially covered by the obscuring material. In this case, the AGN is seen as obscured or unobscured depending on the line of sight, making its UV, optical and X-ray emission highly anisotropic \\citep{anto93}. In particular, UV and X-ray emission can make its way out through the holes in the obscuring cloud and photo-ionize extended regions, which in turn produce powerful optical emission lines. Using the Sloan Digital Sky Survey (SDSS; \\citealt{york00, stou02, adel08}) we have selected hundreds of type 2 quasars at redshifts $z<0.8$ on the basis of their optical emission line properties \\citep{zaka03, reye08}.  IR observations are key to determining the physical conditions of the obscuring material and its geometric properties such as its covering factor and extent. In this paper we present \\spi\\ Space Telescope \\citep{wern04} Infrared Spectrograph (IRS; \\citealt{houc04}) observations of twelve type 2 quasars from the SDSS sample. In Section \\ref{sec_data} we describe our sample selection and data reduction. In Section \\ref{sec_spectra} we present a detailed analysis of the spectral features. We compare our sample with other samples of type 2 quasars in Section \\ref{sec_compare}, we discuss our results in the context of radiative transfer models in Section \\ref{sec_models} and we conclude in Section \\ref{sec_conc}.  Throughout the paper, we adopt a cosmology with $h=0.7$, $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$. Objects are identified as SDSS Jhhmmss.ss$+$ddmmss.s in Figure \\ref{pic1} and Table \\ref{tab_summary} and are shortened to SDSS Jhhmm$+$ddmm in the text. We refer to type 1 and type 2 AGNs based on their optical classification (in type 1 objects, broad lines are detected and in type 2, they are not). When discussing the luminosity dependence of AGN phenomenon, we distinguish between Seyfert galaxies ($L_{\\rm bol}\\ll10^{45}$ erg/s) and quasars ($L_{\\rm bol}\\ga 10^{45}$ erg/s) based on their bolometric luminosities. ",
        "conclusions": "\\label{sec_conc}  We have conducted \\spi\\ IRS spectroscopy of ten type 2 quasars and obtained spectra of two more objects from the \\spi\\ archives. The sample was selected from the SDSS spectroscopic database on the basis of optical emission line properties and was required to have a minimal [OIII]5007\\AA\\ luminosity of $10^9L_{\\odot}$ and a minimal IR flux to obtain good quality spectra, but was otherwise unbiased with regard to any multi-wavelength properties. In particular, the sample appears to represent a wide range of column densities along the line of sight, and half of the objects are likely to be Compton-thick. Optically-selected type 2 quasars show a diversity of IR spectral properties. Most objects show high-ionization fine structure emission lines such as [NeVI]7.63\\micron, essentially ruling out star formation as the dominant energy source. The 10\\micron\\ Si feature appears in absorption in eight of the objects with a large range of strengths, and in another three the feature is not detected. In the one remaining object (SDSS~J0920+4531) the feature may have been detected in emission, peaking at 11\\micron, which corresponds to an unusually low color temperature of $\\simeq 110$K.  Half of the sample show PAH features. One type 2 quasar (SDSS~J0815+4304) shows absorption due to organic carbonaceous grains and possibly water ice absorption. We detect three rotational lines of molecular hydrogen in one object (SDSS~J1641+3858), with excitation temperature $T_{\\rm exc}=700$K and mass $M=3.5\\times 10^7M_{\\odot}$. Molecular hydrogen is also detected (with a lower confidence) in SDSS~J0050$-$0039.  The median IR star-formation luminosity of type 2 quasar host galaxies inferred from PAH emission is $5\\times 10^{11}L_{\\odot}$, but the bolometric output of our objects is dominated by the obscured quasars. These levels of star formation are similar to those found by \\citet{lacy07} and \\citet{mart08} in other samples of type 2 quasars. The high star-formation luminosities of type 2 quasars -- much higher than in field galaxies -- suggest that quasar activity is associated with intensified star formation in its host galaxy, as suggested by merger-driven models of quasar evolution \\citep{hopk06}. Moreover, the median star-formation luminosity that we find is higher than that of type 1 quasar host galaxies, $6\\times 10^{10}-1.4\\times 10^{11}L_{\\sun}$, depending on the quasar luminosity \\citep{shi07}. This trend supports the suggestion that the obscured fraction evolves with time, in the sense that the probability to see the quasar as obscured decreases with time after the merger \\citep{hopk06}. In this case type 2 quasars would appear on average at an earlier stage of the merger and therefore with a higher star-formation rate than type 1 quasars, as observed.  Objects in our sample occupy luminous bulge-dominated or elliptical hosts, whereas IR-selected type 2 quasars from the sample of \\citet{lacy07} inhabit dusty disk galaxies. The two samples have different luminosities, with the objects from our sample being about 0.5 dex more luminous ($\\nu L_{\\nu}[27\\micron]=10^{45.3}$ erg/s in our sample vs. $10^{44.8}$ erg/s in the sample of \\citealt{lacy07}). \\citet{hopk08} argue that at $z\\ll 1$ most AGN activity at low luminosities is due to gas instabilities in the host galaxy, while high luminosity quasars are fueled in major mergers which result in formation of bulges. One possible explanation for the observed difference between host properties is that the two samples bracket a transition between different quasar fueling mechanisms.  Both optically-classified type 2 and type 1 quasars often show spectra with no sign of the 10\\micron\\ Si feature. The fraction of such objects depends on how the sample was selected; overall, about a third of all type 1 and type 2 quasars discussed in this paper do not show the feature. Most quasars that show Si in absorption are optically classified as type 2, while those with Si emission are type 1. Some models of clumpy AGN tori show that the appearance of the feature is primarily determined by the geometrical distribution of the material and its total optical depth, while the viewing angle acts as a secondary parameter. Therefore, the observed correlation between the optical classification and the appearance of the IR spectrum may be a useful discriminator of models of AGN obscuration. We further find that the correlation is significantly weaker in Seyfert galaxies than it is in quasars -- specifically, half of Seyfert 1 galaxies show the 10\\micron\\ Si feature in absorption. This suggests that either the geometry or the physical properties of obscuring material vary with luminosity."
    },
    "0808/0808.4152_arXiv.txt": {
        "abstract": "A class of modified gravity, known as $f(R)$-gravity, has presently been applied to Cosmology as a realistic alternative to dark energy. In this paper we use the most recent Type-Ia Supernova (SNe Ia) data, the so-called \\emph{Union} sample of 307 SNe Ia, to place bounds on a theory of the form $f(R)=R - \\beta/R^n$ within the Palatini approach.  Given the complementarity of SNe Ia data with other cosmological observables, a joint analysis with  measurements of baryon acoustic oscillation peak and estimates of the {\\rm CMB} shift parameter is also performed. We show that, for the allowed intervals of $n$, $\\Omega_{mo}$, and $\\beta$, models based on $f(R) = R - \\beta/R^{n}$ gravity in the Palatini approach can produce the sequence of radiation-dominated, matter-dominated, and accelerating periods without need of dark energy. ",
        "introduction": "One of the key problems at the interface between fundamental physics and cosmology is to understand the physical mechanism behind the late-time acceleration of the Universe. In principle, this phenomenon may be the result of unknown physical processes involving either modifications of gravitation theory or the existence of new fields in high energy physics. Although the latter route is most commonly used, which gives rise to the idea of a dark energy component (see, e.g., \\cite{alc}), following the former, at least two other attractive approaches to this problem can be explored. The first one is related to the possible existence of extra dimensions, an idea that links cosmic acceleration with the hierarchy problem in high energy physics, and gives rise to the so-called brane-world cosmology \\cite{RandSund1999}. The second one, known as $f(R)$ gravity, examine the possibility of modifying Einstein's general relativity (GR) by adding terms proportional to powers of the Ricci scalar $R$ to the Einstein-Hilbert Lagrangian \\cite{Buchdahl}. The cosmological interest in $f(R)$ gravity comes from the fact that these theories can exhibit naturally an accelerating expansion without introducing dark energy.  However, the freedom in the choice of different functional forms of $f(R)$ gives rise to the problem of how to constrain on theoretical and/or observational grounds, the many possible $f(R)$ gravity theories. Much efforts within the realm, mainly from a theoretical viewpoint, have been developed so far~\\cite{Amendola} (see also Refs.~\\cite{Francaviglia} for recent reviews), while only recently observational constraints from several cosmological data sets have been explored for testing the viability of these theories\\cite{Amarzguioui,Tavakol,Fairbairn,Movahed,Fabiocc,tv}.  An important aspect worth emphasizing concerns the two different variational approaches that may be followed when one works with $f(R)$ gravity theories, namely, the metric  and the Palatini formalisms (see, e.g., \\cite{Francaviglia}). In the metric formalism the connections are assumed to be the Christoffel symbols and variation of the action is taken with respect to the metric, whereas in the Palatini variational approach the metric and the affine connections are treated as independent fields and the variation is taken with respect to both. In fact, these approaches are equivalents only in the context of GR, that is, in the case of linear Hilbert action; for a general $f(R)$ term in the action they give different equations of motion.  In the present paper we will restrict ourselves to the Palatini formalism for gravitation and will focus on its application to the flat Friedmann-Robertson-Walker (FRW) cosmological model.  We will derive constraints on the two parameters $n$ and $\\beta$ of the $f(R) = R - \\beta/R^{n}$  gravity theory from the most recent compilations of type Ia Supernovae (SNe Ia) observations, which includes the recent large samples from Supernova Legacy Survey (SNLS), the ESSENCE Survey, distant SNe Ia observed with HST, and others, giving a sample of 307 SNe Ia events \\cite{0804.4142}. We also combine the SNe Ia data with information from the baryon acoustic oscillation ({\\rm BAO}) \\cite{Eisenstein2005} and the {\\rm CMB} shift parameter \\cite{wmap} in order to improve the SNe Ia bounds on the free parameters of the theory. ",
        "conclusions": "$f(R)$-gravity based cosmology has presently been thought of as a realistic alternative to general relativistic dark energy models. In this paper, we have worked in the context of a $f(R) = R - \\beta R^{-n}$ gravity with equations of motion derived according to the Palatini approach. We have performed consistency checks and tested the observational viability of these scenarios by using the latest sample of SNe Ia data, the so-called \\emph{Union} sample of 307 events. Although the current SNe Ia measurements alone cannot constrain significantly the model parameters $n$, $\\Omega_{mo}$ and $\\beta$, when combined with information from BAO and CMB shift parameters, the fit leads to very restrictive constraints on the $n - \\Omega_{mo}$ (or, equivalently, $\\beta - \\Omega_{mo}$) parametric space. At 99.7\\% c.l., e.g., we have found the intervals $n\\in [-0.25, 0.35]$ and $\\Omega_{mo} \\in [0.2, 0.31]$ ($\\beta\\in [2.3, 7.1]$). We note that, differently from results in the metric formalism \\cite{Amendola-b}, the universe corresponding to the best-fit solution for a combined SNe Ia+BAO+CMB $\\chi^2$ minimization ($\\Omega_{mo} = 0.26$ and $n = -0.12$) shows all three last phases of the cosmological evolution: radiation era, matter era and a late time cosmic acceleration."
    },
    "0808/0808.1407_arXiv.txt": {
        "abstract": "Young massive stars in the central parsec of our Galaxy are best explained by star formation within at least one, and possibly two, massive self-gravitating gaseous discs. With help of numerical simulations, we here consider whether the observed population of young stars could have originated from a large angle collision of two massive gaseous clouds at $R \\simeq 1$ pc from \\sgra. In all the simulations performed, the post-collision gas flow forms an inner nearly circular gaseous disc and one or two eccentric outer filaments, consistent with the observations. Furthermore, the radial stellar mass distribution is always very steep, $\\Sigma_* \\propto R^{-2}$, again consistent with the data. The 3D velocity structure of the stellar distribution is however sensitive to initial conditions (e.g., the impact parameter of the clouds) and gas cooling details. In all the cases the amount of gas accreted by our inner boundary condition is large, enough to allow \\sgra to radiate near its Eddington limit during $\\sim 10^5$ years. This suggests that a refined model would have physically larger clouds (or a cloud and a disc such as the CND) colliding at a distance of a few pc rather than 1 pc as in our simulations. ",
        "introduction": "The central parsec of our Galaxy is host to young, massive ``He-I'' stars that can be classified into two stellar discs that orbit the central supermassive black hole, \\sgra, outside the inner arcsecond. Most of the stars are located in a well-defined thin stellar disc that rotates clockwise as seen on the sky \\citep{Levin03,GenzelEtal03,PaumardEtal06}. The rest of the stars can be classified as a second more diffuse disc that is more eccentric and rotating counter clockwise \\citep{GenzelEtal03,PaumardEtal06} although the statistical significance of this feature is disputed by \\cite{LuEtal06}. The stellar populations are relatively co-eval, formed approximately $6 \\pm 1$ million years ago \\citep{PaumardEtal06}. In terms of formation scenarios, the ``in-situ'' model is currently the one best favoured by the Galactic Centre (GC) community; in this picture, stars form inside a self-gravitating gaseous disc \\citep[e.g.,][]{Levin03,NC05,PaumardEtal06}. Theorists had expected gaseous massive discs around supermassive black holes to form stars or planets \\citep[e.g.,][]{Paczynski78,Goodman03} long before the properties of the young stars in the GC became known. Recently, \\cite{NayakshinEtal07} have numerically simulated the fragmentation process of a geometrically thin gaseous disc of mass $\\simgt 10^4$ Solar mass for conditions appropriate for our GC (albeit with a rather simple cooling prescription) and found a top-heavy mass function for the stars formed there. \\cite{AlexanderEtAl08} extended these numerical studies to the fragmentation of eccentric accretion disks.  However, the in-situ model has not so far addressed in detail the origin of the gaseous discs themselves. Similar discs are believed to exist in AGN and quasars, and indeed are invoked as a means of feeding a supermassive black hole's immense appetite for gaseous fuel \\citep[e.g.][]{Frank02}. However, \\sgra is currently not a member of the AGN club, as its bolometric luminosity is around $\\sim 10^{-9}$ of its Eddington limit \\citep[e.g.,][]{Baganoff03a}, implying a quiescent mode of accretion.  The viscous timescale of a thin, marginally self-gravitating disc around \\sgra, given a disc mass of $\\approx 10^4 \\msun$ \\citep[reasons for this choice of mass see in, e.g.,][]{NayakshinEtal06} is very long compared with the age of the stellar systems (approx. $10^9$ yrs compared to $10^6$ yrs). In fact, the gaseous discs would have to evolve even faster as the two stellar systems are largely co-eval. The requisite gaseous discs could therefore not have been assembled by viscous transport of angular momentum.  There is also no strong evidence for other similar star formation events in the region within the last $\\sim 10^8$ years. Therefore, the one-off star formation event appears to be best explained by a one-off deposition of gas within the central parsec. There are several ways in which this could have happened, e.g. a Giant Molecular Cloud (GMC) with a sub-parsec impact parameter (in relation to \\sgra) could have self-collided and become partially bound to the central parsec \\citep[e.g.,][]{NC05}. Alternatively, a GMC could have struck the Circumnuclear Disc (CND) located a few pc away from \\sgra, and then created gas streams that settled into the central parsec.  In this article we explore such a one-off collision event in a very simple setup. We allow two massive, uniform and spherical clouds on significantly different orbits to collide with each other one parsec away from \\sgra, and follow the gas dynamics in some detail. We find that the collision forms streams of gas with varying angular momentum, both in magnitude and direction. Parts of these streams collide and coalesce to form a disc in the inner region. As the gas cools, it becomes self-gravitating, and stars are born. The resulting distribution of stellar orbits is compared with the observational data throughout the article, and particularly in \\S \\ref{sec:discussion}. We find that our model is reasonably successful in explaining the population of young He-I stars in the Galactic Centre. ",
        "conclusions": "\\label{sec:discussion}  In approximate order of robustness, our key results are as follows:  \\begin{enumerate}  \\item Formation of a gaseous nearly circular disc in the inner region of the computational domain is common to all runs, as is the ensuing formation of stars on similar circular orbits. This is natural as the dynamical time in the innermost disc is only $\\sim 60 r^{3/2}$ years. Conversely, the outer gaseous stream becomes self-gravitating much faster than it could circularise, and hence orbits of stars in that region are more eccentric, in reasonable agreement with observations \\citep{PaumardEtal06}.  \\item Radial distribution of stellar mass closely follows the observed $\\Sigma_* \\propto 1/R^2$ profile of the disc stellar populations \\citep{PaumardEtal06}. This was observed for all the runs, although we expect these results to change if colliding clouds moved in similar directions, significantly reducing the angular momentum cancellation in shocks and the thermal ``kick'' velocity due to the shock.  \\item Runs with a comparatively long cooling time parameter, $\\beta=1$, lead to kinematically less dispersed stellar populations than runs with faster cooling. As a result, the longer the cooling time, the more closely the resulting stellar system can be fit by planar systems in velocity space. The innermost stellar disc is then reminiscent of the observed clockwise {\\em thin} stellar system \\cite{PaumardEtal06}.  Rapidly cooling runs produce clumpier gas flows that lead to significant gaseous disc orientation changes, producing less coherent discs; such geometrically thick systems are incompatible with the observations.  \\item With the chosen initial conditions, faster cooling promotes survival of the gaseous streams corresponding to the orbits of the original clouds. These streams fragment and form stars mainly in a clustered mode, although this is expected to depend on the details of the radiation feedback from young stars, which is not modelled in this set of simulations.  \\end{enumerate}  \\noindent The observed well defined, flat, geometrically thin and almost circular clockwise stellar system \\citep{PaumardEtal06} is best created via a gentle accumulation of gas. Several independent major gas deposition events lead to a too warped disc, and/or mixed up systems consisting of several stellar rings or discs co-existing at the same radius. To avoid this happening, the inner disc must be created on a timescale longer than the critical rotation time, which is estimated at $t_{\\rm cr} \\sim$~few~$\\times 10^4$~years.  Deposition of gas onto the inner disc will most likely take place over the collision time, $t_{\\rm coll} \\sim R_{\\rm cl}/v_{\\rm cl}$, where $R_{\\rm cl}$ and $v_{\\rm cl}$ are the cloud's size and velocity magnitude. We therefore require that the collision itself be more prolonged than $t_{\\rm cr}$. Estimating the velocity of the cloud at $v_{\\rm cl} \\sim 150$ km/sec, which is of the order of circular velocities in the GC outside the inner parsec, we find $t_{\\rm coll} = R_{\\rm cl}/v_{\\rm cl} \\sim 10^4 \\hbox{years}\\; R_{\\rm cl, pc}$, where $R_{\\rm cl, pc}$ is the size of the cloud in parsecs. We hence require the cloud to be larger than a few parsecs to satisfy $t_{\\rm coll}\\simgt t_{\\rm cr}$. Note that this size is not necessarily the original size of the cloud if the cloud gets tidally disrupted before it makes the impact. In the latter case we can take $R_{\\rm cl}$ to be the radial distance to the centre of the Galaxy at which the tidal disruption took place. Finally, the location of the collision should not be too far from the central parsec, or else too much angular momentum would have to be lost to deposit a significant amount of gas in the $\\sim 0.1$ pc region.  In addition, the radial distribution of gas and stars in our simulations is too compact, contradicting the observations (no He-I stars inside the inner arcsecond and no gaseous disc there either). We thus believe that a realistic scenario would be a GMC of the order of a few parsecs in size striking the CND a few parsecs away from \\sgra. Alternatively such a cloud could self-collide if the impact parameter with respect to \\sgra\\ is small enough, but the cloud would need to be very structured, e.g., essentially consist of several smaller clouds or filaments."
    },
    "0808/0808.3724_arXiv.txt": {
        "abstract": "Strong tidal interaction with the central star can circularize the orbits of close-in planets. With the standard tidal quality factor $Q$ of our solar system, estimated circularization times for close-in extrasolar planets are typically shorter than the ages of the host stars. While most extrasolar planets with orbital radii $a\\la0.1\\,$AU indeed have circular orbits, some close-in planets with substantial orbital eccentricities have recently been discovered. This new class of eccentric close-in planets  implies that either their tidal $Q$ factor is considerably higher, or circularization is prevented by an external perturbation.  Here we constrain the tidal $Q$ factor for transiting extrasolar planets by comparing their circularization times with accurately determined stellar ages. Using estimated secular perturbation timescales, we also provide constraints on the properties of hypothetical second planets exterior to the known ones. ",
        "introduction": "The median eccentricity of the current sample of $\\sim300$ planets is 0.19, while it is 0.013 for close-in planets with semi-major axis $a<0.1\\,$AU. The circular orbits of close-in planets most likely result from orbital circularization due to tides \\citep[e.g.,][]{Rasio96,Marcy97}. This requires the tidal circularization time $\\tau_{\\rm circ}$ to be short compared to the age of the system $\\tau_{\\rm age}$. Since $\\tau_{\\rm circ}$ is a very steep function of $a$ (see Eq.~\\ref{tcirc} or \\ref{tcirc0}), while $\\tau_{\\rm age} \\sim 1-10\\,$Gyr for most systems, a sharp decline in eccentricity is expected below some critical value of $a$. However, the observed transition seems to occur around $0.03-0.04\\,$AU, whereas the calculated \\(\\tau_{\\rm circ}\\) becomes comparable to \\(\\tau_{\\rm age}\\) at $\\sim 0.1\\,$AU as we see below (Figure~\\ref{fig1}). Since almost one quarter (presently 16/68) of planets within $0.1\\,$AU have $e>0.1$, their high eccentricities demand explanation.  First, we calculate the circularization times for transiting planets, and compare them with the estimated ages of the systems. The circularization time $\\tau_{\\rm circ}=-e/\\dot{e}$, where $\\dot{e} $ is the sum of the eccentricity change due to the tides raised on the star by the planet and those raised on the planet by the star, is \\citep{Goldreich66,Hut81,Eggleton98,Mardling02}: \\begin{equation} \\tau_{\\rm circ}= \\frac{2}{81}\\frac{\\Qp^{\\prime}}{n}\\frac{\\Mp}{M_*} \\left(\\frac{a}{\\Rp}\\right)^{5} \\left[\\frac{\\Qp^{\\prime}}{Q_*^{\\prime}}\\left(\\frac{\\Mp}{M_*}\\right)^2 \\left(\\frac{R_*}{\\Rp}\\right)^{5} F_* + \\Fp \\right]^{-1} \\label{tcirc} \\, \\end{equation} The subscripts p and $*$ represent the planet and star, respectively. The modified tidal quality factor for a planet is defined as $\\Qp^{\\prime}\\equiv3\\Qp/2k_{\\rm p}$, where $k_{\\rm p}$ is the Love number, and $\\Qp$ is the specific dissipation function, which depends on the planetary structure as well as the frequency and amplitude of tides. We also define \\begin{eqnarray} F_*&=&\\left[f_1(e^2)-\\frac{11}{18}f_2(e^2)\\frac{\\Omega_{*,{\\rm rot}}} {n}\\right] \\\\ \\Fp&=&\\left[f_1(e^2)-\\frac{11}{18}f_2(e^2)\\frac{\\Omega_{{\\rm p,rot}}} {n}\\right] \\ , \\end{eqnarray} where $n=\\sqrt{G(M_*+\\Mp)/a^3}$ is the mean motion, $\\Omega_{\\rm rot}$ is the rotational frequency, and \\begin{eqnarray} f_1(e^2)&=& \\left(1+\\frac{15}{4}e^2+\\frac{15}{8}e^4+\\frac{5}{64}e^6\\right)/(1-e^2) ^{13/2}, \\\\ f_2(e^2)&=&\\left(1+\\frac{3}{2}e^2+\\frac{1}{8}e^4\\right)/(1-e^2)^5 \\ . \\end{eqnarray} Generally $F_*$ and $F_p$ are comparable, and thus the stellar damping is negligible unless the planet-to-star mass (radius) ratio is large (small) or $Q_*^{\\prime}\\ll Q_p^{\\prime}$. We define the circularization time due to damping in the planet as \\begin{equation} \\tau_{{\\rm circ},0}=\\frac{2}{81}\\frac{\\Qp^{\\prime}}{n}\\frac{\\Mp}{M_*} \\left(\\frac{\\Rp}{a}\\right)^{-5} \\Fp^{-1} \\  \\label{tcirc0}. \\end{equation} Note that $\\tau_{{\\rm circ},0}$ can be shorter or longer than $\\tau_{{\\rm circ}}$, depending on the sign of $F_*$, which changes at $(\\Omega_{*,{\\rm rot}}/n)_{\\rm crit}= 18/11(f_1/f_2)$. In the limit $e\\rightarrow 0$, this equation leads to the standard expression for the circularization time \\citep[Eq.~4.198 of][]{Murray99}.  Figure~\\ref{fig1} compares the circularization times calculated from Eq.~\\ref{tcirc} and \\ref{tcirc0} with the estimated stellar ages for the systems in Table~\\ref{tb1} \\footnote{Note that, in Eq.~\\ref{tcirc}, it is implicitly assumed that the star and the planet both have zero obliquity. Currently available measurements of the Rossiter--MacLaughlin effect show that the planetary orbits in general are closely aligned with the stellar equator \\citep{Queloz00,Winn05}.  Current exceptions may be the HD\\,17156 and XO-3 systems \\citep{Narita07ap,Hebrard08ap}, which we exclude from our analysis.}. Here we assume $\\Qp^{\\prime}=10^5$, and $Q_*^{\\prime}=10^6$, which are the standard values motivated by measurements in our Solar System \\citep[e.g.][]{Yoder81,Zhang08}, and for main-sequence stars \\citep[e.g.][]{Carone07}. For $\\Omega_{\\rm p,rot}$, we assume that planets with circular orbits are perfectly synchronized, i.e., $\\Omega_{\\rm p,rot}/n=1$, since the spin-orbit synchronization times are $\\sim10^{-3}\\,\\tau_{\\rm circ}$ \\citep {Rasio96}. On the other hand, planets with eccentric orbits should spin down until they reach quasi-synchronization \\citep{DobbsDixon04}; in practice we adopt a planetary spin frequency such that the rate of change of spin frequency is zero \\citep[Eq.~54 of][]{Mardling02}. For the stellar spin, we assume typical periods derived from the observed $v_{\\rm rot}\\sin i$, $P_{*,{\\rm rot}}\\sim 3\\,$--\\,70\\,d \\citep{Barnes01}.   \\begin{figure} \\plotone{f1.eps} \\caption[fig1]{ Circularization times calculated from Eq.~\\ref{tcirc0} (stars), Eq.~\\ref{tcirc} with $\\Omega_{*,{\\rm rot}}/n=3\\,$d (orange circles) and 70\\,d (blue circles), compared with the estimated stellar ages (squares) for systems with transiting planets. Here we assume $\\Qp^{\\prime}=10^5$ and $Q_*^{\\prime}=10^6$. For all transiting planets, the estimated circularization time is shorter than the age of the host star. \\label{fig1}} \\end{figure}    Figure~\\ref{fig1} compares the circularization times calculated from Eq.~\\ref{tcirc} and \\ref{tcirc0} with the estimated stellar ages for the systems in Table~\\ref{tb1}.  For most systems the tidal damping in the star is negligible. Two systems with non-negligible stellar damping are WASP-14 and HAT-P-2. Both have a large planetary mass (see Table~\\ref{tb1}), and thus the first term in Eq.~\\ref{tcirc} is significant. Clearly, the estimated circularization times are always shorter than stellar ages, which implies that all these planets should have been circularized by now if their tidal $Q$ values were similar to those of their solar-system analogues. However, our sample contains at least 6 systems with a non-zero orbital eccentricity. Possible explanations are that (1) the structure of these planets is different and their actual $Q$ is greater than what we assumed; or (2) they currently experience external perturbations which maintain their orbital eccentricities against tidal dissipation. ",
        "conclusions": "In this letter, we have investigated the origins of close-in planets on an eccentric orbit.  We place constraints on the tidal Q factor of transiting planets by comparing the stellar age with the tidal circularization time, and find that $10^5\\lesssim \\Qp^{\\prime} \\lesssim 10^9$, which agrees well with current theoretical estimates, can explain these eccentric planets. We also show that it is difficult to explain the high eccentricities of these planets by invoking a current interaction with an unseen second planet.  Our results suggest that at least some of the close-in eccentric planets may be simply in the process of getting circularized.  This work was supported by NSF Grant AST-0507727."
    },
    "0808/0808.2208_arXiv.txt": {
        "abstract": "We conduct a Markov Chain Monte Carlo study of the Dvali-Gabadadze-Porrati (DGP) self-accelerating braneworld scenario given the cosmic microwave background (CMB) anisotropy, supernovae and Hubble constant data by implementing an effective dark energy prescription for modified gravity into a standard Einstein-Boltzmann code.  We find no way to alleviate the tension between distance measures and horizon scale growth in this model.  Growth alterations due to perturbations propagating into the bulk appear as excess CMB anisotropy at the lowest multipoles.  In a flat cosmology, the maximum likelihood DGP model is nominally a $5.3\\sigma$ poorer fit than $\\Lambda$CDM. Curvature can reduce the tension between distance measures but only at the expense of exacerbating the problem with growth leading to a $4.8\\sigma$ result that is dominated by the low multipole CMB temperature spectrum.  While changing the initial conditions to reduce large scale power can flatten the temperature spectrum, this also suppresses the large angle polarization spectrum in violation of recent results from WMAP5.  The failure of this model highlights the power of combining growth and distance measures in cosmology as a test of gravity on the largest scales. ",
        "introduction": "On the self-accelerating branch of the Dvali-Gabadadze-Porrati (DGP) braneworld model \\cite{DGP}, cosmic acceleration arises from a modification to gravity at large scales rather than by introducing a form of mysterious dark energy with negative pressure \\cite{De01}.  In this model our universe is a ($3+1$)-dimensional brane embedded in an infinite Minkowski bulk with differing effective strengths of gravity. The relative strengths define a crossover scale on the brane beyond which ($4+1$)-dimensional gravity and bulk phenomena become important.  A number of theoretical and observational problems of the self-accelerating branch of DGP have been recently uncovered. This branch suffers from pathologies related to the appearance of ghost degrees of freedom \\cite{LuPR03,NicR04,Ko04,GoKS06,ChGNP06,DefGI06,Dv06,KoyS07}. Ghosts can lead to runaway excitations when coupled to normal modes and their existence can invalidate the self-accelerating background solution as well as linear perturbations around it.   Observational problems also arise if one posits that ghosts and strong coupling do not invalidate the idea of self-acceleration itself, {\\it i.e.}~that our Hubble volume behaves in a manner that is perturbatively close on scales near the horizon to the modified Friedmann equation specified on this branch.  These problems fall into two classes, those due to the background expansion history and those due to the growth of structure during the acceleration epoch.  In a flat spatial geometry the DGP model adds only one degree of freedom to fit acceleration, the crossover scale.  Like $\\Lambda$CDM, the extra degree of freedom can be phrased as the matter density relative to the critical density.  Remarkably, with this one parameter, the $\\Lambda$CDM model can fit three disparate sets of distance measures: the local Hubble constant and baryon acoustic oscillations, the relative distances to high-redshift supernovae (SNe), and the acoustic peaks in the cosmic microwave background (CMB).  The flat DGP model cannot fit these observations simultaneously and the significance of the discrepancy continues to grow.  The addition of spatial curvature can help alleviate some, but not all, of this tension \\cite{FaG06,MaM06,SoSH07}.  On intermediate scales that encompass measurements of the large scale structure of the Universe, the early onset of modifications to the expansion also imply a substantial reduction in the linear growth of structure which is itself slightly reduced by the modification to gravity \\cite{LuSS04,KoM06}.  If extended to the mildly nonlinear regime where strong coupling effects may alter growth, this reduction is in substantial conflict with weak lensing data \\cite{WaHMH07}.  On scales approaching the crossover scale, which are probed by the CMB, the unique modification to gravity in the DGP model itself strongly alters the growth rate.  Here the propagation of perturbations into the bulk requires a ($4+1$)-dimensional perturbation framework \\cite{De02}.  The approximate iterative solutions introduced in Rfn.~\\cite{SaSH07} have been recently verified to be sufficiently accurate by a more direct calculation \\cite{CaKSS08} but still are computationally too expensive for Monte Carlo explorations of the DGP parameter space.  More recently, a ($3+1$)-dimensional effective approach dubbed the parameterized post-Friedmann (PPF) framework \\cite{Hu08,HuS07} has been developed that accurately encapsulates modified gravity effects with a closed effective dark energy system.  The PPF approach enables standard cosmological tools such as an Einstein-Boltzmann linear theory solver to be applied to the DGP model.  In this {\\it Paper}, we implement the PPF approach to DGP and conduct a thorough study of the tension between CMB distance, energy density, and growth measures, along with the SNe and local distance measures, across an extended DGP parameter space.  We find that even adding epicycles to the DGP model does not significantly improve the agreement with the data. Curvature, while able to alleviate the problem with distances, exacerbates the problem with growth. Changing the initial power spectrum to remove excess power in the temperature spectrum destroys the agreement with recent polarization measurements from the five-year Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{Noetal08}.  The outline of the paper is as follows.  In \\S \\ref{sec:theory} we review the impact of the DGP modifications on distance measures and the growth of structure emphasizing an effective dark energy PPF approach that is detailed in Appendix A and compared with a growth-geometry splitting approach in Appendix B.  We present the results of the likelihood analysis in \\S \\ref{sec:results} and discuss these results in \\S \\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  We have conducted a thorough Markov Chain Monte Carlo likelihood study of the parameter space available to the DGP self-accelerating braneworld scenario given CMB, SNe and Hubble constant data.  To carry out this study, we have introduced techniques for characterizing modified gravity and non-canonical dark energy candidates with the public Einstein-Boltzmann code CAMB that are of interest beyond the DGP calculations themselves.  We find no way to alleviate substantially the tension between distance measures and the growth of horizon scale fluctuations that impact the low multipole CMB temperature and its relationship to the polarization.  In particular we show that the maximum likelihood flat DGP model is a poorer fit than $\\Lambda$CDM by $2\\Delta \\ln L = 28$, nominally a $(2\\Delta\\ln L)^{1/2}=5.3\\sigma$ result. Interestingly, a substantial ($\\sim 30\\%$) contribution comes from the change in the growth near the crossover scale where perturbations leak into the bulk.  Adding in spatial curvature to the model can bring the distance measures back in agreement with the data but only at the cost of exacerbating the problem with growth.  The net result is that the maximum likelihood only improves by $2\\Delta \\ln L = 5$ with one extra parameter and the difference from $\\Lambda$CDM is $2\\Delta \\ln L=23$, which is still $4.8\\sigma$ discrepant with most of the difference arising from the changes to growth.  Furthermore, while the excess power at large angles can be reduced by changing the initial power spectrum to eliminate large scale power, existing WMAP5 CMB polarization measurements already forbid this possibility.  Specifically, by introducing any finite cut off to the initial power spectrum to flatten the temperature power spectrum, the global likelihood decreases.  While it is still possible that the resolution of the ghost and strong coupling issues of the theory can alter these consequences, it is difficult to see how they can do so without altering the very mechanism that makes it a candidate for acceleration without dark energy.  The failure of this model highlights the power of combining growth and distance measures in cosmology as a test of gravity on the largest scales."
    },
    "0808/0808.2722_arXiv.txt": {
        "abstract": "We present the results of binary population simulations of carbon-enhanced metal-poor (CEMP) stars. We show that nitrogen and fluorine are useful tracers of the origin of CEMP stars, and conclude that the observed paucity of very nitrogen-rich stars puts strong constraints on possible modifications of the initial mass function at low metallicity.  The large number fraction of CEMP stars may instead require much more efficient dredge-up from low-metallicity asymptotic giant branch stars. ",
        "introduction": "One of the most striking results of recent large surveys for very metal-poor stars in the Galactic halo is the large proportion of highly carbon-enriched objects among them. These carbon-enhanced metal-poor (CEMP) stars, usually defined as metal-poor stars with $\\mathrm{[C/Fe]} > 1.0$, make up at least 10 per cent and probably as much as 20--25 per cent of very metal-poor stars with $\\mathrm{[Fe/H]} < -2$ (\\cite[Frebel \\etal\\ 2006]{2006_Frebel}; \\cite[Lucatello \\etal\\ 2006]{2006_Lucatello}).  The majority (about 80 per cent, according to \\cite[Aoki \\etal\\ 2007]{2007_Aoki}) of CEMP stars are also enriched in Ba and other heavy elements produced by slow neutron captures (the $s$-process) in asymptotic giant branch (AGB) stars.  For these so-called CEMP-s stars a likely scenario is pollution by mass transfer from a more massive AGB companion in a binary system, which has since become a white dwarf. Supporting evidence for this scenario comes from radial velocity monitoring, which suggests that all CEMP-s stars could statistically be binaries (\\cite[Lucatello \\etal\\ 2005a]{2005_Lucatello_binfrac}).  The remaining fraction of CEMP stars that do not show s-process enrichments (the CEMP-no stars) are typically more metal-poor than the CEMP-s stars and so far have shown no evidence for binarity (\\cite[Aoki \\etal\\ 2007]{2007_Aoki}). These stars exhibit a variety of abundance patterns, and may have formed instead from material ejected from rapidly rotating massive stars (\\cite[Meynet \\etal\\ 2006]{2006_Meynet}) or from faint core-collapse supernovae (\\cite[Umeda \\& Nomoto 2005]{2005_Umeda}). Confusing such a clear distinction are a surprisingly large number of CEMP stars enhanced in both $s$-process and $r$-process (rapid neutron-capture) elements, whose origin is still quite unclear (\\cite[Jonsell \\etal\\ 2006]{2006_Jonsell}).  Within the mass transfer scenario, the large proportion of CEMP-s stars requires the existence of a sufficient number of binary systems with primary components that have undergone AGB nucleosynthesis. In recent studies (\\cite[Lucatello \\etal\\ 2005b; Komiya \\etal\\ 2007]{2005_Lucatello_IMF,2007_Komiya}) it has been argued that this requires a different initial mass function (IMF) at low metallicity, weighted towards intermediate-mass stars. If true, this in turn has important consequences for the chemical evolution of the halo and, by implication, of other galaxies. However, the model calculations on which these estimates are based still contain many uncertainties regarding the evolution and nucleosynthesis of low-metallicity AGB stars, the efficiency of mass transfer, and the evolution of the surface abundances of the CEMP stars themselves.  In this contribution we explore the effect of some of these uncertainties on the number fraction of CEMP stars by means of a binary population synthesis study. Apart from carbon, we concentrate on nitrogen and fluorine enrichments as possible tracers of the origin of CEMP stars, the latter motivated by the recent discovery of fluorine in a CEMP star (see Sect.~2).  In Sect.~3 we present the first results of our binary population synthesis simulations, and in Sect.~4 we give our conclusions. ",
        "conclusions": "The detection of a large fluorine overabundance in the CEMP star HE\\,1305+0132 is well explained within the AGB binary mass transfer scenario (\\cite[Lugaro \\etal\\ 2008]{2008_Lugaro}). On the other hand, models of rapidly rotating massive stars do not produce fluorine (\\cite[Meynet \\etal\\ 2006]{2006_Meynet}). Therefore fluorine appears to be a useful discriminant between different scenarios proposed for the origin of CEMP stars.  Our population synthesis results indicate that (rather independent of model assumptions) at least 80\\,\\% of CEMP stars formed by AGB mass transfer should be enriched in fluorine ($\\mathrm{[F/Fe]} > 1.0$), i.e.\\ most CEMP-s stars should also be FEMP stars.  Our binary population synthesis models show that the paucity of NEMP stars among metal-poor halo stars is incompatible with a strongly modified IMF at low metallicity, heavily weighted towards intermediate-mass stars, as has been suggested by \\cite{2007_Komiya} in order to explain the high proportion of CEMP stars. Another possible explanation for the ubiquity of CEMP-s stars is that low-metallicity AGB stars undergo much more efficient dredge-up than shown by the detailed evolution models available to date."
    },
    "0808/0808.1337_arXiv.txt": {
        "abstract": "We present our analysis on the clustering properties of star-forming galaxies selected by narrow-band excesses in the Subaru Deep Field. Specifically we focus on H$\\alpha$ emitting galaxies at $z = 0.24$ and $z = 0.40$ in the same field, to investigate possible evolutionary signatures of clustering properties of star-forming galaxies. Based on the analysis on 228 H$\\alpha$ emitting galaxies with $39.8 < \\log L({\\rm H}\\alpha) < 40.8$ at $z = 0.40$, we find that their two-point correlation function is estimated as $\\xi = (r/1.62^{+0.64}_{-0.50}{\\rm Mpc})^{-1.84\\pm0.08}$. This is similar to that of H$\\alpha$ emitting galaxies in the same H$\\alpha$ luminosity range at $z = 0.24$, $\\xi = (r/1.88^{+0.60}_{-0.49}{\\rm Mpc})^{-1.89 \\pm0.07}$. These correlation lengths are smaller than those for the brighter galaxy sample studied by Meneux et al. (2006) in the same redshift range. The evolution of correlation length between $z = 0.24$ and $z = 0.40$ is interpreted by the gravitational growth of the dark matter halos. ",
        "introduction": "Formation and evolution of galaxies and its dependence on the environment, e.g., the formation of large-scale structure, are fundamental problems of the modern cosmology. Studying the evolution of the clustering of galaxies is useful to understand the evolution of galaxies and that of large-scale structures. Galaxies are considered to form within dark matter halos. Owing to the WMAP observations (Spergel et al. 2007), the cosmological parameters are well determined and the evolution of the clustering of dark matter halos is now well understood both analytically (Mo \\& White 1996; Sheth \\& Tormen 1999) and from N-body simulations (Jenkins et al. 1998; Kauffmann et al. 1999; Springel et al. 2005). However, physical relationship between galaxy populations and the dark matter halos is not clear, since galaxy formation processes include gas cooling, star formation, and feedback.  At the present universe, galaxies selected from optical observations are considered to trace the distribution of underlying dark matter halos (Verde et al. 2002). However, the strength of the clustering depends on properties of galaxy populations: e.g., luminosity (Zehavi et al. 2005), color or spectral type (Zehavi et al. 2002), and stellar mass (Li et al. 2006). The dependence of clustering on the properties of galaxies gives constraint on theories of galaxy formation. To constrain the model of galaxy formation, it is useful to investigate such dependences in detail. In this paper, we concentrate the dependence of clustering on the luminosity of galaxies. Norberg et al. (2001) fit the relative bias $b/b^*$ as a function of $L/L^*$, $b/b^* = 0.85 + 0.15 L/L^*$, where $b^*$ is the bias for $L^*$ galaxies and the bias is defined as the square root of the ratio of the galaxy and dark matter correlation functions, $b \\equiv (\\xi_{\\rm g}/\\xi_{\\rm DM})^{1/2}$. Using the SDSS data, Tegmark et al. (2004) gave another fitting function, $b/b^* = 0.85 + 0.15 L/L^* - 0.04(M-M^*)$. The relative bias at the low-luminosity end is smaller than the Tegmark's relation in Zehavi et al. (2005). Therefore the relative bias of low-luminosity galaxies is still controversial.  Recently, the dependence of clustering on the luminosity is studied up to $z \\sim 1.5$ from the VIMOS-VLT Deep Survey (Marinoni et al. 2005; Meneux et al. 2006; Pollo et al. 2006), the DEEP2 Galaxy Redshift Survey (Coil et al. 2006), and the Canada-France Legacy Survey (CFHTLS, McCracken et al. 2008). Brighter galaxies are more strongly clustered than low luminosity galaxies and the dependence of relative bias on the luminosity at $z \\sim 0.8$ is in agreement with that in the local Universe (Marinoni et al. 2005). The observed decreasing of the clustering strength is consistent with simple gravitational growth picture (Meneux et al. 2006). However, the evolution of the clustering of low-luminosity galaxies is not clear till now, since those objects are too faint for spectroscopic surveys.  In this paper, we present the clustering properties of H$\\alpha$ emitters at $z = 0.24$ and $z = 0.40$, which are low-luminosity active star-forming galaxies in the Subaru Deep Field. Their mean absolute magnitude is $\\langle M_B \\rangle \\sim -17.0$, which is 1 mag fainter than the previous observations (Pollo et al. 2006; Meneux et al. 2006). We therefore study the evolution of galaxies with low luminosities. Using the narrowband imaging survey, we select star-forming galaxies within the restricted redshift to fainter magnitudes. Throughout this paper, magnitudes are given in the AB system. We adopt a flat universe with $\\Omega_{\\rm matter} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70 \\; {\\rm km \\; s^{-1} \\; Mpc^{-1}}$. ",
        "conclusions": "Figure \\ref{r0z} shows the correlation length $r_0$ as a function of redshift. For comparison, we also show those at $z < 0.15$ in the 2dFGRS (Norberg et al. 2001) and at $0.2 < z < 1.5$ in the VIMOS-VLT Deep Survey (VVDS) (Meneux et al. 2006) together with our results of H$\\alpha$ emitters at $z = 0.40$ and 0.24. We also show the correlation length of H$\\alpha$ emitters at $z = 0.24$ in the COSMOS field (Shioya et al. 2008). The two-point correlation function of H$\\alpha$ emitters at $z=0.24$ in the COSMOS field, $\\xi = (r/1.87^{+0.21}_{-0.20}{\\rm Mpc})^{-1.88\\pm0.03}$, is very similar to that for our sample at $z=0.24$, although we note that their average $B$-band absolute magnitude, $\\langle M_B \\rangle = -18.5$, is 1.5 mag brighter than our sample.  First, the correlation lengths of our H$\\alpha$ emitters [$39.8 < \\log L({\\rm H}\\alpha) < 40.8$] at $z = 0.24$ and 0.40 are smaller than those of early- and late-type galaxies at $0.2 < z < 0.5$ derived by Meneux et al. (2006) from VVDS. Since the correlation length depends on the luminosity of galaxies, we need to make a fair comparison using absolute magnitudes in the same rest-frame wavelength. The mean $B$-band absolute magnitudes, $\\langle M_B \\rangle$, of our sample with $39.8 < \\log L({\\rm H}\\alpha) < 40.8$ are $\\langle M_B \\rangle = -16.9$ and $\\langle M_B \\rangle = -17.0$ for $z = 0.40$ and $z = 0.24$, respectively. On the other hand, the mean $B$-band absolute magnitudes of Meneux's sample are $\\langle M_B \\rangle = -19.1$ and $-18.1$ for early-type and late-type galaxies, respectively. The smaller correlation length of our H$\\alpha$ emitters is interpreted by their small luminosity.  We note that the clustering amplitude of our luminous H$\\alpha$ emitter sample at $z = 0.40$ is weaker than those of the samples of Meneux et al. (2006). It may imply that the small clustering amplitude of our H$\\alpha$ emitter sample depends not only on the low luminosity but also the spectral type, e.g., the clustering amplitude of red galaxies is larger than that of blue galaxies. As we shown in Figure \\ref{Color}, the color of most of our sample is consistent with Scd - Irr type in Coleman et al. (1980), which is the same as type 3 \\& 4 in Meneux et al. (2006). The clustering amplitude may depends on the equivalent widths of emission lines rather than the global SED. Our sample is biased toward galaxies with large emission-line equivalent width, which are considered to be in a very active phase of the star formation (e.g., Leitherer et al. 1999). We note that there is no significant difference in the correlation length between our two samples although it would be expected that the brighter sample could have a larger correlation length than that for the sample with $39.8 < \\log L({\\rm H}\\alpha) < 40.8$. One possibility is that our brighter sample galaxies are due to brightening by the current star formation activity.  Next, we compare the correlation length of the H$\\alpha$ emitters at $z = 0.40$ with that of the H$\\alpha$ emitters at $z = 0.24$. Although these are consistent within the error, the ratio of the $r_0$ at $z = 0.40$ to that at $z = 0.24$ (0.86) is also consistent with the expected value from the gravitational growth of dark matter halos. This fact implies that the H$\\alpha$ emitters at $z = 0.40$ exist in the same environment where the H$\\alpha$ emitters at $z = 0.24$ exist. If we assume the gravitational growth of dark matter halos to $z \\sim 0$, the correlation length of this kind of galaxies evolve to $r_0 \\sim 2.3$ Mpc. The ratio of the correlation length to that of the $L^*$ galaxies ($M^*=-20.5$) is 0.36. This ratio is smaller than the prediction of Norberg's law. It is suggested that the clustering amplitude of low-luminosity galaxies become smaller for lower luminosity galaxies.  \\vspace{1pc} We would like to thank the Subaru Telescope staff for their invaluable help. We would like to thank Nobunari Kashikawa who is the chief of SDF, and the other SDF members. This work has been financially supported in part by grants of JSPS (Nos. 15340059 and 17253001)."
    },
    "0808/0808.1101_arXiv.txt": {
        "abstract": "It has been known for a long time that the clustering of galaxies changes as a function of galaxy type. This galaxy bias acts as a hindrance to the extraction of cosmological information from the galaxy power spectrum or correlation function. Theoretical arguments show that a change in the amplitude of the clustering between galaxies and mass on large-scales is unavoidable, but cosmological information can be easily extracted from the shape of the power spectrum or correlation function if this bias is independent of scale. Scale-dependent bias is generally small on large scales, $k<0.1\\hompc$, but on smaller scales can affect the recovery of $\\Omega_mh$ from the measured shape of the clustering signal, and have a small effect on the Baryon Acoustic Oscillations. In this paper we investigate the transition from scale-independent to scale-dependent galaxy bias as a function of galaxy population. We use the Sloan Digital Sky Survey DR5 sample to fit various models, which attempt to parametrise the turn-off from scale-independent behaviour. For blue galaxies, we find that the strength of the turn-off is strongly dependent on galaxy luminosity, with stronger scale-dependent bias on larger scales for more luminous galaxies. For red galaxies, the scale-dependence is a weaker function of luminosity. Such trends need to be modelled in order to optimally extract the information available in future surveys, and can help with the design of such surveys. ",
        "introduction": "\\label{sec:intro}  Galaxies are not expected to form a Poisson sampling of the distribution of matter in the Universe. Indeed, it has been known for some time that different populations of galaxies demonstrate different clustering strengths \\citep{davis76,dressler80,park94,pd94,seaborne99,norberg01,norberg02,zehavi02,zehavi05,li06}, showing that they cannot all have a simple relationship linking their distribution with that of the matter. This galaxy bias severely limits our ability to extract cosmological data from galaxy surveys \\citep{percival07,sanchez08}.  The large-scale shape of the linear matter power spectrum is dependent on $\\Omega_mh$ because of the change in the evolution of the Jeans scale after matter-radiation equality \\citep{silk68,peebles70,sunyaev70,bond84,bond87,holtzman89}. However, changes in the general shape of the power spectrum, such as that caused by this physical process, are hard to separate from galaxy bias, which also imprints a signature that changes smoothly with scale. Galaxy bias can also affect mode-coupling in the power spectrum by changing the non-linear scale, which can lead to small changes in the Baryon Acoustic Oscillation positions \\citep{crocce08,matsubara08}. Such effects are beyond the scope of our work, and we concentrate here on the broad changes to the shape of the galaxy power spectrum.  The simplest model of galaxy bias is local, linear, deterministic bias, $\\delta_g({\\bf x})=b_{\\rm lin}\\delta_{\\rm lin}({\\bf x})$ where $\\delta_g$ is the galaxy overdensity field, and $\\delta_{\\rm lin}$ is the linear mass overdensity field. In this model, the bias $b_{\\rm lin}$ is constant in space, but can change for different galaxy populations. A more generic model may also include a stochastic element which enters the description as an additional term, $\\delta_g({\\bf x})=b_{\\rm lin}\\delta_{\\rm lin}({\\bf x})+\\epsilon$. Or the bias could be non-local, where it depends on a smoothed version of the density field. In this paper, we will only be concerned with the relation between the galaxy and mass power spectra, and therefore define a practical measure of galaxy bias \\begin{equation}  \\label{eq:bias} P_g(k) = b(k)^2P_{\\rm lin}(k). \\end{equation} Such a model can incorporate some of the complexities discussed above, but is local in $k$-space, so cannot include any mode coupling terms, which are expected to be present.  As we move to large scales, the galaxy bias is expected to tend towards a constant value. The simplest model for this is the peak-background split model \\citep{BBKS,cole89}. Here galaxy formation depends on the local density field. Large scale density modes can alter the local galaxy number density by pushing pieces of the density field above a critical threshold. On large scales we expect a linear relationship between the large scale mode amplitudes and the change in number density, so the shape of the galaxy and mass power spectra are the same. Interestingly, such a linear relationship is broken if the density field has a non-Gaussian component \\citep{dalal08,slosar08}. On small scales we expect galaxy clustering to be different from that of the mass, as pairs of galaxies inside single collapsed structures become important \\citep{seljak00,peacock00,cooray02}.  Although large-scale bias is seen as a nuisance when extracting cosmological information from galaxy power spectra, it does provide constraints on possible galaxy formation models. \\citet{wild05} proposed bivariate lognormal models of relative bias motivated by observations of galaxy distributions. The data were found to support a small (everywhere $<$5\\%), but significant, amount of stochasticity and non-linearity in these models on all scales, implying that galaxy formation is not solely a function of local density; these effects must be understood in order to properly utilise the next generation of galaxy surveys. A related work, \\citet{conway05}, also supported these findings.  In order to use the power spectrum to extract cosmological information, we need to investigate the transition between scale-independent and scale-dependent galaxy bias. This transition is known to be a function of the galaxy population chosen. Comparing work in \\citet{cole05} and \\citet{percival07}, which used the same techniques, and the analysis of \\citet{sanchez08}, shows that there are deviations between the shapes of the 2dFGRS and SDSS galaxy power spectra. If this results from galaxy bias and the fact that the 2dFGRS selected galaxies in a blue band whilst the SDSS selected galaxies in a red band, then we should expect that similar changes show up in red and blue galaxies drawn from a single survey. Using the SDSS, \\citet{percival07} showed that at $k=0.2\\hompc$, the shape of the power spectrum is a strong function of the $r$-band luminosity. In this paper we extend this work by splitting in galaxy colour and luminosity, and by fitting models to the resulting power spectra, in order to see if the SDSS contains sufficient galaxies to fully explain the trend observed between SDSS and 2dFGRS galaxies.  Previous work examining the observed dependence of the large scale bias on luminosity was carried out by \\citet{norberg01} who proposed the phenomenological model $b=b_{1} + b_{2}L/L_{\\star}$, \\citet{tegmark04} extended this to better fit the available data by including an extra parameter in the form $b=b_{1} + b_{2}L/L_{\\star}+b_{3}(M-M_{\\star})$.  \\citet{swanson08} reexamined both these models for samples split by galaxy colour. \\citet{wild05} also examined the colour dependence of relative bias. \\citet{zehavi05} presented colour dependent and luminosity dependent measurements of the projected correlation function of SDSS galaxies. In all cases subsets of the data existed where some aspect of the bias model in question was shown to be a function of colour or luminosity.  We first describe the SDSS catalogues we use in Section~\\ref{sec:cats}, and how we model the radial selection function of our subsamples using luminosity function fits in Section~\\ref{sec:sel_func}. Section~\\ref{sec:calc_spectra} describes our calculation of power spectra and uncertainties, and the bias models to be fitted are presented in Section~\\ref{sec:bias}. Results of fitting the bias models are given in Section~\\ref{sec:results}. In Section~\\ref{sec:bias_lum} we present simple models for the luminosity dependence of our bias parameters and the results of fitting these to our data. Conclusions and discussion are presented in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have quantified the bias of red and blue galaxies as a function of luminosity on scales $k<0.4\\hompc$. We find clear differences in the large-scale asymptotic bias between blue and red galaxies, similar to that found by \\citet{zehavi05} and \\citet{swanson08}, and shown in Fig.~\\ref{fig:bias_param_vs_mag}. At large scales red galaxies have are more biased than blue galaxies for all luminosities. The bias of blue galaxies is a strong, monotonically increasing function of luminosity, with more luminous galaxies being the most biased. For red galaxies, the picture is more complicated, with an increase in the bias of faint red galaxies (\\citet{zehavi05} find a similar result but stress the scale dependence of the observation). This trend is significant at the $7.2-\\sigma$ level.We find a stronger bias for luminous blue galaxies, compared with that of \\citet{swanson08}. The reason for this is unknown, although it is worth noting that our catalogue is larger than that used by \\citet{swanson08}, as they used additional cuts, constructing volume limited subsamples, compared with our sample selection procedure (see section \\ref{sec:cats}). We performed a simple test of the dependence of this difference on observed scale, repeating our analysis for scales $0.038<k<0.070\\hompc$ which approximates the \\citet{swanson08} measurement scale of $\\sim20\\mpcoh$, the discrepancy in the bias of the brightest blue galaxies was not removed. However, the change of scales brought the best fit red galaxy linear bias model into better agreement with the \\citet{swanson08} faint, red galaxies; the model is poorly constrained at the faint end due to the lack of unique data points in that range. It is worth noting that, due to different colour cut criteria, the brightest absolute magnitude bins in \\citet{swanson08} are inherently more red than ours, this might lead one to expect larger bias measurements for those bins, the opposite of the observed trend. The bright blue sample contains the smallest numbers of galaxies, so the errors on the relative bias should be the largest of any sample. However, our expected errors are insufficient to fully explain the discrepancy, and there remains no obvious reason for this difference.  The work on constant bias models has been extended by considering how the turn-off from constant large-scale bias depends on galaxy colour and luminosity. We have compared two models for this turn-off: the $P$-model and $Q$-model given in Eqns.~\\ref{eqn:ScaleBias_eqn1} \\&~\\ref{eqn:ScaleBias_eqn2} respectively. We find that there is little to choose between the two in terms of how well they can fit the current data, and both provide adequate fits to the power spectrum trends observed as a function of galaxy luminosity and colour at the current level of data precision. Although there is no observational motivation, \\citet{hamann08} argued that the $P$-model has a physical motivation, and therefore offers a more attractive model of bias.  We use the values of $P$ and $Q$ obtained for these two models to quantify the degree of divergence from a constant bias model. Note that these parameters change both the position and amplitude of the turn-off. We find that the best-fit values of $P$ and $Q$, shown in Fig.~\\ref{fig:bias_param_vs_mag}, are a strong function of luminosity for blue galaxies. Red galaxies show far weaker evolution with luminosity, and are consistent with the hypothesis of no change in the scale at which the bias can no longer be described as a constant, to current data precision. Interestingly, this trend in the turn-off from apparent scale-independent behaviour does not match that of the amplitude of the large-scale bias. If amplitude and turn-off scale were linked, we might have expected the $Q$ and $P$ parameters to match for red galaxies of high or low luminosity, where the constant bias component matches, but be different for intermediate luminosity galaxies. We do not observe such a trend, suggesting that there might not be a simple link between the amplitude of bias and the scale at which the 1-halo term becomes important.  It is clear that the simple $P$ and $Q$-models will become insufficient to model the observational data, both on very small scales, and as the data improve. They are, after all, simply motivated to fit observed trends, and do not encompass all of the physics involved. One alternative and more complex prescription for bias prediction is that presented in \\cite{yoo08} which can predict $P_{obs}(k)$ given a set of cosmological parameters and a HOD model, the parameters of which have been determined from the observed, projected correlation function (the correlation function provides additional information to that of the power spectrum as they are sensitive to different scales). However, the potential limits of simple, phenomenological models, such as the $P$ and $Q$-models, does not stop us being able to use them to investigate trends in the data, as we have done in this paper. The current data clearly show that, on relatively large scales (our small scale limit of $0.4\\hompc$ approximately corresponds to $17\\mpcoh$), blue and red galaxies have very different bias properties. It is the blue galaxies that are faint in the $r$-band whose relative clustering on different scales most closely resembles that of the mass, and the luminous blue galaxies the least. This supports the hypothesis that the shapes of the 2dFGRS and SDSS main galaxy power spectra of \\citet{cole05} and \\citet{tegmark04,percival07} differ because the average galaxy bias of each sample differs, caused by different sample selections. The clustering of the 2dFGRS galaxies, on average, would be expected to be a better tracer of the linear clustering signal out to smaller scales than the SDSS galaxies.  This highlights the importance of sample selection for future galaxy surveys, and the importance of understanding galaxy bias for extracting cosmological information from power spectrum shapes from such surveys. Red and blue galaxies show very different trends in their large-scale bias and the scale-dependent smaller-scale bias as a function of luminosity. It is clear that galaxies need to be split into sub-populations by more than just the luminosity in a single band in order to properly understand and model bias."
    },
    "0808/0808.3274_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "{\\sf Conclusions}} We have presented and contrasted two different approximations for the non-linear evolution of $\\alpha$. While in the local approximation we assume that the size of the perturbation which gives rise to spatial variation of $\\alpha$ is much smaller than the Hubble radius, $H^{-1}$, in the infinite wavelength approximation the opposite is required for self-consistency. Hence, it is not surprising that the results obtained using the two models are so different. The crucial question is which one provides the right answer when applied to cosmology. Clearly, the infinite wavelength approximation is not a good approximation in general because we are not interested in an homogeneous perturbation with a size larger than the Hubble radius.  Note that in this approximation not only the gradients are neglected but it is also assumed that the Hubble parameter, $H$, is the local one (not the global $H$). In addition, the fact that the matter density is homogeneously distributed inside a given spherical region does not imply that $\\alpha$ will also be homogeneous within that region.  Therefore we are convinced that local approximation provides the right answer since the non-linear effects are only expected to be important in this context on scales much smaller than $H_0^{-1}$ (scales smaller than the typical galaxy cluster size). On such small scales, we therefore predict that the spatial variations of $\\alpha$ generated in the simplest models should be too small to be of cosmological interest.  These results confirm that it is difficult to generate significant large-scale spatial variations of $\\alpha$ as it is shown in previous chapter for the special cases of cosmic strings and magnetic monopoles \\cite{Nosso,Nosso2} even when we account for the evolution of the fine structure constant in the non-linear regime, when density perturbations become much larger than unity. This results were also confirmed by other authors in Ref. \\cite{Barrow3}--\\cite{Barrow2}.  \\chapter{\\label{capconclusions}{\\sf Outlook}} In this thesis we have studied the cosmological consequences of topological defects, in particular for the dark energy scenario and varying fundamental constants theories. In the following, we summarize our main results and conclusions and address future works.  In Chapter \\ref{cap2}, we have considered the possibility of a domain wall network as a source of dark energy to explain the recent acceleration of the universe. \\begin{itemize} \\item We have presented an analytic model that describes very well the domain wall evolution. We have taken into account the curvature and the interactions of the walls with radiation and other particles and derived a strong bound on the curvature of the walls, which shows that viable candidate networks must be fine-tuned and non-standard. \\item We have studied several typical domain wall models in two spatial dimensions, and discussed the conditions under which various types of defect junctions can exist.We have highlighted the situations where only stable Y-type junctions or X-type ones are formed. In addition we have shown the case where both types co-exist. \\item Using geometrical, topological and energy arguments we distinguished which features of a particular model are crucial to construct the best candidate potential for frustration. However, our results strongly lead us to a no frustration conjecture. In all our numerical simulations, including those of the ideal model, we have found a dynamical behavior consistent with a convergence towards a scaling solution and no evidence of any behavior that could eventually lead to frustration. \\item The results of 3D simulations clearly indicate that no frustrated network can be formed dynamically out of random initial conditions that would mimic a realistic phase transition. Therefore we have presented the most compelling evidence to date that domain wall networks can not be the dark energy. \\end{itemize}  In this context, it is interesting to study the dependence of our results on the number of space-time dimensions. Another relevant problem is the role of the junctions in the dynamics of the network. We recall that the type of junctions formed will depend both on energetic considerations and on the topology of the minima of the potential in field space. In the future we intend to consider further examples of particular models and study this issue with analytic and numerical tools.  In Chapter \\ref{cap3} we have investigated cosmic strings in the context of varying-$\\alpha$ theories. \\begin{itemize} \\item We have considered a generic gauge kinetic function and shown that there is a class of models of this type for which the classical Nielsen-Olesen vortex is still a valid solution. However, in general, the solutions of the local cosmic strings depart from the standard case. \\item We have shown that the Equivalence Principle constraints impose tight limits on the allowed variations of $\\alpha$ on cosmological scales induced by cosmic string networks of this type. \\end{itemize} In Chapter \\ref{cap4} we have investigated the variations of the fine-structure constant in the vicinity of a static local magnetic monopole. \\begin{itemize} \\item We have introduced the Bekenstein-type models in the Non-Abelian Yang-Mills theory and studied various models with varying-$\\alpha$ and confirmed that despite the existence of a class of models for which the 't Hooft-Polyakov standard solution is still valid. \\item We have verified that, in general, the solutions of the local magnetic monopoles depart from the standard case. \\item We have shown that Equivalence Principle constraints impose tight limits on the variations of $\\alpha$ induced by magnetic monopoles which confirms the difficulty to generate significant large-scale spatial variation of the fine structure constant found in the previous section, even in the most favorable case where these variations are seeded by magnetic monopoles. \\end{itemize} In Chapter \\ref{cap5} we have investigated the evolution of the fine-structure constant in the non-linear regime. \\begin{itemize} \\item We have presented and contrasted two different approximations for the non-linear evolution of $\\alpha$. While in the local approximation we have assumed that the size of the perturbation which gives rise to spatial variation of $\\alpha$ is much smaller than the Hubble radius, $H^{-1}$, in the infinite wavelength approximation the opposite is required for self-consistency. \\item We have shown that the infinite wavelength approximation is not a good approximation in general because we are not interested in an homogeneous perturbation with a size larger than the Hubble radius.  Note that in this approximation not only the gradients are neglected but it is also assumed that the Hubble parameter, $H$, is the local one (not the global $H$). In addition, the fact that the matter density is homogeneously distributed inside a given spherical region does not imply that $\\alpha$ will also be homogeneous within that region. \\item We have concluded that local approximation provides the right answer since the non-linear effects are only expected to be important in this context on scales much smaller than $H_0^{-1}$ (scales smaller than the typical galaxy cluster size). On such small scales, we therefore predict that the spatial variations of $\\alpha$ generated in the simplest models should be too small to be of any cosmological interest. \\item We have presented results which confirm that it is difficult to generate significant large-scale spatial variations of $\\alpha$ as it is shown in previous chapter for the special cases of cosmic strings and magnetic monopoles even when we account for the evolution of the fine structure constant in the non-linear regime, when density perturbations become much larger than unity.  \\end{itemize}   Overall, we have assumed that the spatial variations of $\\alpha$ in the vicinity of cosmic strings, magnetic monopoles and compact objects are sourced only by the gauge kinetic function $B_F$ coupled to the Maxwell term. A generalization of this model, we intend to investigate in the near future, includes different source of variation of the fine-structure constant due to a discrete spontaneous symmetry breaking which produces domain walls dividing the universe in patches with different values of $\\alpha$. Of course, the symmetry breaking scale must be sufficiently low in order for the domain walls to be cosmologically benign. This is an interesting model, in which domain walls may be responsible for cosmologically relevant spatial variations of the fine-structure constant as well as other fundamental parameters.  \\chapter*{{\\sf Appendix - Numerical Codes - Domain Wall Network Evolution}} \\small"
    },
    "0808/0808.0936_arXiv.txt": {
        "abstract": "We consider reheating after inflation in theories with a non-standard kinetic term. We show that the equation that exhibits parametric resonance is the Hill-Whittaker equation, which is a particular case of the more general Hill equation. This equation is in general more unstable than the Mathieu equation, such that narrow resonance preheating may be efficient in theories with a non-canonical kinetic term. ",
        "introduction": "\\label{sec:1} Reheating is the period after inflation during which the energy stored in the inflaton is transferred to the matter and radiation of which the universe is made up. During this stage, the inflaton oscillates about its minimum and decays into relativistic particles. Through coupling with other fields, these oscillations can give rise to parametric instability and lead to an explosive growth in particle density. This phase, called \\emph{preheating} \\cite{TB,KLS1,STB,GKLS,KLS2}, plays a crucial role in the dynamics of reheating of an inflationary universe.  It has been shown \\cite{BV,FB2} (see also \\cite{FB,BKM}) that parametric amplification of entropy modes during reheating can lead to a change in the curvature power spectrum on scales larger than the Hubble radius (but smaller than the horizon) without violating causality. Indeed entropy perturbations act as a source for large scale curvature perturbations. The effect is negligible for single-field inflation - where entropy perturbations are suppressed on super-Hubble scales - but can be large in the case of multiple field inflation \\cite{GWBM}. The study of this mechanism in multiple field models demands attention since it can significantly modify the usual inflationary picture.  It has been shown \\cite{Malik} that for the simplest preheating model, the contribution to the large scale power spectrum of curvature perturbations from entropy modes is negligible. However, very little work has been done to this date to show that this is also the case for more complicated scenarios. With the large number of inflationary models proposed today, one can expect the study of this mechanism to become a powerful tool.  In this paper, we study preheating in two-field inflation models $\\{ \\phi, \\chi \\}$ with a non-standard kinetic term for $\\chi$. Such a kinetic term arises inevitably in string-inspired models where the K\\\"{a}hler potential $\\mathcal{K}$ has a non-trivial geometry. Examples of this are brane inflation models, where the inflaton is the distance between two branes in higher dimensional spacetime, and modular inflation, where inflation is caused by the dynamics of the deformations of a six dimensional Calabi-Yau manifold (see e.g. \\cite{Linde,Burgess,Cline,Tye,Eva} for reviews on string inflation models).  Using the general form of a two-field Lagrangian with a non-standard kinetic term introduced in \\cite{diMarco} (see also \\cite{lalak}), we show that the entropy fluctuations during reheating obey the Hill-Whittaker equation, a modified version of the Mathieu equation (the equation that exhibits parametric instability). Indeed, the non-canonical part of the kinetic term enters as an oscillatory damping term into the equation, modifying its stability behaviour. We show that the Floquet theorem applies for this new type of equation such that the modes $\\chi_k$ undergo parametric resonance for a wide range of parameters. Regions of instability appear as bands in the parameter space. We show that they are larger in the case  of the modified Mathieu equation. Therefore preheating is more likely to happen in models with non-standard kinetic terms. ",
        "conclusions": "In this paper, we have studied the simplest preheating model with a non-standard kinetic term. We observed that the non-standard kinetic term introduces an oscillating friction term in the Mathieu equation that modifies the stability behavior of the perturbations in the reheat field $\\chi$. The equation that is obtained is the Hill-Whittaker equation, which in the form in which it appears here is more unstable than the Mathieu equation.  Thus parametric resonance is more likely to be efficient in theories with a non-standard kinetic term. In Section \\ref{nosymbreak}, we showed that for the simplest preheating model, there exists a function $b(\\phi)$ that makes parametric resonance efficient whatever the coupling $g$. Thus one should be careful when dealing with theories with non-standard kinetic terms in the context of reheating as models that do not seem to give rise to preheating at first sight (because they have a small value of $q$) can become resonant due to the oscillatory friction/anti-friction term introduced by the non-standard kinetic term in the action. Note, however, that for small values of $b(\\phi)$, as is realized in the Roulette Inflation model, the effects of the non-standard kinetic term are negligible.  As discussed in a companion paper, equations like those studied in this paper occur in the Roulette Inflation model. There, the perturbations in the reheat field $\\chi$ correspond to entropy fluctuations that are parametrically amplified through the mechanism discussed here. These entropy modes seed curvature perturbations on large scales, altering homogeneity and scale invariance. In a companion paper, we were able to put constraints on the parameters of the Roulette Inflation model so as to match the results from the COBE experiment. Similar work was conducted in the cases of the $D3/D7$ brane inflation model \\cite{BDD} and the KKLMMT model \\cite{BFL}. In both cases it was found that under certain assumptions, during reheating entropy fluctuations can seed a secondary curvature mode that dominates over the primary mode (the purely adiabatic linear perturbation theory mode).  Our results show that the study of parametric resonance during reheating can be used to put constraints on the many inflationary models proposed today. We expect the power of this method to be fully appreciated in the future."
    },
    "0808/0808.2927_arXiv.txt": {
        "abstract": "We investigate the spacetime of anisotropic stars admitting conformal motion. The Einstein field equations are solved using different ansatz of the surface tension. In this investigation, we study two cases in details with the anisotropy as: [1]  $p_t = n p_r$ [2] $p_t - p_r = \\frac{1}{8 \\pi}( \\frac{c_1}{r^2} + c_2)$ where, n, $c_1$ and $c_2$ are arbitrary constants. The solutions yield expressions of the physical quantities like pressure gradients and the mass. ",
        "introduction": "Most of the stars in the galaxies are the main sequence stars which evolve by burning lighter elements into heavier nuclei. Stars massive than $\\sim10M_\\odot$ explode into supernova leaving the core behind which then collapses to form a compact object. The cores are supported by the degenerate pressure of its constituent particles and possess the densities of the relativistic scales i.e. $R_s=2GM/rc^2\\sim1$ (which is the Schwarzschild radius of the star) where $M$ is the mass and $r$ is radius of the compact object. At these densities, relativistic effects dominate and the physical quantities like gradients of pressure and mass are determined by the Tolman-Oppenheimer-Volkoff (TOV) equations for an isotropic and homogeneous compact star. In these stars, the matter is found in stable ground state where the quarks are confined inside the hadrons. It is suggested that the quarks if de-confined into individual $u$, $d$ and $s$ quarks can also yield a stable ground state of matter, which is called `strange matter' \\cite{witten}. Stars composed of mostly strange matter are therefore termed as `strange stars' \\cite{depaolis} (see \\cite{xu,glen} for review on this topic). Generally, superdense stars with mass to size ratio exceeding 0.3 are expected to be composed of strange matter \\cite{tike}. The prime motivation for the existence of strange stars was to explain the exotic phenomena of gamma ray bursts and soft gamma ray repeaters \\cite{cheng,cheng1}. Now with the observations of the Rossi X-ray Timing Explorer, it is convincingly shown that astrophysical source SAX J1808.4-3658 is more likely a strange star \\cite{li}. The transition from the normal hadronic to the strange matter occurs at sufficiently high densities or corresponding low temperatures as $T\\propto V\\propto\\rho^{-1}$. These conditions can mostly be found inside the cores of fast rotating pulsars, P-stars (composed of $u$ and $d$ quarks and are in $\\beta$ equilibrium with electrons) or magnetic field powered magnetars. The density distribution inside these stars need not be isotropic and homogeneous (as proposed in the TOV model) if strange matter truly exists, then stars composed of entirely strange matter can also be found. Recently several authors have studied compact stars with anisotropic matter distribution \\cite{prad,bohm,dev,kg,yav}. We provide star models admitting conformal motion with different anisotropy.  To search the natural relation between geometry and matter through the Einstein's Equations, it is useful to use inheritance symmetry. The well known inheritance symmetry is the symmetry under conformal killing vectors (CKV) i.e. \\begin{equation} L_\\xi g_{ik} = \\psi g_{ik}, \\label{Eq3} \\end{equation} where $L$ is the Lie derivative operator, $\\xi$ is the four vector along which the derivative is taken and $\\psi$ is the conformal killing vector. Note that if $\\psi=0$ then Eq. (1) gives the Killing vector, if $\\psi=constant$ it gives homothetic vector and if $\\psi=\\psi(\\textbf{x},t)$ then it yields conformal vectors. Thus CKV provides a deeper insight into the spacetime geometry. Moreover, if the conformal factor $\\psi=0$, it implies that the underlying spacetime is conformally flat which further implies that the Weyl tensor also vanishes. These conformally flat spacetimes represent gravitational fields without the source of matter producing these fields.  The plan of the paper is as follows: In the second section, we model our gravitational system and formulate the field equations. Next, we shall solve these field equations using different ansatz of our parameters. Then we present the graphical representation of our results. Finally we conclude with the discussion of our results. ",
        "conclusions": "In the previous section, we have given the pictorial representation of the parameters involved in the three cases. The constant of integration is fixed at $C_3=1$ while the radius of the compact star is chosen to be $R=10$km. All the other physical parameters like the metric functions $e^\\lambda$ and $e^\\nu$, radial pressure $p_r$, transverse pressure $p_t$, surface tension $p_r-p_t$, mass function $m(r)$, density $\\rho(r)$ and the radial pressure gradient $dp_r/dr$ are plotted against the radial parameter $r$ in all the cases.\\\\ In figures 1-9, we have taken different values of $n$ in the range (0,1) to determine parameters from case-1. We observe that the radial and transverse pressures gradually decrease as $r$ approaches to the surface $R$. This result is consistent with our boundary condition of vanishing pressures all at the stellar surface. It comes as no surprise that the pressure and the mass profile of the compact stars is all along similar to the normal hydrogen burning stars. More specifically, the mass profile of the star shows sharp increase for $n=\\frac{5}{16}$ while it is small for values $n<\\frac{5}{16}$. The surface tension $p_t-p_r$ of the star also decreases along with increasing $r$. Note that the surface tension varies along negative range due to higher radial pressure. Thus transverse pressure although exists but is smaller then the radial pressure along the allowed range of parameter $r$. The pressure gradient $\\nabla_rp_r$ or $dp_r/dr$ also decreases outward along $r$. For small values of parameter $n$, the gradient is steeper then for large values.  The density profile of the compact star falls for increasing $r$. The conformal factor $\\psi(r)$ assumes maximum value at the star's center while it approaches zero near the stellar surface. Therefore, the vanishing conformal factor at the star's surface could be used as an alternative boundary condition for solving conformal field equations.\\\\ Similarly, the parameters obtained in case-2 are plotted in figures 10-17 using different values of parameter $c_1$. The constant parameter is fixed at $c_2=0.001$. The pressure profile of the star in this case is similar to the earlier in case-1 having steep slopes with increasing $r$. Also, the density profile and the conformal factor show convergence for large values of $r$. The mass profile is steady for $c_1=-0.001$ while it is steepest for $c_1=-0.9.$ Similarly the declining pressure gradient profile is observed for large $r$ which is consistent with our boundary condition Eq.(25).\\\\ In figures 18-25, we have made a comparison of various stellar parameters obtained in cases 1 and 2. It is observed that the metric function $e^\\lambda$, conformal factor $\\psi$, radial pressure $p_r$ and the transverse pressure $p_t$ of both cases converge as $r$ tends to $R$. The mass parameter $m(r)$ is steeper in the second case then the first one. Therefore the mass will increase fast radially in the linear model of the surface tension. Moreover, the surface tension profile of both models goes asymptotically parallel to each other. The density profile in the first case goes to zero while it remains constant in the second case in the asymptotic limit of large $r$.\\\\ One can note that any static spherically symmetric space time admitting conformal motion suffers $r=0$ singularity. In our model also, we had to tolerate $r=0$ singularity. Therefore the model does not exist at $r=0$. However, our solutions are valid for some radius $r>0$. We also note that there is a singularity in the mass density,  in spite of, the total mass is finite. So, the models ( case - I and case - II ) are physically acceptable. In addition to this, one can find that the matter density and fluid pressure are non negative and gradient $\\frac{dp_r}{dr}$ is decreasing with r. Thus our models ( i.e. interior solutions of the gravitational field equations ) are fully physically meaningful. As our models suffer $r=0$ singularity, the minimum value of r is taken to be a positive quantity ( km ). Since radius of the star is taken to 10 km ( this is justified as the radius of some of the observed strange stars are nearly equal to 10 km ), so one can note that the total mass of the star is m ( r=10). Our results have been shown in fig. 7 and fig.15. \\\\ In our current analysis, we have used an idealization of spherical symmetry, stationary and static compact star. In general, from the astrophysical point of view, most of the observed or predicted compact star candidates are no longer static but rotating about a unique axis that may be oblique from the magnetic axis of the star. More exotic stars including magnetars (driven by mostly magnetic fields) and pulsars (mostly driven by higher angular momentum) are well known examples of rotating compact stars. In our forth coming work, we plan to work out a similar analysis presented here for the fast rotating or ultra-fast rotating compact stars. We also plan to work on the rotating stars with the slow rotation approximation as well."
    },
    "0808/0808.0922_arXiv.txt": {
        "abstract": "{Extragalactic globular clusters  have been studied in elliptical galaxies and in a few luminous spiral galaxies, but little is known about globular clusters  in low-luminosity spirals.   } { Past observations with the ACS have shown that NGC~45  hosts a large population of globular clusters (19), as well as several young  star clusters. In this work we aim to confirm the bona fide globular cluster status for 8 of 19 globular cluster candidates and to derive metallicities, ages, and velocities.} {VLT/FORS2 multislit spectroscopy in combination with the Lick/IDS system was used to derive velocities and to  constrain  metallicities and [$\\alpha/$Fe] element ratio of the globular clusters.} {  We confirm the 8 globular clusters as bona fide globular clusters. Their velocities indicate halo or bulge-like kinematics, with little or no overall rotation. From absorption indices such as H$\\beta$, H$\\gamma$, and H$\\delta$  and the combined [MgFe]$'$ index, we found that the globular clusters are metal-poor [Z/H]$\\leq$-0.33 dex  and  [$\\alpha/$Fe]$\\leq$0.0 element ratio. These results argue in favor of a population of globular clusters formed during the assembling of the galaxy.} {} ",
        "introduction": "Globular clusters  are present in almost all kinds of galaxies. Observations of  extragalactic globular cluster systems have shown that globular cluster systems can often be divided into (at least) two sub-populations, although the origin of these remains unclear. In the Milky Way and M31, the metal-poor globular cluster sub-populations display halo-like kinematics and spatial distributions \\citep[e.g.][]{zinn85,ashmanzepf1998,Barmby2000,perrettM31}, while the metal-rich globular clusters may be associated with the bulge and/or thick disk \\citep[e.g.][]{dante95,Barbuy98,cote1999,Bica2006}. The globular cluster sub-populations in elliptical galaxies show many similarities to those in spirals, and some of the metal-rich clusters may have formed in galaxy mergers \\citep[e.g.][]{ashman1992}. Some of the metal-poor (``halo'') clusters may have been accreted from dwarf galaxies \\citep{dacosta} or in proto-galactic fragments from which the halo assembled \\citep{searle}. However, there is evidence based on metallicity that globular clusters could not have been formed from the destruction of dwarf galaxies \\citep[see][and references there-in]{Koch2008}. A major challenge is to establish how each one of these mechanisms may fit into the paradigm of hierarchical structure formation \\citep[e.g.][]{santos}.   One important step towards understanding the roles of merging and accretion processes is to extend our knowledge about globular clusters to many different galaxy types, such as dwarf galaxies and late-type spirals, in a range of environments. Studies of globular clusters in spiral galaxies are more difficult than in early-type galaxies because the globular cluster systems are generally poorer and appear superposed on an irregular background. Consequently, most studies of extragalactic globular clusters have focused on elliptical galaxies.   In spite of the similarities, there are also important differences between globular cluster systems of large ellipticals and those of spirals like the Milky Way. Elliptical galaxies generally have many more globular clusters per unit host galaxy luminosity (i.e., higher globular cluster \\emph{specific frequencies} \\citet{harrisvandenbergh} than spirals), and on average the globular cluster systems of ellipticals are more metal-rich \\citep{markus1999}. The best studied globular cluster systems in spiral galaxies are those associated with the Milky Way and M31. Globular clusters in M31 appear similar to those in the Milky Way in terms of their  luminosity functions,  metallicities and, size distributions \\citep{cramptonM31,perrettM31,barmbyM31}. M33 has a large number of star clusters \\citep{christianM33_82,christianM33_88,chandar2001} but many of them have young ages and there may be only a dozen or so truly ``halo'' globular clusters \\citep{sarajedini2000}. The Magellanic Clouds are also well-known for their rich cluster systems, but again only few of these are truly old globular clusters. The Large Magellanic Cloud has about 13 old globular clusters \\citep{johnson99LMC}, which however show disk-like kinematics. The Small Magellanic Cloud has only one old globular cluster, NGC~121.   The globular cluster populations in low-luminosity spirals are almost unknown. The question of how these clusters formed (and their host galaxy) remain unanswered. Considering that these kind of galaxies remain almost unperturbed during their live, it is probable that we are observing their first population of globular clusters and therefore, we can approach to the conditions in which the host galaxies were formed.  The nearby Sculptor group hosts several late-type galaxies whose globular cluster systems are potentially within reach of spectroscopic observations with 8 m telescopes in a few hours of integration time. A previous study of the star cluster population in the Sculptor group was done by \\citet{knut}. They observed several globular clusters candidates, finding 19  globular clusters in four galaxies, most of them metal-poor with [$\\alpha/$Fe] lower than the Milky Way globular clusters.   In this paper we concentrate on the late-type, low-luminosity spiral galaxy NGC 45 in Sculptor, in which we have previously identified a surprisingly rich population of old globular cluster candidates in HST/ACS imaging. NGC~45 is classified as a  low-luminosity spiral galaxy with $B=11.37 \\pm 0.11$ and $B-V=0.71$ \\citep{paturel}. It is located at $\\sim 5$ Mpc from us, $(m-M)_0=28.42 \\pm 0.41$ \\citep{bottinelli}, in the periphery of the Sculptor group. In \\citet{Mora} we found 19 globular clusters located in projection with the galaxy bulge, which appear to belong to the metal-poor population. Those 19 globular clusters yield a $S_N$ of $1.4-1.9$, which is high for a late-type galaxy.  In this paper we focus on 8 of those 19 globular clusters. We analyze them through spectroscopy to confirm or reject their globular cluster status and to constraint ages and metallicities. ",
        "conclusions": "Although uncertainties on the age estimates for our sample of globular clusters in NGC~45 remain large, we were able to constrain their metallicities and $\\alpha$/Fe abundance ratios. These showed that the globular clusters in NGC~45 are metal poor, corroborating the metal poor population deducted from the globular cluster colors in \\citet{Mora}.  Assuming  that  the globular clusters are tracer of star formation events, considering that NGC~45 is an isolated galaxy, probably a background object and, not a true member of the Sculptor group \\citep{puche}, it is puzzling how these entities formed in this galaxy. \\citet{alan} suggested that NGC~45 once made a close pass by the  Sculptor group, getting close to NGC~7793 and transferring angular momentum. This would have been excited the globular cluster formation in NGC~45 but there is no  record of this in the globular cluster population. Therefore, the formation of the globular clusters  must be happen  at early times probably when the galaxy assembled.  The globular cluster velocities and the velocity dispersion of the system argue in favor of a real  association between the globular clusters and the galaxy bulge. The velocity distribution of the GC population seems to be dominated by random motions, although it would be desirable to corroborate this statement with further velocity measurements of the remaining 11 candidates.    The sub-solar [$\\alpha$/Fe] values derived here are unlike those typically observed in old GC populations \\citep[e.g.][]{puzia2005}, including those in the Milky Way. However, we note that they are consistent with those derived by \\citet{knut}. It is also of interest to note that dwarf galaxies in the Local Group tend to show less alpha-enhanced abundance ratios than the Milky Way \\citep{Sbordone,Tolstoy2003}, despite the fact that many of them are satellite galaxies (e.g. LMC, SMC, etc). This may point to important differences in the early chemical evolution in these different types of galaxies, and is potentially an argument against the notion that a major fraction of the GCs in large galaxies could have been accreted from minor galaxies similar to those observed today.    One possible explanation for the relatively low [$\\alpha$/Fe] ratios of the GCs in NGC~45 is that the formation and assembly of the halo/bulge component took longer than in major galaxies like the Milky Way. This would allow time for Type Ia supernovae to appear and enrich the gas with greater amounts of Fe."
    },
    "0808/0808.2050_arXiv.txt": {
        "abstract": "We present the results of a LIGO search for short-duration gravitational waves (GWs) associated with Soft Gamma Repeater (SGR) bursts.  This is the first search sensitive to neutron star $f$-modes, usually considered the most efficient GW emitting modes. We find no evidence of GWs associated with any SGR burst in a sample consisting of the 27 Dec.\\ 2004 giant flare from SGR 1806$-20$ and 190 lesser events from SGR 1806$-20$ and SGR 1900+14 which occurred during the first year of LIGO's fifth science run. GW strain upper limits and model-dependent GW emission energy upper limits are estimated for individual bursts using a variety of simulated waveforms. The unprecedented sensitivity of the detectors allows us to set the most stringent limits on transient GW amplitudes published to date. We find upper limit estimates on the model-dependent isotropic GW emission energies (at a nominal distance of 10\\,kpc) between $\\sci{3}{45}$ and $\\sci{9}{52}$ erg depending on waveform type, detector antenna factors and noise characteristics at the time of the burst.  These upper limits are within the theoretically predicted range of some SGR models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.3867_arXiv.txt": {
        "abstract": "{We present Minimal Dark Matter and its univocal predictions for Dark Matter observables. During the idm08 conference, PAMELA presented preliminary results showing an anomaly in the positron fraction: we find a good agreement, with a modest astrophysical boost factor.}  \\FullConference{Identification of dark matter 2008\\\\ August 18-22, 2008\\\\ Stockholm, Sweden\\\\ --\\\\ Preprint number: SACLAY--t08/163}    \\begin{document} ",
        "introduction": "The quest for the identification of the missing mass of the universe has been with us since many decades now~\\cite{sum}. While explanations in terms of modifications of Newtonian gravity or General Relativity become more and more contrived, evidence for the particle nature of such Dark Matter (DM) now comes from many astrophysical and cosmological observations. Non-baryonic new particles that may fulfill the r\\^ole of DM have emerged in the latest decades within many Beyond the Standard Model (SM) theories, most notably supersymmetry. These constructions try to naturally explain the hierarchy between the ElectroWeak (EW) scale and the Planck scale and, in doing so, introduce a host of new particles with EW masses and interactions. Some of these particles can be good DM candidates (e.g.\\ the lightest neutralino). DM stability is the result of extra features introduced by hand (e.g.\\ R-parity), usually necessary also to recover many good properties of the SM that are lost in these extensions (automatic conservation of baryon number, lepton number, etc). Finally, the richness of these theories implies the introduction of many unknown new parameters (e.g.\\ all sparticle masses), so that the phenomenology of the DM candidate remains unclear.  The Minimal Dark Matter (MDM) proposal~\\cite{MDM} originates from different motivations: focussing on the Dark Matter problem only, we add to the SM the minimal amount of new physics (just one extra EW multiplet $\\chi$) and search for the minimal assignments of its quantum numbers (spin, isospin and hypercharge) that make it a Dark Matter candidate without ad hoc extra features, and without ruining the good features of the SM. As detailed in the following section, we do find one optimal candidate, and we here focus on it. Its only free parameter (the DM mass) is fixed from  the cosmological DM abundance, so that any DM observable can be univocally predicted.  Indirect searches are one of the most promising ways to detect Dark Matter. DM particles in the galactic halo are expected to annihilate and produce fluxes of cosmic rays that propagate through the galaxy and reach the earth. Their energy spectra carry important information on the nature of the DM particle (mass and primary annihilation channels). Many experiments searched  for signatures of DM annihilations in the fluxes of $\\gamma$ rays, positrons and antiprotons.  At the idm08 conference, the PAMELA experiment~\\cite{PAMELA} reported preliminary results that seem to be the first strong hint for a DM indirect signal. We here assume that PAMELA data will be confirmed and that on-going re-evaluations of the astrophysical backgrounds will confirm previous studies, such that the PAMELA excess implies WIMP DM (non-WIMP DM candidates such as the gravitino may remain viable if unstable~\\cite{Ibarra}; DM candidates with a relic density due to a baryon-like asymmetry are disfavored). In view of its univocal predictions, MDM could have been immediately excluded, rendering this talk unnecessary. After introducing the MDM model we compare its predictions, as previously computed in~\\cite{MDMindirect}, with the PAMELA results. ",
        "conclusions": ""
    },
    "0808/0808.2994_arXiv.txt": {
        "abstract": "Only $\\sim$10\\% of baryons in the Universe are in the form of stars, yet most models of luminous structure formation have concentrated on the properties of the luminous stellar matter.  Such models are now largely successful at reproducing the observed properties of galaxies, including the galaxy luminosity function and the star formation history of the universe. In this paper we focus on the ``flip side'' of galaxy formation and investigate the properties of the material that is not presently locked up in galaxies.  This ``by-product'' of galaxy formation can be observed as an X-ray emitting plasma (the intracluster medium, hereafter ICM) in groups and clusters. Since much of this material has been processed through galaxies, observations of the ICM represent an orthogonal set of constraints on galaxy formation models. In this paper, we attempt to self-consistently model the formation of galaxies and the heating of the ICM.  We set out the challenges for such a combined model and demonstrate a possible means of bringing the model into line with both sets of constraints.  In this paper, we present a version of the Durham semi-analytic galaxy formation model GALFORM that allows us to investigate the properties of the ICM. As we would expect on the basis of gravitational scaling arguments, the previous model (presented in Bower et al.\\ 2006) fails to reproduce even the most basic observed properties of the ICM.  We present a simple modification to the model to allow for heat input into the ICM from the AGN ``radio mode'' feedback.  This heating acts to expel gas from the X-ray luminous central regions of the host halo.  With this modification, the model reproduces the observed gas mass fractions and luminosity-temperature (L--T) relation of groups and clusters.  In contrast to simple ``preheating'' models of the ICM, the model predicts mildly positive evolution of the L--T relation, particularly at low temperatures. The model is energetically plausible, but seems to exceed the observed heating rates of intermediate temperature clusters. Introducing the heating process into the model requires changes to a number of model parameters in order to retain a good match to the observed galaxy properties.  With the revised parameters, the best fitting luminosity function is comparable to that presented in Bower et al.\\ (2006).  The new model makes a fundamental step forward, providing a unified model of galaxy and cluster ICM formation. However, the detailed comparison with the data is not completely satisfactory, and we highlight key areas for improvement. ",
        "introduction": "The success of galaxy formation models is usually measured by their ability to match key observational properties of the galaxy distribution. However, only a small fraction ($\\sim 10\\%$, Cole et al.\\ 2001; Lin et al.\\ 2003; Balogh et al.\\ 2001; 2008) of the baryons in the universe end up as luminous stars. The vast majority of baryons remain in diffuse form, either because they are unable to condense out of the intergalactic medium (for example, because their host dark matter halos are too small to resist heating from the diffuse inter-galactic background radiation; see, e.g., Gnedin 2000; Benson et al.\\ 2002) or because they are ejected from the star forming regions of galaxies by strong feedback (White \\& Frenk 1991; Benson et al.\\ 2003 [Be03]). In general, it is difficult to observe the intergalactic medium directly: because of its low temperature, its properties must be inferred from metal line studies (e.g., Aguirre et al.\\ 2005); however, within groups and clusters of galaxies the intergalactic medium becomes sufficiently hot (and dense) that it can be observed at X-ray wavelengths (Cavaliere et al.\\ 1976). This X-ray emitting plasma is usually referred to as the intracluster medium (ICM).  There is a long history of work attempting to explain the observed properties of the ICM (e.g., Kaiser 1991; Evrard \\& Henry 1991; Tozzi \\& Norman 2001; Voit et al.\\ 2003;  Bode et al.\\ 2007; McCarthy et al.\\ 2008). Generally, it has been concluded that the observed properties cannot be explained by the gravitational collapse of dark matter haloes alone: the energetics of observed clusters and the scaling of the X-ray emission with system temperature suggest that an additional heat source is required. For example, gravitational collapse predicts that the X-ray luminosity of clusters should scale with temperature as roughly $T^2$ (for $T$ more than a few keV), while the observed relation is much steeper, scaling as $\\sim T^{2.8}$ (e.g., Edge \\& Stewart 1991; Markevitch 1998). The steepening of the relation can be explained by heating the ICM so that its central density is lower in lower temperature systems. This is most efficiently achieved by heating the ICM prior to its collapse so that a high minimum adiabat is set, resisting the gravitational compression of the system. Such preheating models have been explored extensively in the literature (e.g.\\, Kaiser 1991; Evrard \\& Henry 1991; Ponman, Cannon \\& Navarro 1999; Balogh et al.\\ 1999; Borgani et al.\\ 2002; McCarthy et al.\\ 2002; Muanwong et al.\\ 2002). {Indeed the authors of the present paper have been strong proponents of the energetic efficiency of the preheating model. However, the source of the preheating energy is rarely explicitly modelled. It is often hypothesised to be associated with galaxy formation, or the growth of supermassive black holes, but there is an inherent tension in these models. The scaling of system entropy with mass ($K_{\\rm vir}\\propto T_{\\rm vir}\\rho_{\\rm vir}^{-2/3}\\,$\\footnote{As is common in the astronomical literature we indicate entropy by the adiabatic index ($K$) of the gas rather than the thermodynamically correct logarithmic quantity. $T_{\\rm vir}$ and $\\rho_{\\rm vir}$ are the characteristic temperature and gas density of the system.}) makes it difficult to simultaneously preheat the IGM to a sufficiently high adiabat that it is able explain the properties of galaxy clusters and yet retain sufficient low entropy gas in lower mass halo to obtain a realistic galaxy population. This is a generic problem --- few models attempt to explain the properties of the ICM while simultaneously accounting for the observed properties of galaxies (for two exceptions see Wu et al.\\ 2000 and Scannapieco et al.\\ 2001). For example, while Bower et al.\\ (2001) explored the effect of heating during galaxy formation on the properties of the ICM, these models did not take into account the back reaction of this heating on the formation of galaxies. We briefly explored the a self-consistent model in Be03, but found that it was not capable of reproducing the observed galaxy luminosity function. It is, nevertheless, possible that a successful pre-heating model may eventually emerge. Two possible strategies include (1) cooling a large fraction of the baryons prior to the pre-heating epoch and then slowing the consumption of this material to prolong star formation to the present epoch, or (2) linking the preheating level (at $z\\sim 2$) to the mass of the present-day halo. The first scheme is at odds with the strong feedback required in many current galaxy formation models since gas is rapidly re-cycled between the cold disk and the halo. The second scheme might be effective if entropy excesses are strongly amplified during halo mergers (e.g., Borgani et al.\\ 2005). However, such a scheme is currently too ill-defined to be implemented in to the semi-analytic models. While we are currently undertaking a series of numerical experiments to better define the effect of halo mergers on the entropy distribution of the gas they contain (McCarthy et al.\\ 2007; McCarthy et al.\\ in prep), the results are currently difficult to interpret, in part because of the lack of consistency between SPH and Mesh-code simulations of galaxy clusters (Voit, Kay \\& Bryan 2005; Mitchel et al.\\ 2008). It is also important to stress that the energetic efficiency of the pre-heating model is only realised if the heating occurs before the gas is incorporated into virialised haloes. This makes the process intrinsically hard to model in the current semi-analytic framework.}  In view of the above difficulties, it is useful to take a step back from the problem. If the properties of galaxies and the low observed stellar mass fraction are set aside, cooling provides an appealing explanation for the observed scalings of the ICM (Voit \\& Bryan 2001; Muanwong et al.\\ 2001). Because the cooling time is closely related to the adiabat (or entropy) of the gas, lower mass systems (with lower characteristic entropy) tend to cool out a larger fraction of their ICM. This is sufficient to reproduce many of the observed trends in X-ray properties, but the implied stellar fractions are much larger than those observed (for a recent discussion see Balogh et al.\\ 2008). {This suggests that a simpler alternative to the pre-heating model is worth further investigation: we need to arrange for feedback to eject much of the cooling gas from the system before, rather than after, allowing it to form stars. In this paper, we explore such a model, introducing a self-consistent `radio-mode' gas ejection scheme into the Bower et al.\\ (2006, hereafter B06) galaxy formation model. We propagate the ejected gas fractions through the merger hierarchy so that the scheme has elements in common with the preheating scenario discussed above. However, since gas is ejected in virialised haloes by `in-situ' heating, it has none of the energetic efficiency of the pre-heating scenario and the required energy injection will inevitably be large.}  The over-cooling problem is closely related to the problems of shaping the galaxy luminosity function and explaining galaxy ``down-sizing'' and the absence of bright blue galaxies at the centres of clusters. In B06 we showed that these problems could be resolved by including a strong ``radio mode'' of AGN feedback in the models (see also Croton et al. 2006: we use the term ``radio-mode'' to distinguish recurrent, largely mechanical AGN feedback resulting from accretion in hydrostatic haloes, from the more radiatively efficient ``quasar-mode'' of AGN activity which we associate with galaxy mergers and disk instabilities). At late times (low redshifts), massive haloes host galaxies with large black holes so that even a small amount of gas cooling out of the ICM and being accreted onto the black hole results in sufficient energy feedback to offset the cooling. In B06 we assumed that this set up a self-regulating feedback loop that prevented any significant amount of gas cooling. Adding the additional requirement that the feedback loop is only effective in hydrostatic haloes (where the sound crossing time is shorter than the cooling time at the cooling radius) creates a natural scale at which the efficiency of galaxy formation falls. This results in a good match to the luminosity function and other observational constraints on the formation and evolution of galaxies (see Birnboim \\& Dekel 2003 and Keres et al.\\ 2005 for further discussion of importance of distinguishing `hydrostatic and ``rapid cooling'' haloes).  In this paper we take the process a step further. We consider the possibility that sufficiently massive black holes not only prevent cooling in their host haloes, but may also inject sufficient energy to expel gas from the halo. As gas is expelled, the central density drops, the cooling time becomes longer and the cooling rate, and hence energy feedback, become smaller. The system will move to a new lower density configuration where the energy feedback just balances the cooling rate. The concept is appealing since the final configuration is set by the cooling time in the halo, while the ejection of gas avoids the excess production of stars. It combines the simplicity of the scheme suggested by Voit \\& Bryan (2001) while offering the potential to give a good match to observed galaxy properties. The model allows us to propagate the effects of heating at early epochs to later times, but it does not implement ``preheating'' in the way envisaged by many previous papers. {As a result, the energy requirements of the model we present are larger than in preheating schemes. Since observational estimates of the pV work required to inflate X-ray cavities suggest that jet powers are only comparable to cluster cooling luminosities (above 3~keV) this is a significant draw-back (e.g., Birzan et al. 2004; 2008; Dunn\\& Fabian 2006; Best et al. 2007) unless these estimates severely underestimate that total jet heating power. The model will also struggle to match the details of the internal properties of clusters --- see the discussion of the entropy profiles of clusters at intermediate radius in McCarthy et al.\\ (2008), for example. }  Nevertheless, an order of magnitude calculation shows that the model is worth further consideration. Examining the properties of clusters in the B06 model, we find that the {\\it total} mass of all the black holes in a cluster of mass $M=3\\times10^{14}\\hMsol$ is typically $\\sim 5\\times10^9\\hMsol$. {In order to estimate the maximum energy contribution from black hole growth, we assume the radio-mode dominates the growth of these large black holes, and the kinetic power of the jet is $0.1\\dot{m}_{\\rm bh}c^2$} (where $\\dot{m}_{\\rm bh}$ is the mass growth rate of the black hole). Under these assumptions, the total heating energy is $\\sim 10^{63}h^{-1}\\erg$, while the potential energy of the baryons is $\\sim GM^2f_{\\rm b}/r_{\\rm vir} = 1.5\\times10^{63}h^{-1}\\erg$ (where $f_{\\rm b}$ is the baryon mass fraction and $r_{\\rm vir}$ is the virial radius of the system).  Since these numbers are comparable, it suggests that black hole heating could eject a substantial fraction of the hot gas from the system. At lower halo masses, the black hole heating would completely dominate the thermal energy of the baryons; while, at higher halo masses, the black hole heating becomes a minor perturbation.  Thus, under these assumptions, the effect of this heating is to establish a new scale of $\\sim 3\\times10^{14}\\hMsol$ on which haloes are able to retain their hot X-ray emitting plasma. In the rest of this paper, we explore this idea in detail, adding flesh to the order of magnitude calculation outlined above. In particular, we take full account of the different channels for black hole mass growth (we assume that black hole growth occurs through the QSO mode does not provide heat the ICM efficiently) and for the effect of heating in subhaloes that are subsequently accreted by the growing cluster.  The methods we adopt here are semi-analytic and based on the techniques described in detail in Cole et al.\\ (2000) (see the recent review by Baugh 2006). We are able to implement feedback on a macroscopic scale without attempting to resolve the detailed physical processes that heat and eject the ICM.  Much simulation work is being devoted to studying the formation of jets and their interaction with the surrounding ICM (Quilis et al.\\ 2001; Churazov et al.\\ 2001; Dalla Vecchia et al.\\ 2004; Heinz et al.\\ 2006; Sijacki \\& Springel 2006). Ultimately, these mechanisms need to be incorporated into cosmological scale simulations of galaxy formation and the ICM (Sijacki et al.\\ 2007; Okamoto et al.\\ 2008).  Unfortunately, the myriad of important physical processes make this direct approach extremely difficult and computationally expensive.  Semi-analytic methods, such as those adopted in the present study, allow a wide range of possible physical processes to be explored with a greatly reduced computational effort. Ultimately, however, the details of the processes we model will need to be justified by high resolution numerical simulations and observations.  The approach we present here is intended to capture the broad-brush energetics and integrated properties of clusters. We assume that the heating effect of the AGN can be captured by a single number that measures the non-gravitational heat input into the system and we adopt a particular form the modification of the system's density profile. In reality the situation is likely considerably more complex: for example, the radial dependence of the heat input may vary between systems (e.g., Heinz et al.\\ 2006 have argued that halo mergers play a vital role in mixing the heat input from jets into ICM), or the energy in an infalling subsystem might be distributed in different ways depending on the shocks generated as it falls into the main halo (McCarthy et al.\\ 2007). As a result of these processes, systems may have different density profiles even though the total energy input is the same and we cannot expect to recover the detailed radial structure of clusters. However, despite this simplification, we will see that the model already captures the global features of observational data well, including the scatter in the X-ray luminosity--temperature (hereafter L--T) correlation.  The structure of the paper is as follows. In Section~2, we describe how we implement the ``radio mode'' heating of the ICM. This section includes a discussion of the parameter values that are modified from those in B06. We present results from the model in \\S3, initially focusing on the X-ray luminosity--temperature (L--T) correlation, and then moving to gas mass fractions and galaxy properties. We present our conclusions and discuss how the model can be developed in \\S4. Through out we use units of $\\hMsol$ for masses. For comparison with X-ray observations, however, the choice of H$_0$ does not scale out of the relations, and we adopt $H_0=73$ $\\hbox{km}\\,\\hbox{s}^{-1}\\,\\hbox{Mpc}^{-1}$. The model assumes $\\Omega_b = 0.045$, $\\Omega_M=0.25$ and $\\Omega_{\\Lambda}=0.75$. ",
        "conclusions": "At the start of this paper, our challenge was to build a model that was simultaneously able to account for the observed properties of galaxies such as the luminosity function and the X-ray properties of groups and clusters. In this way we are seeking a model that accounts for both the number and distribution of stars in the universe and the thermodynamic properties of the material that is left over from the galaxy formation process. The importance of this second aspect is emphasised by looking at the mass fraction that it contains: in galaxy clusters more than 90\\% of the baryons are left over as a by-product. {In currently popular galaxy formation models, including the B06, roughly one eighth of this material has been processed through stars, but a much larger fraction has been accreted by galaxies and then ejected in the form of galactic winds.}  To rise to the challenge of modelling the ICM, we have made a relatively simple modification to the highly successful galaxy formation model presented in B06. In B06, we included a ``radio mode'' of AGN feedback, allowing the energy injected from such sources to offset radiative cooling in hydrostatic haloes. {On its own this model fails spectacularly to reproduce the observed X-ray luminosities of groups and clusters}. In this paper, we have taken the radio-mode feedback process a step further, allowing the radio mode feedback to eject gas from the X-ray emitting regions of hydrostatic haloes.  This modification of the model is largely successfully in reproducing the observed correlations between X-ray luminosity and system temperature. In particular, as well as reproducing the median slope of the relation, the model reproduces the large scatter in the observed luminosities of lower temperature systems. When we take into account the distinction between the model virial temperatures and the observed spectroscopic temperatures the normalisation of the relations is also in good agreement. The steepening of the L--T relation results from gas being ejected from the X-ray emitting regions of lower temperature groups, and the diverse formation histories of these systems drives the large scatter in X-ray properties. The gas mass fractions of the model systems agree reasonably well with the observed trends.  We also examined the evolution of the L--T relation predicted by the model. The observational measurements concentrate on clusters hotter than $5\\keV$. In this regime, the model is in agreement with current observational constraints. However, at lower temperatures the model predicts groups to significantly higher luminosities (at high redshift) than that suggested by simple gravitational scaling. The redshift dependence of the model thus differs from that expected in a simple preheating scenario (where we expect weaker than gravitational evolution). Although the large scatter in the L--T relation below $3\\keV$ makes the evolution hard to measure, this issue clearly deserves further observational effort.  { The major problem for the model is the high level of heating required from the AGN. This is an inevitable consequence of assuming that the heating occurs after the system's collapse (McCarthy et al. 2008). Thus, while the model requires only low values for the efficiency with which matter is accreted onto the central black hole, the predicted heating rates exceed the cooling rates by factors of 10 (at 5 keV) to 100 (below 1 keV) in active systems. This conflicts with observational estimates of cluster heating rates that suggest that radio-mode heating is just sufficient to balance cooling in the more massive systems (Birzan et al. 2004; 2008; Dunn \\& Fabian 2006; 2008). However, Best et al.\\ (2007) finds that the ratio heating rate increases in lower temperature systems, and Nusser et al.\\ (2006) argue that the PdV energy is likely to underestimate the total heat input by a factor 4 to 10. Using the revised radio luminosity calibration of Birzan et al.\\ (2008), we estimate that the model is plausibly compatible with the observed heating rates at $T\\sim 1$~keV and exceeds the observations by a factor 3 at higher temperatures. Clearly a much more detailed comparison of the energetics of the model with observational data is needed, paying careful attention to the observational selection effects and the difficulty in observing bubbles in distant or low surface brightness systems. It may also be possible that a more complex form for the modification of the cluster density profile might result in lower X-ray luminosities for a given energy input. We have investigated whether the model can reproduce the observed L--T relation with lower values of the efficiency parameters $\\epsilon_{SMBH}$ and $\\eta_{SMBH}$. The experiment shows that the model maintains a good match to the lowest energy systems, but that problems occur at intermediate temperatures ($T \\sim 0.5$ keV) where the L--T relation develops a pronounced break that is incompatible with the data. }  The modification of the model alters the properties of the galaxies formed, but we find that small adjustments to the parameters in the B06 model are able to restore reasonably good agreement with observational constraints. In order to achieve a good fit we need to slightly raise the threshold at which haloes become hydrostatic and to decrease the dynamical friction orbital decay rate.  Without the latter modification, excessive merging tends to produce a tail of excessively bright galaxies and a power-law (rather than exponential) break in the luminosity function.  We have demonstrated the success of the model in reproducing the broad-brush observational X-ray properties of groups and clusters. However, a number of issues require closer examination and an improved model that goes beyond a simple ad-hoc modification of the cluster density profile. In particular, we have not attempted to address the detailed entropy profiles of these systems. In its current form, the model is unsuitable for this. To tackle such issues requires two developments. Firstly, we need to consider the distribution of excess gas entropies and to be able to propagate this through the merger hierarchy.  For example, McCarthy et al. (2007) show that a low mass substructure that has experienced little non-gravitational heating tends to drop to the centre of a higher mass halo into which it is accreted with little increase in entropy. In this way it is quite simple to produce a small mass of low entropy (rapidly cooling) material embedded in a halo of high entropy (long cooling time) gas. This complexity is not handled by the currently model since we only track the average non-gravitational energy in a halo. The second aspect of the model that needs consideration is to allow for the possibility that AGN activity and supernovae might eject some gas from haloes even if they are not in the hydrostatic regime. This is a tricky issue that is difficult to assess without resort to direct numerical simulations (eg., Dav\\'e et al. 2008). {Both of these effects will tend to amplify the energy input at earlier epochs making the energetic demands of the model much less daunting. The effect may be particularly relevant at intermediate temperatures where the required heating rates appear most incompatible with the observational estimates.}  In summary, the current model demonstrates that we are well on the way to understanding physical process that set the combined properties of the galaxy population and the thermodynamic history of the intra-cluster medium.  While the model might not yet present a complete solution it provides us with great insight into the physical processes at work."
    },
    "0808/0808.2840_arXiv.txt": {
        "abstract": "We  present the  first study  of early  dark energy  cosmologies using N-body  simulations   to  investigate  the   formation  of  non-linear structure. In contrast  to expectations from semi-analytic approaches, we find that early dark energy  does not imprint a unique signature on the statistics of non-linear structures.  Investigating the non-linear power spectra  and halo  mass functions, we  show that  universal mass functions hold for early dark energy, making its presence difficult to distinguish  from $\\Lambda$CDM.   Since early  dark energy  biases the baryon acoustic oscillation scale, the lack of discriminating power is problematic. ",
        "introduction": "Uncovering the nature of the dominant energy component of the Universe at  the  current epoch,  dark  energy, is  a  central  goal of  modern precision cosmology. The evidence for  the existence of dark energy is strong  (\\citet{riess98}, \\citet{perl99},  \\citet{wmap5}),  however on the theoretical front there are no leading candidates despite a wealth of suggestions  (see \\citet{review}  for a review).   As observational data  becomes more  precise and  extensive due  to current  and future generations  of large surveys  and increasingly  powerful instruments, there  remains a number  of theoretical  challenges to  determine with high accuracy the expected observational consequences of the myriad of possible dark energy models.  Dark  energy  affects  the   Universe  primarily  through  the  global expansion  rate.   This  can  be measured  directly  through  distance measurements  such  as   from  Type  Ia  supernovae  (\\citet{riess98}, \\citet{perl99},  \\citet{kowalski08}) over the  last 10  billion years. However, some dark energy models predict a non-negligible contribution to the early time expansion  behavior, called early dark energy (EDE). To  probe EDE  requires observations  tied to  the very  high redshift universe,  such as  the cosmic  microwave background  (CMB)  or baryon acoustic    oscillations     (BAO).     However,    as     shown    in \\cite{LinderRobbers},  early   dark  energy  can  be   hidden  in  CMB observations, even  while shifting  the intrinsic acoustic  scale upon which BAO measurements rely upon.  Thus failure to detect EDE can bias the cosmological model interpreted from CMB and BAO data.  An alternative and complementary measure of dark energy is through the observation of structure in the  Universe. The change in the dynamical evolution of expansion relative to the $\\Lambda$CDM model, say, alters the  growth  of  structures  in  the inhomogeneous  universe  that  we inhabit.   The statistics of  structure therefore  encodes information about dark energy. Structure measures involve a different weighting of cosmological parameters  than distance measures,  thus the combination breaks  degeneracies.   Moreover,  they  can  test  the  framework  of gravitational growth of  inhomogeneities, important for distinguishing between a  physical dark  energy and changes  to the laws  of gravity. Finally,  growth is a  cumulative process,  depending on  the physical conditions from early  times through the epoch at  which the structure is  observed.    Thus,  observations  of  structure   are  crucial  to understanding cosmology at all  epochs. Because of the implications of early dark energy  for interpreting CMB and BAO  measurements, and for understanding the fundamental physics behind dark energy, the behavior of structure formation and evolution is a key point.  Early dark energy models address the coincidence problem by having the dark energy  evolve relatively slowly compared to  the dominant matter or radiation component; in some  cases motivated by string theory they scale exactly with the background  component.  Indeed, this was one of the first dark energy models \\citep{Wett88}.  EDE is discussed in some detail in  \\cite{DoranRobbers06} and references  therein.  Constraints from  CMB  and primordial  nucleosynthesis  data  restrict the  energy fraction in the  early universe contributed by EDE to  less than a few percent  \\citep{DoranRW07} but  this could  have  significant effects. Some  of   the  implications  for  linear  growth   were  examined  in \\cite{Linder06}     and     \\citet{DoranRobbers06}    and     extended semi-analytically to  non-linear theory by  \\cite{bartDW06}.  The main result of previous studies of the non-linear growth in EDE cosmologies (such  as \\citet{bartDW06}  and  \\citet{FedeliBart07}) is  that for  a fixed   linear  theory   matter  perturbation   amplitude   at  $z=0$, $\\sigma_8$,  there are  more collapsed  structure in  EDE  models than $\\Lambda$CDM.   For  higher   redshifts  this  effect  increases,  for instance \\cite{bartDW06} found that at $z=1$, there are up to an order of magnitude more dark matter  halos for higher halo masses ($M_{halo} \\gtrsim 10^{15} M_{\\odot}$).  This  can be simply interpreted as early dark  energy,  which does  not  appreciably  cluster, suppressing  the growth of  structure at early  times.  Hence, to achieve  the observed level  of current structure,  the early  amplitude of  clustering must have   been   greater.     Other   applications   include   the   high Sunyaev-Zel'dovich amplitude  possibly seen at high  multipoles in the CMB \\citep{SadehRephSilk,WaizmanBart}  and the early  formation of the first stars \\citep{SadehReph08}.  A key aim of this paper  is to use N-body simulations to calculate the predicted number of dark matter halos in EDE cosmologies compared with $\\Lambda$CDM.   In particular,  the non-linear  level of  structure is affected  by the early  growth in  a way  that can  change predictions based on  the usual  linear growth factor  at a particular  epoch.  We examine both the halo mass  function giving the abundance of collapsed objects and the non-linear  matter power spectrum.  N-body simulations are a  robust method  of determining these  statistics and  this paper critically examines predictions previously made analytically.  In Section \\ref{simdetail} we  describe the N-body simulations we used to  probe the  effects of  EDE on  non-linear structure  formation. We present the results for the  abundance of dark matter halos in Section \\ref{halos} and  the results for the non-linear  matter power spectrum in Section \\ref{psresults}. We  discuss the implication of our results and make concluding remarks in Section \\ref{conclude}. ",
        "conclusions": "In this work we have examined  the effects of early dark energy on the non-linear growth  of structure  through N-body simulations.   The key result  is  that  non-linear  large  scale  structure,  including  the abundance of  clusters, is relatively  insensitive to the  presence of EDE.   This result  is in  contrast  to previous  analytic work  which expected a  substantial increase  in non-linear structure  compared to $\\Lambda$CDM with the same $\\sigma_8$.  The implications  of the results presented  here are that  EDE is much more  difficult to detect  than previously  thought.  This  presents a problem  for  dark  energy  cosmology,   not  only  for  the  sake  of understanding dark  energy but because, as shown  in \\citet{DST07} and \\citet{LinderRobbers},  the  presence of  EDE  can significantly  bias distance  measurements  using  baryon  acoustic  oscillations  if  the possibility  of  EDE  is  not  considered.  Fitting  for  EDE  however significantly reduces  the discriminating  power of BAOs.  Ideally it  was hoped  to detect the  presence of EDE  through cluster abundances   (see   \\citet{FedeliBart07}),   such   as   with   future Sunyaev-Zel'dovich,  weak  lensing, X-ray,  or  optical surveys.   This would  provide  an  independent  measure  of EDE  thus  restoring  the discriminating  power   of  BAO  measurements.    However,  this  work demonstrates that this is not the case.  As the  amount of EDE increases,  it becomes easier  to distinguish in that the primordial amplitude  of matter density perturbations must be increased  to  reach the  same  growth  ($\\sigma_8$)  by today.   Such effects in  linear growth and  the CMB anisotropy power  spectra allow that EDE must  contribute less than a few percent  to the early energy density  \\citep{DoranRW07}.   The  acoustic  sound horizon  shifts  by approximately  half   of  this.   However  this   article  shows  that observations of non-linear structure should not be expected to provide a substantial leap in our ability to detect EDE.  This study  has demonstrated that  the halo mass function  formulas of both \\citet{Jenkins01} and \\citet{Warren} are valid for predicting the EDE to $\\Lambda$CDM  results ratio to within $\\sim  10\\%$ or better in the mass  range $\\sim 10^{12}  - 10^{15} M_\\odot$.  Future  studies of EDE should  use these formulas or N-body  simulations when considering halo  abundances,  rather   than  spherical  collapse  motivated  mass functions as have been used in the past.  We have also demonstrated that the Halofit \\citep{Halofit} formula for the non-linear power spectrum, when  using the EDE linear matter power spectrum and growth factor, is accurate for EDE cosmologies for scales larger than  around $k\\sim1\\,h$/Mpc.   On smaller scales  this formula under predicts the  power. Further work  would be needed  to accurately re-calibrate this formula  for EDE on these scales,  perhaps along the lines  of the  way \\citet{2006MNRAS.366..547M}  corrected  Halofit for constant $w$ dark energy cosmologies.  Until such a study is performed we propose, as discussed in  Section \\ref{varys8}, that the results of this study suggest a simple prescription for predicting the non-linear power spectrum for EDE models to within an accuracy of a percent or so relative to the $\\Lambda$CDM prediction."
    },
    "0808/0808.2889_arXiv.txt": {
        "abstract": "{The search for radio emission from extra-solar planets has so far been unsuccessful. Much of the effort in modelling the predicted emission has been based on the analogy with the well-known emission from Jupiter. Unlike Jupiter, however, many of the targets of these radio searches are so close to their parent stars that they may well lie inside the stellar magnetosphere.} {For these close-in planets we determine which physical processes dominate the radio emission and compare our results to those for large-orbit planets that are immersed in the stellar wind.} {We have modelled the reconnection of the stellar and planetary magnetic fields. We calculate the extent of the planetary magnetosphere if it is in pressure balance with its surroundings and determine the conditions under which reconnection of the stellar and planetary magnetic fields could provide the accelerated electrons necessary for the predicted radio emission. } {We show that received radio fluxes of tens of mJy are possible for exoplanets in the solar neighbourhood that are close to their parent stars if their stars have surface field strengths above 1-10G. We show that for these close-in planets,  the power of the radio emission depends principally on the ratio $(N_c/B_\\star^{1/3})^2$ where $N_c$ is the density at the base of the stellar corona, and $B_\\star$ is the stellar surface magnetic field strength.} {Radio emission is most likely to be detected from planets around stars with high-density coronae, which are therefore likely to be bright X-ray sources. The dependence of stellar coronal density on stellar rotation rate and effective temperature is crucial in predicting radio fluxes from exoplanets.}   \\keywords {exoplanets, planetary radio emission, hot Jupiters} ",
        "introduction": "Searches for the radio signatures of extra-solar planets have not yet yielded any detections. The motivation for these searches was the well-known decametre emission from Jupiter which can outshine the radio emission from the quiet Sun at these wavelengths by some four orders of magnitude \\citep{farrell_radio_planets_99}. This emission is believed to originate from non-thermal electrons travelling down the converging field lines at the planet's magnetic poles  \\citep{dulk85,zarka_98}. There are several possible sources for these electrons, including currents generated by corotation breakdown inside the middle magnetosphere, the Io-Jupiter interaction and the reconnection of Jupiter's magnetic field with the solar wind magnetic field \\citep{cowley_01,saur_04}. Here we focus on the third possibility since it is common to all the magnetised planets. If these electrons are accelerated to energies greater than 1-10 keV, they  can be unstable to the electron cyclotron maser instability and as a result they emit at the local electron gyrofrequency $f \\mbox{[MHz] = 2.8B(G)}$. The observed peak of Jupiter's emission at 39.5MHz suggests a field strength of 14.5G at the site of the emission close to Jupiter's surface  \\citep{connerney_98}. The proposal that such a situation might occur in some of the known ``Hot Jupiters\" led to several searches but none of these has revealed the expected emission  \\citep{ryabov_04}. A comprehensive review of the current status of theory and observation is given in \\citet{zarka_07}.  One of the most promising reasons for this lack of detections is that tides raised on planets orbiting close to their host stars rapidly synchronise the axial spin with the orbital motion. In this case, the planet may be rotating slowly, producing only a weak magnetic field.  Since the emission frequency is proportional to the magnetic field strength, this could result in emission at unobservably low frequencies. The other synchronisation that can take place is between the spin of the star and the orbit of the planet. This timescale is set by torques due to the tidal bulges raised on the star by the planet (see, e.g. \\citet{zahn77}) and is generally expected to be longer than the star's main-sequence lifetime even for close-orbiting gas-giant planets.  The observed spin rates of planet-host stars are consistent with this expectation. A notable exception is the late F star tau Boo, which rotates synchronously with the 3.3-day orbit of its planet, whose mass is sufficient to lock the star's rotation within its main-sequence lifetime  \\citep{marcy_51peg_97,lubow_spindown_97}. More recently \\citet{dobbs_dixon_planets_04} have considered the combined effects of tidal torques and angular momentum loss from the star in the form of a wind. They  suggest that synchronisation is more likely among F stars than among the lower mass G and K stars.  The nature of the stellar wind may also have an impact on the expected radio emission, since the wind may compress the stellar magnetosphere and also may strip particles out of the planetary exosphere \\citep{griessmeier_planet_tides_04,griessmeier_planet_winds_05,jaritz_planet_roche_05,lipatov_3Dhybrid_05,stevens_planet_winds_05}.  Fluctuations in the stellar wind pressure and speed that occur as a result of the stellar equivalent of solar coronal mass ejections may, however, enhance any planetary radio emission \\citep{khodachenko_CME_07}.  A recent search for radio emission from systems thought to be embedded in strong stellar winds has yielded no detections, but tight upper limits \\citep{stevens_planet_winds_05,george_radio_planets_07}.  It is clear from this modeling that conditions both exterior to the planet (such as the speed, density and magnetic energy of the stellar wind) and within the planet itself (such as the nature of the dynamo-generated magnetic field) may have an impact on the detectability of exoplanetary radio emission. There is another factor, however, which is the process by which the power carried in the stellar wind is converted into radio emission. Within our own solar system, there is a clear scaling between the magnetic or kinetic power in the solar wind incident on a planet and the observed output radio power \\citep{farrell_radio_planets_99,zarka_01,zarka_07}. This  ``radio-magnetic Bode's law'' can be used to predict the radio emission from a variety of stellar wind or planetary magnetosphere conditions.  In this paper we take a complementary approach. Rather than focussing on the physics of the stellar wind or the planet, we choose simple models for these and focus instead on the process of energy conversion. Our aim is to examine how the radio power scales with different stellar parameters, and  to determine if this scaling is the same for all exoplanets.  \\begin{figure} \\begin{center} \\psfig{figure=8658fig1.ps,width=2.5in} \\caption{When the planet is far from the star, the stellar wind elongates the planetary magnetosphere in the direction away from the star. $R_{\\rm m}$ is the extent of the planetary magnetosphere which is determine by the balance of pressures between the planetary and stellar plasmas. The shapes of the stellar and planetary magnetosphere are shown as dashed lines (not to scale).} \\label{cartoon_1} \\end{center} \\end{figure} \\begin{figure} \\begin{center} \\psfig{figure=8658fig2.ps,width=2.in} \\caption{When the planet is within the stellar magnetosphere, the effective headwind due to the relative motion between the planet and the stellar magnetosphere  elongates the planetary magnetosphere in the azimuthal direction (see also \\citet{zarka_01}).} \\label{cartoon_2} \\end{center} \\end{figure} ",
        "conclusions": "In studying the possible radio emission from exoplanets, we have focussed on the electric field that is generated when the planetary and stellar magnetic fields reconnect. This electric field can accelerate electrons to energies sufficient to power the electron-cyclotron maser instability which is believed to be responsible for much of the radio emission from the planets in our solar system. We have assumed that the mass, radius and magnetic field strength of all exoplanets is the same as Jupiter, but used the observed stellar values. By modeling the electric field developed in this interaction, we have calculated the number of accelerated electrons produced. We find that it scales as the square of the density $N$. Thus the output radio power scales as $N^2$ while the input power from the stellar magnetic field scales as $B^2$. Since for planets immersed in a stellar wind (as is the case for the planets in our own solar system) both the density and the field strength scale as $R_{\\rm orb}^2$, this naturally predicts that the output and input power should be proportional, as is observed for the solar system.  We find, however, that the output power from the inner planets behaves differently. The planets lie inside the closed magnetic field (or magnetosphere) of the star. The magnetospheres of these planets are confined not by the ram pressure of the stellar wind, but by the magnetic pressure of the stellar magnetosphere. The density and magnetic pressure within which they are immersed do not vary with orbital radius in the same way as for the outer planets, and so we would not expect the output and input powers to be proportional. In fact, we find that as the planetary orbital radius decreases, the density rises and so the pool of available electrons increases, but this is exactly balanced by the shrinking of the magnetosphere which reduces the cross-section of the planet. The result is that while the input power might rise, the output power saturates at a value determined by the ratio \\begin{equation} P_{\\rm out} \\propto \\left[ \\frac{N_c}{B_\\star^{1/3}}\\right]^2 \\end{equation} where $N_c$ is the density at the base of the stellar corona and $B_\\star$ is the surface field strength. We conclude, therefore, that the radio power emitted by ``hot Jupiters\" is determined by the way that coronal density scales with stellar field strength as a function of stellar rotation rate and effective temperature."
    },
    "0808/0808.3846_arXiv.txt": {
        "abstract": "Spectral and timing studies of {\\it Suzaku} ToO observations of two SGRs, 1900$+$14 and 1806$-$20, are presented. The X-ray quiescent emission spectra were well fitted by a two blackbody function or a blackbody plus a power law model. The non-thermal hard component discovered by INTEGRAL was detected by the PIN diodes and its spectrum was reproduced by the power law model reported by INTEGRAL. The XIS detected periodicity $P = 5.1998\\pm0.0002$\\,s for SGR\\,1900$+$14 and $P = 7.6022\\pm0.0007$\\,s for SGR\\,1806$-$20. The pulsed fraction was related to the burst activity for SGR\\,1900$+$14. ",
        "introduction": "Magnetars (neutron stars with a super strong surface magnetic field $B \\sim 10^{15}$\\,G -- e.g., \\cite{duncan1992, paczynski1992}) have received considerable attention recently. The magnetar's field strength exceeds the critical field $B_{\\rm c} \\equiv m_{e}^{2}c^{3}/c\\hbar \\approx 4.4 \\times 10^{13}$\\,G, and hence the non-linear effects of quantum electrodynamics must be considered in the processes of radiation transfer. Several studies have now found X-ray sources which are magnetar candidates, namely the soft gamma repeaters (SGRs) and anomalous X-ray pulsars (AXPs). They display X-ray quiescent emission with luminosities $L = 10^{33}$-$10^{35}$\\,erg\\,s$^{-1}$ (e.g., \\cite{woods2006}). The SGRs and a few AXPs also exhibit ``short bursts'' with typical durations $\\Delta{t} \\sim 100$\\,ms and super-Eddington luminosities $L \\sim 10^{40}$\\,erg\\,s$^{-1}$. One of the most remarkable burst phenomena is giant flares from SGRs. Three SGRs have emitted giant flares over the past three decades. The most energetic one had $L \\sim 5 \\times 10^{47}$\\,erg\\,s$^{-1}$, from SGR\\,1806$-$20 on 27 December 2004 \\citep{terasawa2005}. The giant flares have been proposed as candidates for some short gamma-ray bursts (GRB) with durations $\\Delta{t} \\lesssim 2$\\,s (e.g., \\cite{hurley2005}). Delayed X-ray and radio emission following giant flares or short bursts  display similar behavior to GRB afterglows \\citep{frail1999, feroci2003, cameron2005, nakagawa2008}, which may also support this hypothesis.  Two spectral models, a two blackbody function (2BB) and a blackbody plus a power law model (BB$+$PL), have been suggested for the quiescent X-ray emission of the magnetar candidates. A comprehensive study of many short SGR bursts detected by the High Energy Transient Explorer 2 (HETE-2) has shown that 2BB is preferable, although it is an empirical model \\citep{nakagawa2007a}. Although a physical model for the energy spectra remains controversial despite observations by several satellites, strong linear correlations between the lower and higher blackbody temperatures ($kT_{\\mathrm{LT}}$, $kT_{\\mathrm{HT}}$), and between the lower and higher blackbody radii ($R_{\\mathrm{LT}}^2$, $R_{\\mathrm{HT}}^2$) were found, irrespective of the activity states (i.e., the X-ray quiescent emission or the short bursts), using 2BB \\citep{nakagawa2007b}. There are various theoretical models which reproduce the spectra of the quiescent X-ray emission (e.g., \\cite{perna2001, guver2006}). Even more exotic models have been proposed, such as that invoking a p-star \\citep{cea2006}.  One piece of evidence in favor of the magnetar model is that magnetic pressure confines the photon-pair plasma from the giant flare in a (rather) small volume close to the neutron star \\citep{thompson1995}. One complication is the fact that there has not been a secure direct measurement of the magnetic field of a magnetar candidate. One possible such measurement is an absorption feature interpreted to be due to proton cyclotron resonance scattering which has been seen in a spectrum of the precursor emission of a $\\sim1.5$\\,s burst from SGR\\,1806$-$20 \\citep{ibrahim2002}. This kind of feature can also be seen in the spectrum of X-ray quiescent emission from AXP\\,1RXS\\,J170849$-$400910 \\citep{rea2003}. The lack of this feature in most magnetar candidates may be explained by a model which smears the absorption feature due to multiple cyclotron resonant scatterings in a stellar atmosphere and its magnetosphere \\citep{guver2006}.  Several studies have reported periods $P = 5$-$13$\\,s and period derivatives $\\dot{P} = 10^{-11}$-$10^{-13}$\\,s\\,s$^{-1}$ for the magnetar candidates \\citep{woods2006}. It is known that pulse characteristics such as the pulsed fraction and pulse shape vary in time (e.g., \\cite{woods2007}), and involve complex pulse morphology. However, little has been done to elucidate the underlying physics of the pulse characteristics. A recent theoretical study suggests that the pulse characteristics may be explained by trapped fireballs produced by starquakes \\citep{jia2008}.  This paper reports the broadband spectroscopy of two SGRs, 1900$+$14 and 1806$-$20, observed by {\\it Suzaku}. The relation between pulsed fraction and burst rate will also be discussed. The distances are assumed to be 14.5\\,kpc for SGR\\,1900$+$14 \\citep{hurley1999, vrba2000} and 8.7\\,kpc for SGR\\,1806$-$20 \\citep{corbel1997, corbel2004, cameron2005, mcclure2005, bibby2008}. ",
        "conclusions": "The quiescent emission spectra of the SGRs 1900$+$14 and 1806$-$20 measured by the XIS on-board {\\it Suzaku} are well described by either a two blackbody function (2BB) or a blackbody plus a power law model (BB$+$PL). The 2BB temperatures ($kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$) and radii ($R_{\\mathrm{LT}}^2$ and $R_{\\mathrm{HT}}^2$) show good agreement with the $kT_{\\mathrm{LT}}$-$kT_{\\mathrm{HT}}$ and $R_{\\mathrm{LT}}^2$-$R_{\\mathrm{HT}}^2$ correlations reported by \\citet{nakagawa2007b}. A shallow dip is apparent at around 2.3 keV in the SGR\\,1900$+$14 spectra as shown in figure \\ref{fig:spc_summary} (a). However, the XIS hardware team has cautioned that the detector gain has some unresolved uncertainties around 2\\,keV for the window mode\\footnote{The reports are available on http://www.astro.isas.jaxa.jp/suzaku/process/caveats/caveats\\_xrtxis06.html.}, so this dip might not be real. If we treat it as a line feature, however, it could be well represented by a Gaussian shape absorption with $E = 2.3\\pm0.1$\\,keV and $\\sigma = 0.2_{-0.1}^{+0.2}$\\,keV, where $E$ is the line energy and $\\sigma$ is the line width. If the dip is due to the fundamental absorption feature of proton cyclotron resonant scatterings, the dipole magnetic field is $B = 3.7\\times10^{14}(E/2.3\\,\\mathrm{keV})$\\,G. We conducted an absorption line search for both SGRs 1900$+$14 and 1806$-$20 using a narrow Gaussian absorption line model with the line center as a free parameter ranging from 0.2\\,keV to 12\\,keV. No feature was found except for the above mentioned $E \\sim 2.3$\\,keV shallow dip of SGR\\,1900$+$14. The upper limits of the feature are $\\tau < 0.3$ (SGR\\,1900$+$14) and $\\tau < 0.4$ (SGR\\,1806$-$20).  The non-thermal hard X-ray emission of SGR\\,1806$-$20, discovered by INTEGRAL in a 1.6\\,Ms long observation, was detected by the PIN diodes on-board {\\it Suzaku} with a short exposure time ($\\sim20$\\,ks) because of the good sensitivity and the very low background level of this instrument.  The periodicities of SGRs 1900$+$14 and 1806$-$20 (table \\ref{tab:timing_summary}) are detected despite the short exposure time ($\\sim20$\\,ks), again because of the low background level of the {\\it Suzaku} XIS. It has been recognized that the pulsed fractions of SGRs are not stable, and that their pulse shapes are complicated. One example is the fact that the pulsed fraction displayed a clear drop-off around the day of the giant flare from SGR\\,1806$-$20 on 27 December 2004 \\citep{woods2007}. To perform a deeper study of the pulse characteristics of SGR\\,1900$+$14, we compared the pulsed fractions to the X-ray fluxes and burst rate. The pulsed fractions and the X-ray fluxes have been derived both from our work (tables \\ref{tab:timing_summary} and \\ref{tab:spc_summary}) and from the literature \\citep{mereghetti2006, esposito2007, israel2008}, while the burst rate was estimated from the list on the IPN web site\\footnote{The burst list is available on http://www.ssl.berkeley.edu/ipn3/sgrlist.txt.}. Panel (a) in figure \\ref{fig:1900_tvar} shows their time histories over 10 years from 1997 to 2007, where A, B and C indicate the days of the giant flare on 27 August 1998, the intermediate flare on 18 April 2001 and the extraordinary burst activity with more than 40 bursts on 29 March 2006, respectively. The results of the {\\it Suzaku} observation are plotted with star symbols.  There seems to be drop-off of the pulsed fraction around the days of B and C. Panels (b) and (c) in figure \\ref{fig:1900_tvar} give an expanded view of panel (a) for days B and C, respectively. In panels (b) and (c), the pulsed fractions display a clear drop-off after both active days (i.e., B and C). No burst activity was found after the day B for 9 days, and then the quiet burst activity appeared. Since the day C, the burst activity had been gradually disappearing. No bursts were reported by the IPN during the {\\it Suzaku} observation (the star symbol in figure \\ref{fig:1900_tvar}). The known bursts before and after the {\\it Suzaku} observation were detected at 12:19:51 on 29 March 2006 and at 00:23:48 on 5 April 2006, respectively. On the contrary, SGR\\,1806$-$20 displayed one burst during the {\\it Suzaku} observation, and moderate, steady burst activity before and after it. This drop-off is found in response not only to the giant flare but also to the extraordinary burst activity. (This implies that the pulsed fraction is related to burst activity.) One possible explanation is to assume that magnetic confinement stores a vast amount of energy in the neutron star atmosphere or magnetosphere through energy injections such as starquakes, and that the pulsed fraction is then increased. After that the stored energy may be released as bursts either through reconnection, or some sort of instability such as the thermonuclear explosions of X-ray bursts. A recent theoretical study reports that a variety of pulse morphologies can be caused by trapped fireballs released by starquakes \\citep{jia2008}.  We conclude by pointing out that monitoring the time variations of the pulsed fraction and the burst rate is an important element in revealing the burst mechanism. This will be achieved by wide field detectors such as the Monitor of All-sky X-ray Image (MAXI) on-board the international space station \\citep{matsuoka1997}.  \\bigskip We would like to thank K. Hurley for helpful comments and suggestions to improve our paper. We are also grateful to the XIS team for useful discussions on the timing analyses and the calibrations of the window mode. This work is supported in part by a special postdoctoral researchers program in RIKEN. Y.E.N. is supported by the JSPS Research Fellowships for Young Scientists.  \\onecolumn \\begin{figure} \\begin{center} \\FigureFile(160mm,80mm){figure1.eps} \\end{center} \\caption{Quiescent emission spectra of SGR\\,1900$+$14 (a) and SGR\\,1806$-$20 (b).}\\label{fig:spc_summary} \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(160mm,80mm){figure2.eps} \\end{center} \\caption{Folded 0.2-12\\,keV light curves for SGR\\,1900$+$14 (a) and SGR\\,1806$-$20 (b).}\\label{fig:efold_summary} \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(160mm,160mm){figure3.eps} \\end{center} \\caption{Panel (a): time variations of the 2-10\\,keV flux, the pulsed fraction and the burst rate from 1997 to 2007, where A, B and C indicate the days of the giant flare on 27 August 1998, the intermediate flare on 18 April 2001 and the extraordinary burst activity with more than 40 bursts on 29 March 2006, respectively. Panels (b) and (c): expanded plots around the days of B and C, respectively. The {\\it Suzaku} observation is plotted with star symbols.}\\label{fig:1900_tvar} \\end{figure}  \\begin{table} \\caption{A summary of {\\it Suzaku} ToO observations.}\\label{tab:obs_summary} \\begin{center} \\begin{tabular}{llllllll} \\hline\\hline SGR          & SeqNum\\footnotemark[$*$] & \\multicolumn{2}{l}{Observation Date (MJD)} & $T$\\footnotemark[$\\dagger$] & XIS\\footnotemark[$\\ddagger$] & \\multicolumn{2}{l}{Regions\\footnotemark[$\\S$]} \\\\ &                          & Start     & End                            & (ks)         &   & FI  & BI \\\\ \\hline 1900$+$14  & 401022010 & 53826.363 & 53826.911 & 17.0 & 1, 2 & $\\timeform{2.3'}\\times\\timeform{5.0'}/\\timeform{2.3'}\\times\\timeform{6.7'}$ & $\\timeform{2.3'}\\times\\timeform{7.0'}/\\timeform{2.3'}\\times\\timeform{3.2'}$ \\\\ 1806$-$20  & 401021010 & 54189.631 & 54190.063 & 19.6 & 0, 1, 3 & $\\timeform{2.2'}\\times\\timeform{5.3'}/\\timeform{2.2'}\\times\\timeform{3.7'}$ & $\\timeform{2.2'}\\times\\timeform{7.0'}/\\timeform{2.2'}\\times\\timeform{3.5'}$ \\\\ \\hline \\multicolumn{8}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] {\\it Suzaku} sequence number. \\par\\noindent \\footnotemark[$\\dagger$] Net exposure time for each observation. \\par\\noindent \\footnotemark[$\\ddagger$] XIS sensor number. \\par\\noindent \\footnotemark[$\\S$] Extracted source/background regions for front-illuminated (FI) and back-illuminated (BI) CCDs. }\\hss}} \\end{tabular} \\end{center} \\end{table}  \\begin{table} \\small \\caption{A summary of timing analyses.}\\label{tab:timing_summary} \\begin{center} \\begin{tabular}{llll} \\hline\\hline SGR          & Epoch (MJD TDB)    & $P$\\footnotemark[$*$] & $\\dot{P}$\\footnotemark[$\\dagger$]   \\\\ \\hline 1900$+$14  & 53826.001  & $5.1998\\pm0.0002$ & $8.7\\pm1.2$ \\\\ 1806$-$20  & 54189.0009 & $7.6022\\pm0.0007$ & $75\\pm4$ \\\\ \\hline \\multicolumn{4}{@{}l@{}}{\\hbox to 0pt{\\parbox{85mm}{\\footnotesize \\footnotemark[$*$] Pulse periods in units of s. \\par\\noindent \\footnotemark[$\\dagger$] Pulse period derivatives in units of $10^{-11}$\\,s\\,s$^{-1}$. }\\hss}} \\end{tabular} \\end{center} \\end{table}  \\begin{table} \\small \\caption{A summary of spectral parameters. The XIS spectra of SGR\\,1900$+$14 are fitted with 2BB or BB$+$PL, while the XIS and PIN diodes spectra of SGR\\,1806$-$20 are fitted with 2BB$+$HPL or BB$+$PL$+$HPL. The spectral parameters of the non-thermal hard emission of SGR\\,1806$-$20 are not shown in the table because the index and the flux were fixed to $\\Gamma = 1.6$ and $F = 2.9 \\times 10^{-11}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ (see subsection \\ref{spc_ana}), respectively.}\\label{tab:spc_summary} \\begin{center} \\begin{tabular}{llllllllll} \\hline\\hline SGR & Model & $N_{\\mathrm{H}}$\\footnotemark[$*$] & $kT_{\\mathrm{LT}}$ ($kT_{\\mathrm{BB}}$)\\footnotemark[$\\dagger$] & $R_{\\mathrm{LT}}$ ($R_{\\mathrm{BB}}$)\\footnotemark[$\\ddagger$] & $kT_{\\mathrm{HT}}$\\footnotemark[$\\dagger$] & $R_{\\mathrm{HT}}$\\footnotemark[$\\ddagger$] & $\\Gamma$\\footnotemark[$\\S$] & $F$\\footnotemark[$\\|$] & $\\chi^2$ (d.o.f.) \\\\ &      & ($10^{22}$\\,cm$^{-2}$) & (keV) & (km) & (keV) & (km) & & & \\\\ \\hline 1900$+$14 & 2BB      & $2.2\\pm0.3$ & $0.49\\pm0.06$ & $4.6_{-1.0}^{+1.7}$ & $1.5_{-0.2}^{+0.3}$ & $0.36_{-0.11}^{+0.16}$ & $\\cdot\\cdot\\cdot$ & $4.1\\pm0.5$ & 115 (105) \\\\ & BB$+$PL  & $3.2_{-0.7}^{+0.9}$ & $0.29_{-0.10}^{+0.17}$ & $12.3_{-9.3}^{+71.8}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $2.8_{-0.5}^{+0.3}$ & $4.1\\pm0.5$ & 106 (105) \\\\ 1806$-$20 & 2BB$+$HPL      & $7.8_{-0.9}^{+1.2}$ & $0.54_{-0.09}^{+0.08}$ & $3.1_{-1.3}^{+2.7}$ & $2.7_{-1.1}^{+3.4}$ & $0.06_{-0.04}^{+0.19}$ & $\\cdot\\cdot\\cdot$ & $9.9\\pm1.6$ & 196 (192) \\\\ & BB$+$PL$+$HPL  & $8.0_{-1.1}^{+2.0}$ & $0.51_{-0.13}^{+0.09}$ & $3.6_{-1.6}^{+5.9}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $1.6_{-1.1}^{+2.2}$ & $9.9\\pm1.6$ & 196 (192) \\\\ \\hline \\multicolumn{10}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] $N_{\\mathrm{H}}$ denotes the photoelectric absorption with 90\\% confidence level errors. \\par\\noindent \\footnotemark[$\\dagger$] $kT_{\\mathrm{LT}}$, $kT_{\\mathrm{HT}}$ and $kT_{\\mathrm{BB}}$ denote blackbody temperatures with 90\\% confidence level errors. \\par\\noindent \\footnotemark[$\\ddagger$] $R_{\\mathrm{LT}}$, $R_{\\mathrm{HT}}$ and $R_{\\mathrm{BB}}$ denote the emission radius with 90\\% confidence level errors. \\par\\noindent \\footnotemark[$\\S$] $\\Gamma$ denotes the power law index with 90\\% confidence level errors. \\par\\noindent \\footnotemark[$\\|$] $F$ denotes the 2-10\\,keV flux in units of $10^{-12}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ with 68\\% confidence level errors. }\\hss}} \\end{tabular} \\end{center} \\end{table}  \\clearpage"
    },
    "0808/0808.1072_arXiv.txt": {
        "abstract": "{According to previous investigations, the effect of diffusion in the stellar atmospheres and envelopes of hot white dwarfs and subdwarf B (sdB) stars strongly depends on the presence of weak winds with mass-loss rates $\\dot M < 10^{-11} M_{\\odot}/ \\rm yr$.} {As in most of these stars with luminosities $L /L_{\\odot} \\la 100$, no wind signatures have been detected, the mass-loss rates are unknown. In the present paper mass-loss rates are predicted from the original theory of radiatively driven winds.} {The method of solution is modified  so that the usual parametrization of the line force multipliers is not necessary. This is important especially for very thin winds. In addition we checked whether a one-component description is justified. As a consequence of various simplifications, the mass-loss rates are expected to be overestimated.} {Results are presented for effective temperatures in the range $25000 \\, \\rm K \\leq T_{\\rm eff} \\leq 50000 \\, \\rm K$ and for various metallicities between solar and $Z / Z_{\\odot} = 0.01$. For (pre-) white dwarfs and sdB stars a stellar mass of $M_{*} = 0.5 M_{\\odot}$ is assumed. For fixed values of $T_{\\rm eff}$, $M_{*}$, and $Z$, the results predict decreasing mass-loss rates with increasing surface gravity and an increasing dependence of the mass-loss rates on the metallicity. For white dwarfs with $\\log g > 7.0$ no wind solution exists even if the metallicity would be solar. Winds with mass-loss rates around $10^{-11}$ to $10^{-10} M_{\\odot} / \\rm yr$ are predicted for the most luminous sdB stars with surface gravities of $\\log g \\la 5.5$, if the metallicity is not significantly lower than solar. For lower values of $\\dot M$ metals decouple from hydrogen and helium.} {If weak winds with $\\dot M \\la 10^{-12} M_{\\odot} / \\rm yr$ exist, the metals cannot be coupled to hydrogen and helium. This should lead to additional changes in the surface composition, which have not yet been taken into account in the diffusion calculations with and without mass-loss. A possible scenario is the existence of pure metallic winds with mass-loss rates of $\\dot M \\la 10^{-16} M_{\\odot} / \\rm yr$ and with hydrostatic hydrogen and helium.} ",
        "introduction": "The abundance anomalies in several types of chemically peculiar stars are believed to be at least partially due to diffusion in the stellar atmosphere and envelope. The predictions of diffusion calculations strongly depend on the presence of mass-loss. In several papers (Unglaub \\& Bues \\cite{ub98}, \\cite{ub00}, \\cite{ub01}), we investigated the combined effects of gravitational settling, radiative levitation, and weak winds on the chemical composition of hot white dwarfs and subdwarf B (sdB) stars. According to the results, these effects can explain the decreasing abundances of helium and metals during the cooling process of white dwarfs on the upper cooling sequence (with effective temperatures $T_{\\rm eff} \\ga 50000 \\, \\rm K$). The mass-loss rates were estimated either from scaling laws or were a free parameter. The winds were assumed to be ``chemically homogeneous\". If $\\dot M_{l}$ is the mass-loss rate of an element and $\\zeta_{l}$ its mass fraction in the photosphere, then this assumption states that $\\dot M_{l} = \\zeta_{l} \\dot M$, where $\\dot M$ is the total mass-loss rate. Such a wind prevents (if $\\dot M \\ga 10^{-11} M_{\\odot} / \\rm yr$) or retards (if $\\dot M < 10^{-11} M_{\\odot} / \\rm yr$) gravitational settling. The present paper investigates whether the assumed mass-loss rates are consistent with the predictions of the theory of radiation-driven winds and whether these winds may be chemically homogeneous.  In thin winds the one-component description is questionable, because the momentum redistribution via Coulomb collisions between metals (which are preferably accelerated due to the radiative flux) and hydrogen and helium may not be effective enough (e.g. Owocki \\& Puls \\cite{owo02}; Krti\\v cka et al. \\cite{krt03}). This may lead to selective winds in which the metals are expelled from the star, whereas hydrogen and helium are left behind (Babel \\cite{bab95}, \\cite{bab96}). In contrast to chemically homogeneous winds, selective winds should directly change the surface composition.  In the absence of any mass-loss and other disturbing processes (e.g. convective mixing), an equilibrium between gravitational settling and radiative levitation should be expected. Because of saturation effects the radiative force on an element in the stellar atmosphere and envelope depends on its abundance. The radiative force decreases with increasing abundance, so an ``equilibrium abundance\" can be found, for which the radiative force balances the effect of gravitational settling. This theory is described in Chayer et al. (\\cite{chay9a}; \\cite{chay9b}) and references therein. Dreizler \\& Wolff (\\cite{drei99}) and Schuh et al. (\\cite{son02}) have incorporated the diffusion theory into their NLTE stellar atmosphere code, so synthetic spectra have been calculated for an abundance stratification, which is given from the equilibrium condition between gravitational settling and radiative levitation. However, for hot white dwarfs the quantitative agreement with observational results is not satisfactory in many cases. The presence of weak winds may be a possible reason for these discrepancies.  Winds have been detected in pre-white dwarfs with effective temperatures in the range $30000 \\, \\rm  K \\la T_{\\rm eff} \\la 150000 \\, \\rm K$ and surface gravities $3.5 \\la \\log g \\la 6.0$ (cgs units). These objects are in a post-asymptotic giant branch stage of evolution, and they evolve with approximately constant luminosities close to $10^{4} L_{\\odot}$ to higher effective temperatures. As derived from observations and theoretical calculations (see e.g. Kudritzki et al. \\cite{kud06}; Pauldrach et al. \\cite{paul04}; Tinkler \\& Lamers \\cite{tin02}; Herald et al. \\cite{her05}; Koesterke et al. \\cite{koe98}; Koesterke \\& Werner \\cite{jkoe98} and the review of Kudritzki \\& Puls \\cite{kud00}), the mass-loss rates are between $10^{-6}$ and $10^{-9} M_{\\odot} / \\rm yr$. The mass-loss rates of white dwarfs on the cooling sequence (if winds exist at all) are unknown.  The sdB stars can be identified with models of the extreme horizontal branch (EHB) stars with masses of $M_{*} \\approx 0.5 M_{\\odot}$ (Heber \\cite{heb86}; Saffer et al. \\cite{saf94}), effective temperatures of $20000 \\, \\rm K \\leq T_{\\rm eff} \\leq 40000 \\, \\rm K$, and surface gravities of $5.0 \\la \\log g \\la 6.0$. Helium is usually  deficient with number ratios of $10^{-4} \\la \\rm He / \\rm H \\la 0.1$ (e.g. Edelmann et al. \\cite{edel03}; Lisker et al. \\cite{lisk05}). However, these helium abundances are by at least one order of magnitude more than would be expected from the equilibrium condition between gravitational settling and radiative levitation (Michaud et al. \\cite{mic89}). A possible explanation for this discrepancy is the presence of weak winds that retard gravitational settling (Fontaine \\& Chayer \\cite{fon97}; Unglaub \\& Bues \\cite{ub98}). For mass-loss rates $\\dot M \\approx 10^{-13} M_{\\odot} / \\rm yr$, within the lifetimes of sdB stars near the EHB ($\\approx 10^{8} \\rm yr$), the helium abundances would gradually decrease from the solar value to $\\rm He / \\rm H \\approx 10^{-4}$.  The helium deficiencies in sdB stars are accompanied by anomalies in the abundances of heavy elements. Some of the most recent results of spectral analyses are from Geier et al. (\\cite{gei08}), O'Toole \\& Heber (\\cite{otol06}), Edelmann et al. (\\cite{edel06}), Blanchette et al. (\\cite{bla06}), and Chayer et al. (\\cite{chay06}). The hotter sdBs with $T_{\\rm eff} > 30000 \\, \\rm K$ show strong deficiencies in many cases especially of light metals like Al, Mg, O, and Si by more than a factor of $100$ in comparison to the solar abundances, whereas enrichments of elements heavier than iron by a factor near $100$ are not unusual. The abundances of iron and nitrogen are usually close to the solar value. If the scenario with diffusion and chemically homogeneous winds were the correct explanation of these compositions, it should be possible to find a mass-loss rate for which the anomalies of all elements can be explained simultaneously. For the case with solar initial abundance, the calculations of Unglaub \\& Bues (\\cite{ub01}) show that, in the case of chemically homogeneous winds, helium should always be more deficient than at least the elements C, N, and O, which have been taken into account. No mass-loss rate exists that leads to deficiencies in C and O by more than a factor of $100$, whereas helium is only deficient by a factor of ten. This, however, is not an unusual composition in sdB stars (e.g. Heber et al. \\cite{heb00}).  Pulsating and non-pulsating sdB stars with similar stellar parameters coexist in the $T_{\\rm eff}$-$\\log g$ diagram (Charpinet et al. \\cite{char06}). The proposed pulsation mechanism is associated with a local enhancement of iron (or other iron group elements) in a mass depth of about $10^{-7} M_{*}$, which is indeed expected from the equilibrium condition between gravitational settling and radiative levitation (Charpinet et al. \\cite{char97}; Fontaine et al. \\cite{fon03}). However, according to the calculations of Chayer et al. (\\cite{chay04}) and Fontaine et al. (\\cite{fon06}), a weak wind with only $\\dot M = 6*10^{-15} M_{\\odot} / \\rm yr$ would be sufficient to destroy the iron reservoir within a timescale of $10^{7} \\rm yr$. For higher mass-loss rates, this should happen in even shorter time scales. Thus for $\\dot M \\approx 10^{-13} M_{\\odot} / \\rm yr$, which would be required to explain the helium abundances, this pulsation mechanism would become questionable.  Mass-loss has been detected in sdO stars that are more luminous than sdBs (Hamann et al. \\cite{ham81}; Rauch \\cite{rau93}). For sdB stars, up to now there has been no observational proof for the existence of winds. From a quantitative analysis of the $\\rm H \\alpha$ line profiles of 40 sdB stars (Maxted et al. \\cite{max01}), a comparison of synthetic NLTE $\\rm H \\alpha$ line profiles from static model atmospheres with the observations have revealed perfect matches for almost all stars. Only in the four most luminous sdBs have anomalous $\\rm H \\alpha$ lines with a small emission at the line centre been detected, which are possibly  signatures of weak winds (Heber et al. \\cite{heb03}). For the case with $T_{\\rm eff} = 36000 \\, \\rm K$, $\\log g = 5.5$, and $\\log L/L_{\\odot} = 1.51$,  Vink (\\cite{vink04}) found a similar behaviour of $\\rm H \\alpha$ from a spectral analysis of $\\rm H \\alpha$ with his wind code, if the existence of a weak wind with $\\dot M \\approx 10^{-11} \\rm M_{\\odot} / \\rm yr$ is assumed.  For several effective temperatures in the range $25000 \\, \\rm K \\leq T_{\\rm eff} \\leq 50000 \\, \\rm K$ and for various metallicities, mass-loss rates will be predicted according to the theory of radiation-driven winds from Castor et al. (\\cite{cas75}, henceforth CAK). As the usual scaling laws may be unreliable for thin winds, the method of solution has been changed slightly (see Sect. 2), so it is taken into account that the radiative force is limited for small wind optical depth parameters and tends to some maximum value. To simplify the numerical method, later improvements in the CAK theory are neglected. As discussed in Sect. 6, it should be possible to derive at least an upper limit for the mass-loss rates. For each wind model, we checked whether a one-component description may be justified (see Sect. 3). The results for $T_{\\rm eff} = 40000$ and $50000 \\, \\rm K$ and several metallicities, which have been obtained with line force multipliers according to Kudritzki (\\cite{kud02}), are presented in Sect. 4. The results for $T_{\\rm eff} = 25000$, $30000$, and $35000 \\, \\rm K$ (Sect. 5) were obtained with force multipliers according to our own calculations (see Sect. 2.4.2). The predicted mass-loss rates are compared with the results of Vink \\& Cassisi (\\cite{vink02}).  In Sects. 6.3 and 6.4, we discuss how the abundance anomalies in sdB stars could be explained. The arguments are also relevant for hot white dwarfs and chemically peculiar main sequence stars, which are reviewed by Smith (\\cite{smith96}). A subgroup are the HgMn stars. Because they are characterised by very low rotational velocities and weak or non-detectable magnetic fields, they are one of the best natural laboratories for studying the competing processes of gravitational settling and radiative diffusion (Vauclair \\& Vauclair \\cite{vvg82}). The HgMn stars are characterised by enhancements of heavy metals (e.g. Hg, Mn, Pt, Sr, Ga), deficiencies of some light elements (e.g. He, Al, N), and isotopic anomalies of metals. Some of the most recent papers about these stars are from Zavala et al. (\\cite{zav07}) and Adelman et al. (\\cite{adel06}). Similar to sdB stars, weak winds with mass-loss rates of $10^{-14}$ to $10^{-12} M_{\\odot} / \\rm yr$ could lead to abundance anomalies that are different from the ones obtained from the equilibrium condition between gravitational settling and radiative levitation (Landstreet et al. \\cite{land98}). However, for the typical stellar parameters $10500 \\, \\rm K \\leq T_{\\rm eff} \\leq 16000 \\, \\rm K$ and $\\log g \\approx 4.0$, these winds can hardly be chemically homogeneous, because the radiative force is too low (Babel \\cite{bab96}).  In chemically peculiar main sequence stars, the abundance anomalies depend on stellar rotation  (see e.g. Vauclair \\cite{vauc03}), on the presence of magnetic fields (Turcotte \\cite{tur03}) and of convection zones. In hydrogen-rich hot white dwarfs and sdB stars, these effects do not seem to be the most important ones. In general both are slow rotators (Karl et al. \\cite{karl05}; Koester et al. \\cite{koes98}; Heber et al. \\cite{heb97}; Heber \\& Edelmann \\cite{heb04}). O'Toole (\\cite{otol05}) finds no correlation between magnetic field strengths in sdB stars and abundance anomalies. In some (pre-) white dwarfs, magnetic fields have been detected (Jordan et al. \\cite{jor05}; \\cite{jor07} and references therein). However, diffusion is effective in all white dwarfs and not restricted to a subgroup of them. A thin superficial convection zone with a mass depth of about $10^{-12} M_{*}$ may be present in sdB stars, but only for helium abundances $\\rm He / \\rm H \\ga 0.01$ by number (Groth et al. \\cite{gro85}). ",
        "conclusions": "In the $T_{\\rm eff}$-$\\log g$-diagram of Fig. 13 for $25000 \\, \\rm K \\leq T_{\\rm eff} \\leq 40000 \\, \\rm K$ and for $Z/Z_{\\odot} = 1$, $1/3$ and $0.1$, the lines are shown above which chemically homogeneous winds may exist according to the present results. The samples of sdB stars analysed by Maxted et al. (\\cite{max01}) and Lisker et al. (\\cite{lisk05}) shown in the figure have been analysed for peculiar $\\rm H_{\\alpha}$ line profiles, which may be interpreted as wind signatures according to Heber et al. (\\cite{heb03}) and Vink (\\cite{vink04}). As can be seen in the figure, such signatures have only been detected in the most luminous sdB stars. According to the calculations of stellar evolution by Dorman et al. (\\cite{dor93}), these ones are in a post-EHB stage of evolution. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{8019f13.eps}} \\\\[1.0cm] \\caption{Lines in the $T_{\\rm eff}$ - $\\log g$ diagram above which chemically homogeneous winds may exist for $Z/Z_{\\odot} = 0.1$, $1/3$ and $1$, respectively. Squares and circles represent the sdB stars analysed by Maxted et al. (\\cite{max01}) and Lisker et al. (\\cite{lisk05}), respectively. Filled symbols represent the sdB's with peculiar $\\rm H_{\\alpha}$ line profiles, which may indicate the presence of a weak wind.} \\end{figure}  In the figure it can be seen that the majority of sdB stars populate a region in the $T_{\\rm eff}$ - $\\log g$ diagram that is just below the line above which for $Z/Z_{\\odot} = 1$ coupled winds may exist. If multicomponent effects are neglected, the results predict weak winds with mass-loss rates $\\dot M \\la 10^{-12} M_{\\odot} / \\rm yr$. However, for such weak winds the momentum exchange between the metals, on the one hand, and hydrogen and helium, on the other, is not effective enough, because the densities in the wind are too low. Possible winds should be selective winds that lead to additional changes in the surface composition, which have not yet been taken into account in the diffusion calculations with and without mass-loss.  According to the present results coupled winds can exist for the most luminous sdB stars, if the metallicity is not too low. This would explain why only in these stars wind signatures have been detected. The predicted mass-loss rates are of the order of $10^{-11}$ to $10^{-10} M_{\\odot} / \\rm yr$. This agrees with the predictions of the mass-loss recipe of Vink \\& Cassisi (\\cite{vink02}), nevertheless the existence of these winds is still uncertain from the theoretical point of view. In Vink \\& Cassisi's calculations, the terminal velocity has been a free parameter, and their results are for $v_{\\infty} = v_{\\rm esc}$. The assumption of higher terminal velocities would lead to lower mass-loss rates. In the present calculations, the finite disk correction and the shadowing of the flux by the photospheric lines have not been taken into account. Especially the latter effect seems to be very important. As discussed in Sect. 6, both simplifications should lead to the mass-loss rates being overestimated and the terminal velocities being underestimated. Then for a star with given stellar parameters and chemical composition, the densities in the wind would be lower. Consequently decoupling of the metals from hydrogen and helium would occur at lower gravities than predicted, and the lines shown in Fig. 13 would be shifted to lower gravities. Frictional heating may have a similar effect. In addition, the presence of density inhomogenities in the wind, which are not taken into account in the present calculations, may possibly change the mass-loss rates. This effect of ``clumping\" is still investigated (Oskinova et al. \\cite{osk07}). For luminous stars there is some evidence that the mass-loss rates may need to be revised downwards (e.g. Bouret et al. \\cite{bou05}; Fullerton et al. \\cite{ful06}; Martins et al. \\cite{mar05}; Puls et al. \\cite{puls06}).  To confirm the existence of coupled winds in the most luminous sdB stars, more sophisticated calculations are necessary. Because the sdB stars have a variety of abundances, it is important to show which elements contribute primarily to the radiative force. As explained by Puls et al. (\\cite{puls00}), this should be the most abundant elements with the strongest lines in weak winds. According to the present results, the most important element is carbon for $T_{\\rm eff} = 25000$ and $30000 \\, \\rm K$. For higher effective temperatures the situation is less clear.  If during the post-EHB evolution indeed winds with mass-loss rates $\\dot M \\ga 10^{-11} M_{\\odot} / \\rm yr$ were initiated, the outer layers of the star will be removed rapidly. This should lead to the abundance anomalies gradually decreasing, so elements that are not detectable in sdB stars within the EHB band may reappear on the surface during the post-EHB evolution. This may explain the result of Edelmann et al. (\\cite{edel06}) that, in those sdB stars with $T_{\\rm eff} \\ga 32000 \\, \\rm K$ near the EHB band, no Si, Mg, and Al have been detected, whereas these elements are present in their more luminous counterparts with similar effective temperatures, which are in a post-EHB stage of evolution.  The results confirm the present picture of white-dwarf chemical evolution. The variety of compositions from hydrogen-rich to helium-carbon-oxygen-rich observed in pre-white dwarfs is due to the dredge up of processed material (Werner \\& Herwig \\cite{wer06}). Mass-loss rates between about $10^{-11}$ and $10^{-6} M_{\\odot} / \\rm yr$ are predicted for pre-white dwarfs with $M_{*} = 0.5 M_{\\odot}$, $30000 \\, \\rm K \\leq T_{\\rm eff} \\leq 50000 \\, \\rm K$, and $3.4 \\la \\log g \\la 5.0$, dependent on stellar parameters and metallicities. These  mass-loss rates are large enough to prevent diffusion, at least if the metallicity is not reduced by more than about a factor of ten in comparison to the solar value. No wind solution exists for white dwarfs on the cooling sequence with similar effective temperatures. Even for solar metallicity, the maximum possible radiative acceleration is too low by about one order of magnitude. Thus somewhere during the evolution with $T_{\\rm eff} > 50000 \\, \\rm K$, the mass-loss rates must decrease to very low values so that chemically homogeneous winds can no longer exist. The possibility that metallic winds still exist may explain the result of e.g. Good et al. (\\cite{good05}) that neither diffusion calculations that assume chemically homogeneous winds nor diffusion calculations that assume the absence of any mass-loss can explain the measured abundances in DA and DAO white dwarfs very well.  As discussed in Sect. 6.3, the existence of pure metallic winds with mass-loss rates $\\dot M \\la 10^{-16} M_{\\odot} / \\rm yr$ and with hydrostatic hydrogen and helium may be a promising scenario to explain the abundance anomalies of the metals in the various types of chemically peculiar stars considered in this paper. Then, however, the problem remains, why the abundances of helium in general are larger than predicted from diffusion calculations, which assume an equilibrium between gravitational settling and radiative levitation. Possibly the stellar atmospheres are not completely undisturbed after all.  Hot (pre-) white dwarfs already have a variety of helium (and metal-) abundances before the onset of gravitational settling is expected. A dependence on the star's history may also play some role in sdB stars. According to the analyses of Edelmann et al. (\\cite{edel03}) and Lisker et al. (\\cite{lisk05}) two distinct sequences of sdB stars seem to exist, which are characterised by an offset in the helium abundances. This phenomenon can hardly be explained with a single atmospheric effect."
    },
    "0808/0808.0524_arXiv.txt": {
        "abstract": "Current numerical studies suggest that the first protogalaxies formed a few stars at a time and were enriched only gradually by the first heavy elements.  However, these models do not resolve primordial supernova (SN) explosions or the mixing of their heavy elements with ambient gas, which could result in intervening, prompt generations of low-mass stars.  We present multiscale 1D models of Population III supernovae in cosmological minihalos that evolve the blast from its earliest stages as a free expansion. We find that if the star ionizes the halo, the ejecta strongly interacts with the dense shell swept up by the H II region, potentially cooling and fragmenting it into clumps that are gravitationally unstable to collapse.  If the star fails to ionize the halo, the explosion propagates metals out to 20 - 40 pc and then collapses, heavily enriching tens of thousands of solar masses of primordial gas, in contrast to previous models that suggest that such explosions 'fizzle'.  Rapid formation of low-mass stars trapped in the gravitational potential well of the halo is unavoidable in these circumstances.  Consequently, it is possible that far more stars were swept up into the first galaxies, at earlier times and with distinct chemical signatures, than in present models.  Upcoming measurements by the \\textit{ James Webb Space Telescope} (\\textit{JWST}) and \\textit{Atacama Large Millimeter Array} (\\textit{ALMA}) may discriminate between these two paradigms. ",
        "introduction": "Adaptive mesh refinement (AMR) and smooth particle hydrodynamics methods indicate that primordial stars form in the first cosmological dark matter halos to reach masses of a few 10$^5$ $\\Ms$ at $z \\sim$ 20 - 30, and that due to inefficient H$_2$ cooling they are likely very massive, 30 - 500 $\\Ms$ (\\cite[Abel et al. 2002]{abn02}, \\cite[Bromm et al. 2001]{bcl01}, \\cite[Nakamura \\& Umemura 2001]{nu01}). With surface temperatures of $\\sim$ 10$^5$ K and ionizing emissivity rates of 10$^{50}$ s$^{-1}$, these stars profoundly alter the halos that give birth to them, creating H II regions 2.5 - 5 kpc in radius and sweeping half of the baryons in the halo into a dense shell that grows to the virial radius of the halo by the end of the life of the star, as first pointed out by \\cite[Whalen et al. (2004)]{wan04} and \\cite[Kitayama et al. (2004)]{ket04}.  Recent work shows that a second lower-mass star can form in the relic H II region of its predecessor in the absence of a supernova (\\cite[Yoshida et al. 2007]{yet07}).  While single stars form consecutively in halos, gravitational mergers consolidate them into larger structures.  Numerical attempts to follow this process find that primordial supernovae expel their heavy elements into low-density voids, where star formation is not possible.  The metals return to their halos of origin on timescales of 50 - 100 Myr by accretion infall and mergers and, having been diluted by their expulsion into the IGM, are taken up into later generations of stars relatively slowly.  However, the large computational boxes required to follow the formation of halos from cosmological initial conditions prevents them from fully resolving the explosions.  Furthermore, the blasts themselves are not properly initialized, being set up as static bubbles of thermal energy rather than the free expansions that actually erupt through the atmosphere of the star.  If the explosion is initialized in an H II region and one is not concerned with the transport of metals or fine structure cooling, thermal 'bombs' yield an acceptable approximation to gas motion on kiloparsec scales.  However, this approach cannot capture the mixture of heavy elements with ambient gas or the subsequent formation of dense enriched clumps due to metal line cooling for two reasons. First, the thermal bubble does create large pressure gradients that accelerate gas outward into the surrounding medium, but because the ejecta has no initial momentum and dynamically couples to surrounding gas at low density ($\\lesssim$ 1 cm$^{-1}$), not all the metals are launched out into the halo.  Our numerical experiments demonstrate that unphysically large amounts of heavy elements remain at the origin of the coordinate grid when the explosion is implemented in this manner.  Second, the early evolution of the thermal pulse is different from that of a free expansion, which exhibits the formation of reverse shocks and contact discontinuities that would destabilize in 3D and cause mixing at fairly small radii.  Since dynamical instabilities in the blast that are mediated by line cooling are highly nonlinear, it is crucial to follow them from their earliest stages, something that cannot be accomplished with thermal bubbles.  If instead the progenitor fails to ionize the halo, densities at the center remain high, in excess of 10$^8$ cm$^{-3}$.  When the supernova is initialized with thermal energy in this environment, all the energy of the blast is radiated away before any of the surrounding gas can be driven outward.  Since the energy of the explosion is deposited as heat at the center of the grid, the temperature skyrockets to billions of degrees, instantly ionizing the gas collisionally. Bremsstrahlung and inverse Compton cooling time scales at these densities and temperatures are extremely short, leading earlier studies to erroneously conclude that such blasts 'fizzle'.  In reality, even though 90\\% of the energy of the explosion may be lost in the first 2 - 3 years, the momentum of the free expansion cannot be radiated away, guaranteeing some propagation of ejecta out into the halo. We present a new series of Population III supernova models in both neutral and ionized cosmological minihalos to address these difficulties.  \\vspace{-0.25in} ",
        "conclusions": "Our 1D survey of primordial supernova remnant energetics in cosmological halos strongly suggest that when metals and metal line cooling are included in the next generation of 3D models, dynamical instabilities will strongly mix the surrounding gas with heavy elements.  This may lead to a second, prompt generation of stars forming in either the enriched dense shell of the relic H II region or deeper inside a neutral halo.  If so, the first protogalaxies may have had far more stars than in current models.  Large infall rates in trapped explosions may also be efficient at fueling the growth of the compact remnant of less massive progenitors, possibly providing the seeds of the supermassive black holes found in most large galaxies today.  \\vspace{-0.15in}"
    },
    "0808/0808.2467_arXiv.txt": {
        "abstract": "The energy conditions play an important role in the description of some important properties of the Universe, including the current accelerating expansion phase and the possible recent phase of super-acceleration. In a recent work  we have provided a detailed study of the energy conditions for the recent past by deriving bounds from energy conditions and by making the confrontation of the bounds with supernovae data. Here, we extend and  update these results in two different ways.  First, by  carrying out a new statistical analysis for $q(z)$ estimates needed for the confrontation between the bounds and supernovae data. Second, by providing a new picture of the energy conditions fulfillment and violation in the light of the recently compiled \\emph{Union} set of $307$ type Ia supernovae and by using two different statistical approaches. ",
        "introduction": "In the study of the classical energy conditions~\\cite{EC-basics_refs} in cosmological context, an important viewpoint is the confrontation of their predictions with the observational data. Since the pioneering papers by Visser~\\cite{M_Visser1997} a number of articles have been published concerning this confrontation by using model-independent energy-conditions \\emph{integrated} bounds on the cosmological observables such as the distance modulus and lookback time~\\cite{Santos2006}~--~\\cite{Lima2008} (see also the related Refs.~\\cite{EnergCond_rel}). Energy conditions constraints on modified gravity models, such as the so-called $f(R)$--gravity, have also been investigated in Ref.~\\cite{Santiago2006} and more recently in Ref.~\\cite{SARC2007}.  In a recent work~\\cite{Lima2008}, we have shown that the fulfillment (or the violation) of these \\emph{integrated} bounds at a given redshift $z$ is not sufficient (nor necessary) to ensure the fulfillment (or the violation) of the energy conditions at $z$. This amount to saying that the local confrontation between the prediction of the \\emph{integrated} bounds and observational data is not sufficient to draw conclusions on the fulfillment (or violation) of the energy conditions at $z$. This crucial drawback in the confrontation between \\emph{integrated} bounds and cosmological data has been overcome in Ref.~\\cite{Lima2008}, where new \\emph{non-integrated} bounds have been derived, and confronted with type Ia supernovae (SNe Ia) data of the \\emph{gold}~\\cite{Riess2007} and \\emph{combined}~\\cite{Combined} samples.  In this letter, to proceed further with the investigation of the interrelation between energy conditions on scales relevant for cosmology and observational data, we extend and  update the results of Ref.~\\cite{Lima2008} in two different ways. First, carry out a new statistical analysis for $q(z)$ estimates necessary for the confrontation between the \\emph{non-integrated} bounds and supernovae data. Second, we give a new picture of the energy conditions fulfillment and violation for recent past ($z\\leq1 $) by using the recently compiled \\emph{Union} sample~\\cite{Kowalski2008} with $307$ type Ia supernovae along with the new as well as the previous~\\cite{Lima2008} statistical tools. ",
        "conclusions": "In a previous work~\\cite{Lima2008} we provided a picture of the violation and fulfillment of the energy conditions in the recent past by deriving  \\emph{non-integrated} bounds from energy conditions in terms of the deceleration parameter and the normalized Hubble function in the context of FLRW cosmology, and made the confrontation of the bounds with SNe Ia data through estimates of $q(z)$ from $1\\sigma - 3\\sigma$ confidence regions on the plane $E(z) - q(z)$ calculated with \\emph{gold}~\\cite{Riess2007} and \\emph{combined}~\\cite{Combined} samples.  Here, we have extended and updated the confrontation between the \\emph{non-integrated} bounds and supernovae data. First, by using the fact that, in the flat case, the \\textbf{NEC}, \\textbf{SEC}, and \\textbf{DEC} bounds do not dependent on $E(z)$, we have carried out a new statistical analysis in which the $q(z)$ estimates are obtained by marginalizing over the parameters $E(z)$ along with the $q'_l$'s of $q(z)$ [see Eq.~\\eqref{eq:q_z}]. Second, we have updated the previous work~\\cite{Lima2008} by providing a new picture [see Fig.~\\ref{qxz}(a) and [Fig.~\\ref{qxz}(b)] of the energy conditions fulfillment and violation from the recently compiled Union set of $307$ SNe Ia along with two different statistical tools.  On general grounds, our analyses indicate a possible recent phase of super-acceleration in which the \\textbf{NEC} is violated within $3\\sigma$ confidence level for $z \\lesssim 0.075$ [Fig.~\\ref{qxz}(a)] and $z\\lesssim 0.085$ [Fig.~\\ref{qxz}(b)], and that the \\textbf{DEC} is fulfilled with $3\\sigma$ in the redshift interval $(0, 0.765)$ [Fig.~\\ref{qxz}(a)] and $(0, 0.74)$ [Fig.~\\ref{qxz}(b)]. Regarding the \\textbf{SEC}, our analyses show that, for both statistical approaches employed, the best-fit curve of $q(z)$ crosses the \\textbf{SEC}--bound curve at $z \\backsimeq 0.51$, and that \\textbf{SEC} is violated with $3\\sigma$ within small low redshift intervals [Fig.~\\ref{qxz}(a) and [Fig.~\\ref{qxz}(b)].  Finally, an interesting fact that comes out of our \\textbf{SEC} analysis with $1\\sigma$ C.L., obtained by using the recent SNe Ia \\emph{Union} set, is that for the new $q(z)$ estimate [calculated by marginalizing over $E(z)$] the deceleration to acceleration transition expansion phase of the universe took place in the redshift interval ($\\simeq 0.4, \\simeq 0.64\\,$)."
    },
    "0808/0808.0462_arXiv.txt": {
        "abstract": "We consider accretion onto newborn black holes (BHs) following the collapse of rotating massive stellar cores, at the threshold where a centrifugally supported disk gives way to nearly radial inflow for low angular momentum. For realistic initial conditions taken from pre-supernova (pre-SN) evolution calculations, the densities and temperatures involved require the use of a detailed equation of state and neutrino cooling processes, as well as a qualitative consideration of the effects of general relativity. Through two-dimensional dynamical calculations we show how the energy release is affected by the rotation rate and the strength of angular momentum transport, giving rise to qualitatively different solutions in limits of high and low angular momentum, each being capable of powering a gamma-ray burst (GRB). We explore the likelihood of producing Fe-group elements in the two regimes and suggest that while large and massive centrifugally supported disks are capable of driving strong outflows with a possible SN-like signature, quasi-radial flows lack such a feature and may produce a GRB without such an accompanying feature, as seen in GRB060505. ",
        "introduction": "\\label{sec:intro}  Gamma-ray bursts (GRBs) have remained as an outstanding problem in astrophysics since their discovery in 1967 \\citep{kso73}.  Major advances in their understanding came through observational breakthroughs following BATSE, which provided an extensive all-sky catalog \\citep{fm95} and the identification of counterparts at lower energies at high redshift \\citep{m97} through Beppo Sax \\citep{vkw00} and Swift \\citep{gcn07}. Isotropic equivalent energies range from $10^{49}$~erg~s$^{-1}$ to $10^{52}$~erg~s$^{-1}$, coming from sources at redshifts as high as $z\\simeq 6$ that are randomly distributed over the sky \\citep{m92}. GRBs are commonly divided into short (SGRBs, less than 2 seconds in duration, typically at redshift $z\\simeq1$ and frequently in host galaxies with low star formation rates (SFRs)) and long (LGRBs, longer than 2 seconds, residing at $z\\simeq1-5$ in star forming galaxies) events. An additional distinction is that those of the short variety tend to have harder spectra, hence the short-hard/long-soft denomination \\citep{d96}. Recently there appears to be a substantial and, in some respects, puzzling overlap between the two populations, and it has been suggested that a third group or even a new classification is necessary \\citep{g06,gy06,dv06nat,f06}.  The energy release, short duration, and variability time scales associated with GRBs favor accretion onto compact objects, or magnetically powered events from such objects, as their ultimate source. Many models have been suggested over the years, a number of which were discarded once the global energy budget was fixed through the resolution of the distance scale to the sources. Among the early proposals still considered are compact binary mergers \\citep{ls74,ls76,bp86,bp91,e89,n92}, magnetars \\citep{u92,t94,mr97} and the collapse of massive stellar cores \\citep{w93}, see \\citep{m02,piran04,n07,lrr07} for reviews.  In general, accretion onto a central object can proceed and release gravitational binding energy efficiently if the gas can cool or otherwise lose its internal energy e.g., through advection, and thus move down in the potential well \\citep{s64,z64}. At the usual astrophysical scales, this occurs through photons, and produces phenomena ranging from active galactic nuclei (AGN) to low mass X-ray binaries (LMXBs). If the accretion rate is too high though, the fluid may become so optically thick to the photons that they become trapped and are unable to cool the gas. At even higher accretion rates, a different mechanism enters the picture, namely, neutrino cooling. This regime is termed ``hypercritical\", since it is far above the usual Eddington limit for photons, and is relevant, for example, in post-supernova cores and the associated fall back, as was probably the case in SN1987A \\citep{c89,hc91}. It is this phenomenon that is now thought possible of generically powering GRBs. The question at the forefront of attention has been how to achieve these conditions in an astrophysical setting.  There has been mounting evidence in the past few years linking LGRBs at low redshift with type Ic SNe, with spatial and temporal coincidences for a number of events \\citep{wb06,kaneko08}. \\citet{g98} linked for the first time a LGRB with a SN (GRB980425 with SN1998bw). After 27 days the optical spectrum of GRB980425 resembled that of a H-deficient type Ic SN. To date there have been at least 6 other LGRB/SN connections: XRF020903 \\citep{sod05}; GRB021211 with SN2002lt \\citep{dv03}; GRB030329 with SN2003dh (whose optical spectrum afterglow was strikingly similar to that of SN1998bw \\citep{s03}); GRB031203 with SN2003lw \\citep{m04,rr05}; GRB050525A with SN2005nc \\citep{dv06apj}; and GRB060218 with SN2006aj \\citep{c06}. In addition, LGRBs in general are associated with low metallicity star-forming regions in their host galaxies \\citep{p04,g05,sol05,fruchter06,f06} suggesting a link to massive stars \\citep{bkd02,wb06}.  The observational association with SNe has given support to the idea that LGRBs are produced in the collapse of massive stellar cores. \\citet{mw99} developed this in great detail in the collapsar scenario, assuming a black hole (BH) is formed at the center and evaluating the potential to produce a GRB. The infalling stellar material that is not ejected in the usual SN forms a massive, dense and hot accretion disk which releases the necessary energy to power the burst. They distinguished between Type I and Type II collapsars: the former entails the direct collapse of the iron core to a BH, while in the latter initially a proto-neutron star lies at the center of the star, and later collapses to a BH once it has accreted a sufficient amount of mass. A GRB may be produced even if there is no BH at the center in a magnetically dominated explosion during the proto-neutron star phase of the collapse \\citep{dessart08}. Here, we will focus on Type I collapsars. The absence of H lines in the observed GRB/SN spectra implies that the progenitor has somehow lost its envelope, perhaps through interaction with a binary companion and/or strong winds, and is thus essentially a Wolf-Rayet (WR) star \\citep{izzard04,cantiello07}.  In the standard version of the collapsar, the GRB/SN link is a quite natural consequence, and observed associations have thus provided a strong motivation in this sense to study it further. Typical LGRBs have $z\\simeq1-4$, which is too far for a SN counterpart to be detected, and indeed, the cases in which there is such a correspondence all lie at fairly low redshift. Two recent events however, GRB060505 and GRB 060614 \\citep{f06,g06,gy06,dv06nat}, have pointed perhaps to a different situation in which the GRB is {\\em not} accompanied by a SN. Their low redshift allowed deep searches which convincingly ruled out the presence of an underlying type Ic stellar explosion of the kind seen in SN1998bw and SN2003dh, and particularly, the host of GRB060505 is a star-forming galaxy, similar to that of typical LGRBs \\citep{f06}. We explore here how the initial angular momentum distribution in the star, and possibly the vigor of angular momentum transport within the centrifugally supported accretion disk, may furnish a key ingredient to understand this behavior.  An important point is that despite considerable effort (see \\citet{wh06}), the stellar rotation rate in pre-SN cores is not fully determined. It can sensitively depend on evolutionary details, such as mass loss on the main sequence and the stellar magnetic field, which both lead to important spin down in later stages \\citep{spruit02, hws05}. Binary interactions may also affect the rotation rate through tidal interactions \\citep{detmers08}. We note that \\citet{yoon05}, and \\citet{wh06}, have identified a channel for massive stars in which mass and angular momentum losses are greatly reduced by complete mixing in the main sequence and the consequent absence of a giant phase.  What we do know is that the specific angular momentum, $J$, needed to form a disk around the BH is at least the value of the innermost stable circular orbit $R_{isco}$ ($R_{isco}=3r_{\\rm g}$, where $r_{\\rm g}=2GM_{\\rm BH}/c^{2}$ is the Schwarzschild radius) times the velocity of the particles (which are very close to the speed of light). So, a first estimate for the critical J would be $J_{crit}=R_{isco}c=3r_{\\rm g}c=6GM/c\\sim10^{16}$~cm$^{2}$~s$^{-1}$ for a one solar mass BH (the true critical value is $J_{\\rm crit}\\simeq2r_{\\rm g} c$, see Sextion~3.1).  Most studies of collapsars and neutrino-dominated accretion flows \\citep{mw99,pwf99,h00,npk01,pb03,pmab03,lrrp05} have considered angular momentum distributions that are well above this limit and essentially guarantee the formation of a centrifugally supported accretion disk (in stellar core collapse calculations as well as disk evolution studies). In practice, there is also something akin to a maximum value of angular momentum for the collapsar model to work, since for very rapid rotation the accretion disk forms at large radii and the binding energy cannot be effectively dissipated as neutrinos \\citep{mw99,lrr06}.  Clearly in the limit of slow rotation the accretion disk will form very near the BH, and so general relativity (GR) will play an important role in its evolution.  Our general objective is to explore the morphology, energy release, and observable signatures of hypercritical accretion flows from massive stellar core collapse for a range of angular momentum values covering the transition from centrifugally supported disks to near radial inflow (lying below those usually considered in collapsar calculations) in order to better understand the possible production of LGRBs as well as their putative progenitors. We improve upon previous studies \\citep{mw99,pmab03,lrr06} in significant ways by considering more detailed thermodynamics in the equation of state, an improved treatment of neutrinos, and through realistic initial conditions taken from evolutionary calculations of \\citet{wh06}.  We first describe the input physics and numerical setup in Section~\\ref{sec:input}, followed by our results in Section~\\ref{sec:results}. Prospects for GRB production and the link between SN and GRBs are given in Section~\\ref{sec:conc}. \\\\ ",
        "conclusions": "\\label{sec:conc}  \\subsection{Limitations and comparison to other work}  As with all numerical work, the choices made in carrying out the simulations reflect intentions and biases, and the current investigation lacks in several aspects. For example, being a two-dimensional calculation, self-gravity has been considered only in an approximate manner, and stability issues that relate to this or are intrinsically three-dimensional (such as spiral arms) are ignored.  Magnetic fields have not been included and are in all likelihood important in several aspects, two of which deserve special mention. The first is simply that magnetic fields are thought to be at the origin of the magneto rotational instability (MRI) \\citep{bh98} responsible for angular momentum transport and dissipation, which we have considered through the $\\alpha$ parameter. The MRI likely operates in collapsars, although the details of how it does so are likely to be different than what we can infer about it from theory and observation of accretion disks in CVs and X-ray binaries in a more leisurely state of affairs. The second point is related to the importance of magnetic pressure in possible outflows emanating from a disk. Detailed GR MHD calculations \\citep{devilliers05} show that a flow starting from an equilibrium torus develops the MRI and produces an inner dense disk, a funnel wall, and a corona around the central BH. Gas pressure dominates largely in the first, so our hydrodynamical solution to the inner disk is likely to be a realistic approximation, while magnetic pressure is relatively more important in the other two, and being polar structures they are important for the driving of outflows and energy release in the context of GRBs \\citep{mckinney07}. For very rapidly rotating BHs, \\citet{krolik05} have shown that the magnetic field can actually have a dynamical effect on the mass accretion rate, suppressing it in the inner regions (in their case the Kerr parameter was $a=0.998$), although a full evaluation of this deserves further study.  GR, approximated in this work simply by use of the Paczy\\'{n}ski--Wiita potential, is likely to play a role as well. The location of the horizon and last stable orbit are functions of the rotation of the BH, lying at $r_{\\rm H}=2GM/c^{2},GM/c^{2}$ and $r_{\\rm ISCO}=6GM/c^{2},3GM/c^{2}$ for $a=0,1$ respectively, and this is important since most of the energy is released at small radii. For our purposes, however, they are not likely to be crucial since the relevant radii shift appreciably only for very rapidly rotating holes (with $a$ above 0.7 or so), and the progenitor stars do not easily reach such high values (model 16TI has $a=0.44$). A secondary aspect of GR is related to the efficiency for $\\nu\\overline\\nu$ annihilation and the corresponding energy release. Some energy is directly lost to the BH because of the strong gravitational field, while focusing increases the annihilation efficiency. Recently \\citet{birkl07} have carefully computed the energy deposition rates and efficiencies for annihilation in the Kerr metric for various BH spin rates and flow geometries. They find that the power output is affected by approximately up to a factor of 2, depending on whether the disk is geometrically thin or thick (or even for spherical configurations of the neutrinosphere). Thus the general energy scale is affected, but not crucially so.  On the other hand, our effort has been directed here at improving the thermodynamical treatment and that of neutrinos, as well as in considering realistic initial conditions derived from stellar evolution calculations. The choice of inner boundary means that the dense, inner disk producing most of the luminosity is well resolved and treated appropriately. A final point that deserves improvement is clearly the distribution of angular momentum, since having a constant value at all radii is not realistic. We have chosen it for now to gauge the effect of other changes when compared to previous work, and provide a guide in further investigation, which will fully explore situations in which the radial part is a characteristic function of radius (D. Lopez-Camara et al. 2008, in preparation).  With this in mind, we may consider the implications for GRB production and GRB/SN associations following core collapse and prompt BH formation from this set of calculations.  \\subsection{Global energetics}  The fact that we have placed the inner numerical boundary fairly close to the BH (at $r_{\\rm in}=20$~km) allows us to directly and more realistically compute the energy release (obviously at the cost of a shorter simulation in time). This is similar to what was done by \\citet{lrr06} and so we can directly compare the two sets of results. The two most important differences between the two studies are in the use of a better equation of state and neutrino treatment here, and in the use of initial models taken from stellar evolution calculations. The neutrino luminosities are higher in the present study, partly because of the more realistic physics employed, but also because of the initial conditions, which directly result in higher accretion rates and densities.  Since the cooling rate from captures scales with the density, this fact alone will raise the luminosity. The rates are quite sensitive to the temperature, but the rise in the postshock material is mostly due to the infall kinetic energy per unit mass and is fixed by the potential well of the central BH.  It is clear as well, from the results presented for simplified cases in the adiabatic and isothermal limits, and from the computation of the Kelvin--Helmholtz timescale, that correctly accounting for the energy sinks is a crucial ingredient if one wishes to estimate the global energy release and flow properties.  Finally, the emergent neutrino spectrum is a function of the angular momentum of the infalling gas, through the selection of the dominant cooling mechanism. This is admittedly a difficult issue to resolve observationally, but may impact upon the energy deposition rate.  The two most important variables that determine the global morphology are the rotation rate, quantified here through the distribution of specific angular momentum, and the strength of angular momentum transport, parameterized through the prescription of Shakura \\& Sunyaev. It is interesting to note that for high angular momentum (larger than the threshold for disk formation), viscous dissipation is responsible for the production of large-scale flows, namely by producing an outward moving shock after a certain delay (a few tenths of a second). The effect is to perturb the flow into the inner disk and decrease the luminosity. The impact on the flow at large radii remains to be explored in greater datail due to the limitation of our outer boundary.  \\subsection{Implications for GRB production}  The compression for low angular momentum cases releases a large amount of energy, not more than 1 order of magnitude smaller than for the case with a disk, and we believe this to be an important issue: slowly rotating models are in principle capable of releasing an amount of energy that could produce GRBs, although admittedly perhaps only at the faint end of the luminosity distribution. For a 1\\% efficiency of conversion into pairs through $\\nu\\overline{\\nu}$ annihilation at $L_{\\nu}=10^{53}$~erg~s$^{-1}$, the models shown here can produce annihilation luminosities of about $10^{48}-10^{49}$~erg~s$^{-1}$, depending on the value of $J_{0}$. Performing realistic runs for longer times is a priority, but clearly over long timescales, having slow rotation is not necessarily a handicap regarding the energy release. The simple reason is that, as known for a long time, and first quantified in this scenario by \\citet{lrr06}, it is most efficient in terms of energy extraction to have the material release its energy as close as possible to the BH, or equivalentely, as deep as possible in the potential well, but not too close so that it is still shocked by the presence of a centrifugal barrier. Magnetic mechanisms may also power fully or partially a potential burst, and we note here only that the internal energies are high enough, as in the most common hypercritical accretion scenarios, to do so if a fraction of this is transferred to a magnetic field (at a level of 10\\% of equipartition).  The obvious implication in terms of the type of progenitor one can consider for GRBs is that single stars at the threshold for disk formation through centrifugal support are capable of giving a GRB. Having a substantial amount of rotation will obviously guarantee an accretion disk, but perhaps not all cases require such special conditions (as for example, torquing of the pre-SN star by a binary companion). The isotropic equivalent energy of GRB060505 is $\\log[E_{\\gamma, \\rm iso}]\\simeq 49.5$, clearly within the range obtainable with the models presented here even at low rotation rates.  \\subsection{Nucleosynthesis in the outflow and GRB/SN association}  Centrifugally supported collapsar disks are expected to produce strong winds \\citep{mw99}, driven quite generically by a combination of viscous, neutrino, and magnetic effects (the first two of which we have explicitly included here). \\citet{pth04} in particular, computed the expected nucleosynthesis of $^{56}$Ni in these outflows, using the steady state disk models of \\citet{pwf99} as an initial condition and considering various mass accretion and dissipation rates. The essential feature of these models is the computation of the change in entropy per baryon, $s_{b}/k$ in the outflow, starting from the midplane of the accretion disk and reaching a large distance at which the velocity approaches an asymptotic value $v_{\\rm w}/c \\simeq 0.1-0.2$. From our hydrodynamical determination of the disk structure using an improved equation of state and thermodynamics, we computed the equivalent entropies in the disk and applied the formalism of \\citet{pth04} to compute the total change in $s_{b}/k$ (assuming the wind reaches similar velocities, unresolved in our simulations spanning 0.5~s). The mass fraction converted to $^{56}$Ni, $X_{\\rm Ni}$, then depends sensitively on $s_{b}/k$ and the ratio $\\beta=\\dot{M}_{\\rm o}/v_{\\rm w}^{3}$, where $\\dot{M}_{\\rm o}$ is the mass outflow rate and $v_{\\rm w}$ is the terminal wind velocity. Once the outflow is launched we can obtain an estimate of $\\dot{M}_{\\rm o}$ directly from the momentum field in the simulation, and $v_{\\rm w}/c$ is taken as $\\simeq 0.2$. We find that for $J_{0} > J_{\\rm crit}$, and independently of the strength of angular momentum transport, substantial Ni synthesis occurs (X$(^{56}\\mbox{Ni})\\approx 0.5$) in an outflow with $\\dot{M}\\approx 0.3$~M$_{\\odot}$~s$^{-1}$. This would imply a total ejection of $\\simeq M_{\\odot}$ in $^{56}$Ni in a GRB lasting 10~seconds, in the right range to account for the outflows seen, e.g., in SN2003dh. The relatively narrow range in $J_{0}$ spanned in our calculations does not allow for a good extrapolation to higher rotation rates, but clearly the amount of mass involved is significant, and would likely grow substantially at higher $J_{0}$ where more masive and more radially extended disks are expected (\\citet{pth04} estimate several $M_{\\odot}$ for the massive, centrifugally supported disks taken from \\citet{pwf99}). A configuration in which no significant outflows occur because the rotation rate is too low ($J \\le J_{\\rm crit}$) would not be expected in this scenario to produce any significant ejection of radioactive elements capable of producing a SN-like signature. It could, as shown here, still release enough energy through neutrinos or magnetic mechanisms to power a classical GRB.  The result, if it were to occur at low redshift, could possibly be an event resembling GRB060505."
    },
    "0808/0808.2651_arXiv.txt": {
        "abstract": "We present an analysis of the chemical and ionization conditions in a sample of $100$ weak {\\MgII} absorbers identified in the VLT/UVES archive of quasar spectra. In addition to {\\MgII}, we present equivalent width and column density measurements of other low ionization species such as {\\MgI}, {\\FeII}, {\\AlII}, {\\CII}, {\\SiII} and also {\\AlIII}. We find that the column densities of {\\CII}  and {\\SiII} are strongly correlated with the column density of {\\MgII}, with minimal scatter in the relationships. The column densities of {\\FeII} exhibit an appreciable scatter when compared with the column density of {\\MgII}, with some fraction of clouds having $N(\\FeII) \\sim N(\\MgII)$, in which case the density is constrained to n$_{\\H} > 0.05$~{\\cc}. Other clouds in which $N(\\FeII) << N(\\MgII)$ have much lower densities. From ionization models, we infer that the metallicity in a significant fraction of weak {\\MgII} clouds is constrained to values of solar or higher, if they are sub-Lyman limit systems. Based on the observed constraints, we hypothesize that weak {\\MgII} absorbers are predominantly tracing two different astrophysical processes/structures.  A significant population of weak {\\MgII} clouds, those in which $N(\\FeII) << N(\\MgII)$, identified at both low ($z \\sim 1$) and high ($z \\sim 2$) redshift, are likely to be tracing gas in the extended halos of galaxies, analogous to the Galactic high velocity clouds. These absorbers might correspond to $\\alpha$-enhanced interstellar gas expelled from star-forming galaxies, in correlated supernova events. The $N(\\MgII)$ and $N(\\FeII)/N(\\MgII)$ in such clouds are also closely comparable to those measured for the high velocity components in strong {\\MgII} systems. An evolution is found in $N(\\FeII)/N(\\MgII)$ from $z = 2.4$ to $z = 0.4$, with an absence of weak {\\MgII} clouds with $N(\\FeII) \\sim N(\\MgII)$ at high-$z$. The $N(\\FeII) \\sim N(\\MgII)$ clouds, which are prevalent at lower redshifts ($z < 1.5$), must be tracing Type Ia enriched gas in small, high metallicity pockets in dwarf galaxies, tidal debris, or other intergalactic structures. ",
        "introduction": "\\label{sec:1} The {\\HI} gas directly associated with galaxies that intercept the line of sight to background quasars appears as optically thick Lyman Limit systems in the quasar spectrum. The prominent metal lines associated with these intervening absorbers are typically observed to be kinematically broad ($\\Delta v \\sim 100 - 400$~{\\kms}), strong, and often saturated \\citep[e.g.,][]{ss92,cwcvogt01} . Studying the properties of a large population of such strong {\\MgII} absorbers is a technique used for constraining the evolution of metals in the interstellar media, gaseous halos and coronae of galaxies over a large history of the universe \\citep{lanzetta87,cwc96}. Apparently distinct from these strong {\\MgII} absorbers are the population of quasar absorption line systems in which the low ionization metal lines are weak.  These systems are separated from the strong ones based on the standard definition of the rest-frame equivalent width of $\\MgII~\\lambda 2796$~{\\AA} line being $W_r(2796) < 0.3$~{\\AA}. This is not a firm criterion for division, but has been followed as a convention on the following basis. The survey of \\citet{ss92}, which identified a large population of strong {\\MgII} absorbers used a sample of intermediate resolution spectra ($\\Delta \\lambda \\sim 5$~{\\AA}) which had an equivalent width threshold of $\\sim 0.3$~{\\AA}. Later surveys of higher sensitivity and spectral resolution found that the equivalent width distribution of {\\MgII} systems at $z \\sim 1$ increases steeply for $W_r(2796) < 0.3$~{\\AA} such that $\\sim 67\\%$ of all {\\MgII} absorbers (down to $0.02$~{\\AA}) from that epoch are in fact weak \\citep{weak1, anand07}. It later became clear that such an empirical basis for the classification of {\\MgII} systems into {\\it strong} and {\\it weak} does bear some physical significance in that the two classes might be tracing two or more different populations of objects \\citep{weak1, weak2, charlton03}.  The class of weak {\\MgII} quasar absorption systems have several remarkable properties that are unique. To begin with, unlike the strong systems, the weak {\\MgII} systems are optically thin in neutral hydrogen and produce metal lines that are narrow [$b$(Mg) $\\sim 4$~{\\kms}] and often unsaturated \\citep{weak1}. If weak {\\MgII} absorbers are sub-Lyman limit systems with $10^{15.8} < N(\\HI) < 10^{16.8}$~{\\cmsq}, they would account for a significant fraction ($> 25\\%$) of the high column density regime of the {\\Lya} forest \\citep{weak2}. Surveys of quasar fields to identify host galaxies have not often found weak {\\MgII} systems at close impact parameters (physical distance, D $< 30$~Kpc) of luminous star forming galaxies ($L > 0.05$~$L_*$) \\citep{cwc98, cwc05, milni06}. This is a surprising result particularly in light of the fact that, in a substantial number of weak systems, the metallicity of the low ionization gas where the {\\MgII} absorption arises is constrained to values greater than 0.1Z$_\\odot$.  In some cases the best constraints require metallicities that are between Z$_\\odot$ and 10Z$_\\odot$ \\citep{weak2, charlton03, misawa07}. Thus, even though they have {\\HI} column densities that are $\\sim 4$ orders of magnitude smaller than DLAs, weak {\\MgII} absorbers are produced in gas clouds with metallicities that are 0.5 - 2 dex higher than the average metallicity of DLA absorbers.  The astrophysical systems associated with weak {\\MgII} absorbers have not been identified yet. Several possibilities exist, which partly account for the observed statistical and physical properties of weak {\\MgII} systems. Examples include extragalactic high velocity clouds \\citep{anand07}, dwarf galaxies \\citep{lynch06}, gas clouds expelled in  super winds from dwarfs \\citep[e.g.][]{zonak04, stocke04, keeney06} and/or massive starburst galaxies and metal enriched gas in intergalactic star clusters \\citep{weak2}. Recently, a number of authors have reported the detection of several high metallicity ($> 0.1$~Z$_\\odot$) gas clouds that are residing in the intergalactic medium, both in the local universe \\citep{aracil06, tripp06} and at high redshift [$z > 2$, \\citet{simcoe04, schaye07}]. The observed column densities of {\\MgII}, {\\CII}, {\\SiII} and {\\FeII} in these gas clouds, and the metallicities inferred for them are comparable to several weak {\\MgII} absorbers studied so far \\citep{weak2,lynch07,misawa07}.  Photoionization models of specific weak {\\MgII} systems have shown that they possess a two phase structure. The low ionization gas, traced by such species as {\\MgII}, {\\FeII}, {\\CII}, {\\SiII}, etc. has a gas density of n$_{\\H} > 0.1$~{\\cc} that is roughly 2 - 3 orders of magnitude higher than the density of the associated high ionization gas traced by {\\CIV} lines \\citep{charlton03, lynch07, misawa07}. The column density measured for weak {\\MgII} lines is typically $N(\\MgII) \\sim 10^{12} - 10^{13}$ {\\cmsq}. For such relatively small values, the high number density of ions derived from the photoionization modeling constraints the thickness of the low ionization gas to $\\sim 10$~pc. The weak {\\MgII} population occupies a significant volume of the universe with a cross-section similar to the absorption cross section of luminous ($L_*$) galaxies \\citep{weak1, anand07}. Thus, given their small thickness, if their gas clouds had a spherical geometry then they would be a million times more numerous than luminous galaxies at $z \\sim 1$ in order to reproduce the observed cross-section on the sky. However, an analysis comparing the relative incidence of high and low ionization gas in a sample of weak {\\MgII} and {\\CIV} systems at low redshift ($z < 0.3$) favors a filamentary or sheet-like configuration for the absorber's physical geometry \\citep{milni06}, instead of millions of individual, {\\it spherical} low ionization {\\MgII} clouds of $\\sim 10$~pc size, embedded in a $\\sim $~kpc higher ionization halo, traced by {\\CIV}.  In addition, \\citet{lynch06} and \\citet{anand07} discovered an evolution in the redshift number density ($dN/dz$) of weak {\\MgII} absorbers over the interval $0.4 < z < 2.4$. The $dN/dz$ was found to peak at $z = 1.2$, and subsequently decline (from a no-evolution trend,) towards higher redshift. The equivalent width distribution function was also found to be different between low and high redshift. At $z \\sim 1$, the equivalent width distribution of weak {\\MgII} is significantly higher than an extrapolation of the exponential distribution for strong {\\MgII} absorbers. In contrast, at $z \\sim 2$, the equivalent width distribution of weak {\\MgII} clouds is only slightly in excess of an extrapolation of the strong {\\MgII} distribution. In the context of these observed changes in the absorber statistics between $z \\sim 1$ and $z \\sim 2$, it becomes important to investigate if the changes are reflective of an underlying change in the absorbers' physical or chemical properties. Such an investigation may yield valuable clues into the physical nature of the kind of astrophysical processes/structures that produce weak {\\MgII} absorption at low and high redshifts. The large sample size considered here, provides the scope for such an analysis.  In this paper we present constraints on the chemical abundances, metallicity and ionization conditions for a sample of 100 weak {\\MgII} absorbers identified in spectra that were extracted from the VLT/UVES archive. We compare the observed line properties between the various chemical elements in order to constrain the metallicity, density and line-of-sight thickness of the low ionization gas. Previous studies have focused on individual weak {\\MgII} systems, whereas here we have attempted to derive the range of properties for a large ensemble of weak {\\MgII} systems, which have only recently been discovered (Narayanan {\\etal} 2007). In \\S~\\ref{sec:2} and \\S~\\ref{sec:3}, we explain the measurement of line parameters and the comparison between the prominent metal lines in these systems. In \\S~\\ref{sec:4}, we present the Cloudy photoionization constraints on the densities and metallicities of these absorbers. An observed trend in the {\\FeII} to {\\MgII} ratio with redshift is presented in \\S~\\ref{sec:5}, and a comparison of weak {\\MgII} clouds with the high velocity subsystems in a sample of strong {\\MgII} systems in \\S~\\ref{sec:6}. We conclude the paper with a summary of the significant results (in \\S~\\ref{sec:7}) and a detailed discussion (\\S~\\ref{sec:8}) on the nature of the gaseous structures selected by weak {\\MgII} absorption. ",
        "conclusions": ""
    },
    "0808/0808.1652_arXiv.txt": {
        "abstract": "We discuss some of the recent developments in the quark star physics along with the consequences of possible hadron to quark phase transition at high density scenario of neutron stars and their implications on the Astroparticle Physics. ",
        "introduction": "The underlying quark structure of hadrons suggets the possibility of a quark-hadron phase transition at extreme conditions of high tempereature and/or density. Since a compact object like neutron star (NS) provide the natural scenario of high density, the suggestion for the existence of quark core inside such massive compact objects was put forward by Ivanenko and Kurdgelaidze (1969). The existence of 3-flavour quark star or strange star (QS), made up of $u$, $d$ and $s$ quarks, was suggested by Itoh (1970). In general, the two flavour matter can not be more stable compared to nucleonic matter. The presence of $s$ quarks, along with $u$ and $d$ quarks, provides an additional fermi well which would result in the lowering of the energy of the 3- flavour quark matter or strange quark matter (SQM) compared to 2- flavour quark matter. Since $s$ quark has larger mass, the situation would be more favourable at higher densities. Such possibilities were recognized by Bodmer (1971) and led Witten (1984) to conjecture that SQM may be the true ground state of strongly interacting matter at high densities. Since then, this field has become an active area of research. The possibilty of a quark -hadron phase transition in the early universe (high temperature and small chemical potential) scenario and attempts to produce such a matter in the laboratory through heavy ion collision reactions (Alam, Sinha \\& Raha, 1996) using accelerators have made this field more interesting as well as intriguing.  One of the major difficulty in the theoretical calculations in these areas of research is the fact that the Quantum Chrmodynamics (QCD) perturbative series shows poor convergence except for very small coupling at very high temperatures ($ \\alpha_s < 0.5$, $T \\sim 10^{5} T_c$). The perturbative treatment at high density also fails. The lattice calculations are still not reliable in the high density regime applicable to NS. Hence the high density systems are studied using the QCD inspired phenomenological models. Numerous model calculations have predicted a stable quark matter system within finite parameter ranges (Ghosh \\& Sahu, 1993).  Presently, the major technical advancements in both ground as well as satellite based observations are producing huge amount of data. There exist large number of observational data on mass-radii of NS. But, except very soft equations of state (EOS) most of the other EOS can explain these static properties within the error bars (Schaffner-Bielich, 2007). So it is necessary to study different dynamical features of quark matter in its various forms so as to predict unambiguous signatures of QS and validate them using the available observational data.  In this article, we will review the present status of theoretical understandings of the role of quark matter in the development of astroparticle physics, more specifically, the physics of compact objets with quark core (hybrid stars or HS) and QS in the light of various observational studies along with the efforts for the search of strange matter in cosmic rays. We divide the article in few sections. Section 2 deals with symmetry structures of quark matter at high densities. Compact stars' pulsating modes and their importance is discussed in section 3. QCD phase transition and its consequences are discussed in section 4. Section 5 deals with Strangelets. Summary and discussions are given at the end. ",
        "conclusions": "QCD transition inside the NSs and possible formation of QSs is one of the most interesting areas of research. We have discussed the different facets of quark star physics along with the associated observational scenarios. Though an observational signature of phase transition and SQM in nature is still elusive, the study of these phenomena is extremely important for a better understanding of the physics of strongly interacting matter. Moreover, there are many questions remain to be answered before we can make any conclusions."
    },
    "0808/0808.3461_arXiv.txt": {
        "abstract": "{ We present a CCD BV photometry of the possible binary open star cluster NGC 7031/NGC 7086. The aim is to confirm or disprove their common nature on the grounds of their age and distance. An age of 224 $\\pm$ 25 Myr  and distance 831 $\\pm$ 72 pc was determined for NGC 7031 and 178 $\\pm$ 25 Myr, 955 $\\pm$ 84 pc for NGC 7086, respectively. Based on these differences in age and distance we conclude that the two clusters are most likely not formed together from one and the same Giant Molecular Cloud and thus are not a true binary cluster. } ",
        "introduction": "A binary open star cluster is an object consisting  of two open clusters. They can be basically described as: (i) binary physical systems with common origin formed together from one and the same Giant Molecular Cloud (GMC), having comparable ages and chemical compositions (we call this a true binary cluster); (ii) binary physical systems arising from clusters formed in different parts of the Galaxy and forming a pair through gravitational capture. Cluster pairs formed through tidal  capture have different ages and/or \\\\chemical composition.  The existence of star cluster pairs in our neighbouring galaxies, the Magellanic Clouds, was confirmed by several authors (e.g. Bhatia \\& Hatzidimitriou 1988; Bica et al.1992; Vallenari et al. 1998; Dieball \\& Grebel 2000; de Oliveira et al. 2000). Dieball 2002 proposed a catalogue of binary and multiple cluster candidates in the Large Magellanic Cloud with 473 members. There are more than 1600 open clusters in our Galaxy but only one well established double or binary cluster namely $h + \\chi$ Persei  (see e.g. Uribe et al. 2002, Kharchenko et al. 2005). Our Galaxy seems to show a lack of binary or multiple clusters as compared to the Magellanic Clouds. It is still an open question why such objects do not exist in our Galaxy. The reasons could be that they have already merged or dissipated, or that they simply do not form. Several lists of binary open cluster candidates have been proposed and studied by various authors: Lyng\\aa{} \\& Wramdemark 1984; Pavlovskaya \\& Filippova 1989; Tignanelli et al. 1990; Subramaniam et al. 1995; Loktin 1997; Muminov 2000. One of the most complete and well studied lists is the one of Subramaniam et al. 1995 with 18 candidate pairs, including clusters NGC 7031 and NGC 7086 studied here. Basic parameters of NGC 7031 and NGC 7086 as given in Dias et al. 2002 and WEBDA\\footnote{http://www.univie.ac.at/webda} database are presented in Table 1.  \\begin{table} \\caption{Basic parameters of the clusters NGC 7031 and NGC 7086.} \\begin{tabular}[h]{ccc} \\hline\\noalign {\\smallskip} Parameter            \t& NGC 7031   & NGC 7086    \\\\ \\hline\\noalign {\\smallskip} R.A.(2000)           \t& 21:07:12\t& 21:30:27    \\\\ Decl.(2000)          \t& +50:52:30\t& +51:36:00  \\\\ Distance (pc)        \t& 900\t\t& 1298           \\\\ Ang. diam (arcmin)   & 14.0 \t\t& 12.0            \\\\ E(B-V) (mag)\t\t& 0.854\t\t& 0.807          \\\\ log(age)\t\t\t& 8.138\t\t & 8.142         \\\\ \\hline \\end{tabular} \\end{table}  Table 2 presents previous photographic studies of the open clusters NGC 7031 and NGC 7086. The scatter, especially in the ages, is quite large for each cluster; all studies agree on a larger distance to NGC 7086 in comparison with the distance proposed for NGC 7031.  For a good selection criterion for a binary cluster Dieball 2002 considers: (i) the maximum centre-to-centre separation is $\\approx$ 20 pc and (ii) the age difference between the components of a binary cluster is either $\\leq$ 10 Myr or their ages agree well within the uncertainties of their age determination. Based on N-body simulations Portegies Zwart \\& Rusli 2007 conclude that a cluster pair with a smaller initial separation tends to merge in $\\lesssim$ 60 Myr due to loss of angular momentum from escaping stars, while clusters with a larger initial separation tend to become even more widely separated due to mass loss from the evolving stellar populations. This suggests that 20 pc is a good selection criterion.  The aim of our investigation is to determine more precisely cluster parameters such as reddening, distance and age using CCD photometry and applying criterion for binarity in order to confirm or disprove their binarity.  \\begin{table*} \\caption{ Previous studies of the open clusters NGC 7031 and NGC 7086.} \\begin{tabular}[h]{ccccc} \\hline\\noalign {\\smallskip} Name \t    & Citation \t\t\t\t     \t& Distance & Age & E(B-V) \\\\ &       \t   \t\t\t     \t\t&    pc        & Myr &  mag    \\\\ \\hline\\noalign {\\smallskip} NGC 7031   & Svolopoulos (1961)\t     \t\t& 760\t  &        & 1.03  \\\\ ........ & Hoag et al. (1961)\t    \t\t& 900\t  & 137 & 0.85\t \\\\ ........ & Lindoff (1968)\t\t     \t\t& 910 \t  & 56   & \t\t \\\\ ........ & Hassan \\& Barbone (1973) \t& 710\t  & 56   & 0.71\t \\\\ ........ & Ruprecht et al. (1981) \t\t&\t\t  & 280 &\t\t \\\\ ........ & Janes \\& Adler (1982)    \t\t& 700\t  &        & 0.71\t \\\\ NGC 7086   & Hassan (1967)\t\t     \t\t& 1170\t  & 600 & 0.69\t \\\\ ........ & Lindoff (1968)\t\t     \t\t& 1205\t  & 85   &\t\t \\\\ ........ & Janes \\& Adler (1982)\t     \t& 1200\t  &        & 0.70\t \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "We have presented BV CCD photometry of the closely projected open star clusters NGC 7031 and NGC 7086. The results have been summarized in Table 4. Our estimations of the age difference between NGC 7031 and NGC 7086, and difference in distance along the line of sight are 46 Myr and 124 pc, respectively. These results do not match with the criterion for a binary cluster and based on these results we conclude that the two clusters are most likely not formed together from one and the same GMC and they are not a true binary cluster.  However, the ages of the two clusters roughly agree \\\\within the errors, and if we assume a mean distanse of  893 pc and an angular separation of 1$\\degr$, this corresponds to a separation of 15 pc between the two clusters.These results are close enough for true binary. We caution that the distance differences  seem to indicate that the cluster separation is probably larger than 120 pc, thus it seems unlikely the two clusters form a true binary cluster, but it cannot be completely ruled out. Follow-up observations to determine the radial velocities of the clusters members and the motion of the clusters itself could put more light to this problem."
    },
    "0808/0808.1839_arXiv.txt": {
        "abstract": "Using measured radial velocity data of nine double lined spectroscopic binary systems NSV 223, AB And, V2082 Cyg, HS Her, V918 Her, BV Dra, BW Dra, V2357 Oph, and YZ Cas, we find corresponding orbital and spectroscopic elements via the method introduced by Karami \\& Mohebi (2007a) and Karami \\& Teimoorinia (2007). Our numerical results are in good agreement with those obtained by others using more traditional methods. ",
        "introduction": "\\label{intro} Determining the orbital elements of binary stars helps us to obtain fundamental information, such as the masses and radii of individual stars, that has an important role in understanding the present state and evolution of many interesting stellar objects. Analysis of both light and radial velocity (hereafter RV) curves, derived from photometric and spectroscopic observations, respectively, yields a complete set of basic absolute parameters. One historically well-known method to analyze the RV curve is that of Lehmann-Filh\\'{e}s (cf. Smart, 1990).   In the present paper we use the method introduced by Karami $\\&$ Mohebi (2007a) (=KM2007a) and Karami $\\&$ Teimoorinia (2007) (=KT2007), to obtain orbital parameters of the nine double-lined spectroscopic binary systems: NSV 223, AB And, V2082 Cyg, HS Her, V918 Her, BV Dra, BW Dra, V2357 Oph, and YZ Cas.  The NSV 223 is a contact system of the A type and the mass ratio is believed to be small (Rucinski et al. 2003a,b). The large semiamplitudes ,$K_i$, suggest that the orbital inclination is close to $90^\\circ$. The spectral type is $F7V$ and the period is $0.366128$ days (Rucinski et al. 2003a,b). The AB And is a contact binary. The spectral type is $G8V$. The period is $0.3318919$ days. It is suggested that observed period variability may be a result of the orbital motion in a wide triple system. The third body should then have to be a white dwarf (Pych et al. 2003,2004). V2082 Cyg is most probably an A-type contact binary with a period of $0.714084$ days. The spectral type is $F2V$ (Pych et al. 2003,2004). In HS Her, the effective temperatures were found to be $T_1=15200\\pm750K$ and $T_2=7600\\pm400K$ for the primary and secondary stars, respectively (Cakirli et al. 2007). The secondary component appears to rotate more slowly. The presence of a third body physically bound to the eclipsing pair has been suggested by many investigators. The two component are located near the zero-age main sequence, with age of about $32$ Myr (Cakirli et al. 2007). It is classified as an Algol-type eclipsing binary and single-lined spectroscopic binary. The spectral type of more massive primary component is $B4.5V$. The effective temperature is about $15200\\pm750K$ for the primary component and $7600\\pm400K$ for the secondary component. The period is $1.6374316$ days (Cakirli et al. 2007). The V918 Her is an A-type contact binary. The spectral classification is A7V. The period of this binary star is $0.57481$ days (Pych et al. 2003,2004). The BV Dra  and BW Dra have a circular orbit. From spectrophotometry the components of BV Dar are classified as F9 and F8 while the components of BW Dra are G3 and G0. The period of BV Dra is $0.350066$ days and for BW Dra is $0.292166$ days (Batten \\& Wenxian 1986). The V2357 Oph was classified as a pulsating star with a period of $0.208$ days. The spectral type is G5V (Rucinski et al. 2003a,b). The spectral type of primary component of YZ Cas is A1V and for the secondary is F7V. The period of this binary is $4.4672235$ days (Lacy 1981).  This paper is organized as follows. In Sect. 2, we give a brief review of the method of KM2007a and KT2007. In Sect. 3, the numerical results are reported, while the conclusions are given in Section 4. ",
        "conclusions": "Using the method introduced by KM2007a and KT2007, we obtain both the orbital elements and the combined spectroscopic parameters of nine double lined spectroscopic binary systems. Our results are in good agreement with the those obtained by others using more traditional methods. In a subsequent paper we intend to study the other different systems. \\\\ \\\\ \\noindent{{\\bf Acknowledgements}}. This work has been supported financially by Research Institute for Astronomy $\\&$ Astrophysics of Maragha (RIAAM), Maragha, Iran."
    },
    "0808/0808.2159.txt": {
        "abstract": "We are conducting a large program to classify newly discovered Milky Way star cluster candidates from the list of Froebrich, Scholz \\& Raftery \\cite{2007MNRAS.374..399F}. Here we present deep NIR follow-up observations from ESO/NTT of 14 star cluster candidates. We show that the combined analysis of star density maps and colour-colour/magnitude diagrams derived from deep near-infrared imaging is a viable tool to reliably classify new stellar clusters. This allowed us to identify two young clusters with massive stars, three intermediate age open clusters, and two globular cluster candidates among our targets. The remaining seven objects are unlikely to be stellar clusters. Among them is the object FSR\\,1767 which has previously been identified as a globular cluster using 2MASS data by Bonatto et al. \\cite{2007MNRAS.381L..45B}. Our new analysis shows that FSR\\,1767 is not a star cluster. We also summarise the currently available follow-up analysis of the FSR candidates and conclude that this catalogue may contain a large number of new stellar clusters, probably dominated by old open clusters. ",
        "introduction": "Star clusters are the building blocks of the stellar component of galaxies. Identifying and characterising clusters is thus a crucial step for our understanding of structure formation and mass assembly in the Milky Way. The benefits of wide-field studies of stellar clusters are twofold:  i) Open clusters are currently the most important sites of star formation and early stellar evolution. In particular, the formation of massive stars is intrinsically linked to the formation of star clusters. For example, the cluster mass is thought to correlate with the mass of the most massive star in the cluster (Weidner \\& Kroupa \\cite{2006MNRAS.365.1333W}). Investigating age spread, morphology, and mass segregation in a large sample of young open clusters has the potential to put limits on models for massive star formation (see review by Beuther et al. \\cite{2007prpl.conf..165B}). Furthermore, by probing the distribution of masses in a diverse sample of clusters allows us to constrain the impact of environment on the outcome of star formation and to probe the diversity and the origin of the Initial Mass Function -- fundamental problems in current star formation theory (see review by Bonnell et al. \\cite{2007prpl.conf..149B}).  As open clusters dissolve, the stars migrate into the field. Therefore, the study of old open clusters can shed light on the timescales for cluster disruption and the underlying physical processes such as tidal interactions with giant molecular clouds, evaporation of low-mass objects, mass loss due to stellar evolution, as well as mass segregation. The currently known sample of old open clusters is very incomplete (e.g. Bonatto \\& Bica \\cite{2007A&A...473..445B}). It is hence an important task to enlarge the sample of known and well classified galactic open clusters.  ii) Cluster surveys have the potential to discover new Globular Clusters (GC), a particular interesting avenue given the outstanding importance of this type of cluster. As emphasised by Harris \\cite{1996AJ....112.1487H}, Milky Way GCs have proven throughout the last century to be ``irreplaceable objects in an amazingly wide range of astrophysical studies''. The most important issue is the continuing debate about the key processes in galaxy formation and evolution. GCs have been considered for many years to be the most valuable tracers of the oldest stellar population in our galaxy.  Recent advances point to a complex picture of the genesis of our Galaxy, driven by a mixture of processes including rapid protogalactic collapse, accretion, cannibalism, galaxy collisions, and star bursts. The rich source of historical details provided by the (inhomogeneous) Milky Way GC system is among the most promising approaches to disentangle the many processes (West et al. \\cite{2004Natur.427...31W}). A complete census of the Milky Way GC system is therefore very important. More recently, a few examples of new GC classes were detected in nearby galaxies: the so-called FF (`faint fuzzies') clusters in two lenticular galaxies (Brodie \\& Larsen \\cite{2002AJ....124.1410B}), even more extended GCs in the halo of M31 (Huxor et al. \\cite{2005MNRAS.360.1007H}), very massive GCs in NGC\\,5128 (Martini \\& Ho \\cite{2004ApJ...610..233M}) and ultracompact objects (perhaps bridging the gap in parameter space to dwarf galaxies) in the Fornax cluster (Mieske et al. \\cite{2002A&A...383..823M,2004A&A...418..445M}). There are no known analogues for such unusual GCs in our galaxy where they should have been discovered unless they are hidden by dust in the classical `Zone of Avoidance', the least complete area for the Milky Way GC sample.  Recent studies suggest that the currently known sample of Milky Way GCs is incomplete (see Bonatto et al. \\cite{2007MNRAS.381L..45B} or Bica et al. \\cite{2007A&A...472..483B}), particularly at the low-luminosity end (the Palomar-type GCs). The number of missing GCs close to the Galactic Plane ($|z| < 0.5$\\,kpc) and within 3\\,kpc from the Galactic Centre has been estimated to $\\sim 10\\pm3$ (Ivanov et al. \\cite{2005A&A...442..195I}).  Driven by the aforementioned science goals, we have recently carried out a systematic large-scale cluster survey based on star density maps derived from the 2MASS database (Froebrich, Scholz \\& Raftery \\cite{2007MNRAS.374..399F}, hereafter FSR). The FSR survey covers the entire Galactic Plane ($|b| < 20^\\circ$) and detected a total number of 1788 potential star clusters, from which 767 have been known beforehand and 1021 are unknown cluster candidates. The contamination of those candidates has been estimated to be about 50\\,\\%, indicating that the catalogue may contain up to 500 new star clusters.  In particular, the FSR survey revealed several promising GC candidates in the Zone of Avoidance. Four of them have been discussed already in detail elsewhere (Froebrich et al \\cite{2007MNRAS.377L..54F}, \\cite{2008MNRAS.383L..45F}; Bica et al. \\cite{2007A&A...472..483B}, Bonatto et al. \\cite{2007MNRAS.381L..45B}). Given the small number ($\\sim 150$) of known galactic GCs on the one hand and the diversity of the GC species on the other hand (see above), every new GC is of value.  In this paper, we present detailed follow-up analysis for 14 cluster candidates from the FSR survey, selected to be among the best candidates for new GCs. The paper is structured as follows. In Sect.\\,\\ref{dataanalysis} we present our new observations and the reduction of the data. The detailed analysis and results for each individual cluster, including the appearance of the cluster, the contamination with field stars and the isochrone fitting to determine the cluster properties are presented in Sect.\\,\\ref{results}. Finally in Sect.\\,\\ref{discussion} we discuss and conclude our findings. ",
        "conclusions": "Das Paper zeigt dreierlei: a) Einmal ist die %  vorgestellte Methode gut geeignet, um Haufen zu erkennen. Du verwendest ja viel %  Zeit auf das Vorstellen der Methode, da sollte man hinterher auch sagen, wie gut %  sie funktioniert. Gleichzeitig dann warnen, dass tiefe Daten notwendig sind, %  Bsp. FSR1767. b) Zum anderen sieht es so aus, als seien auch bei den am besten %  aussehendsten Kandidaten von FSR etwa die Haelfte keine Haufen, also sind die in %  FSR angegeben 50\\% fuer das Gesamtsample vermutlich zu optimistisch. c) Unter %  den gefundenen Haufen finden sich aber ein paar recht interessante, z.B. der %  vermutete GC und auch die ganz jungen offenen. %"
    },
    "0808/0808.3828_arXiv.txt": {
        "abstract": "We investigated the light propagation by means of the Robertson-McVittie solution which is considered to be the spacetime around the gravitating body embedded in the FLRW (Friedmann-Lema{\\^i}tre-Robertson-Walker) background metric. We concentrated on the time delay and derived the correction terms with respect to the Shapiro's formula. To relate with the actual observation and its reduction process, we also took account of the time transformations; coordinate time to proper one, and conversely, proper time to coordinate one. We applied these results to the problem of increase of astronomical unit reported by Krasinsky and Brumberg (2004). However, we found the influence of the cosmological expansion on the light propagation does not give an explanation of observed value, $d{\\rm AU}/dt = 15 \\pm 4$ [m/century] in the framework of Robertson-McVittie metric. ",
        "introduction": "} The appearance of radar/laser gauging techniques and spacecraft ranging has enabled us to measure the distance from the Earth to other inner planets and the moon directly and precisely. Recent years, these techniques have been drastically improved, for instance, the planetary radar measurement achieves the observational accuracy of the interplanetary distance within a few 100 [m], the spacecraft ranging a few [m] and the lunar laser ranging a few [cm]. In these observations, the round-trip time of light/signal is measured by the atomic clocks on the Earth. Nowadays, the advent of laser cooled atomic clocks has led to the improvement of realization of SI second and it is expected that this advance will go on. Therefore the arrival time measurement is the most accurate observation among the all kinds of astronomical and astrophysical ones.  Due to the improvement of these techniques, it is important to develop the rigorous light propagation model which contains the various general relativistic and relating effects, and this subject is actively investigated by many authors, e.g. \\cite{shapiro1964,richter1983,hellings1986,kopeikin1997,kopeikin1999, kopeikin2002,klioner2003a,klioner2003b,defellice2004,defelice2006,asada} and references therein. These theoretical developments also play a crucial role in testing the gravitational theories \\cite{will1,will2}.  On the other hand, an arrival time measurements have also devoted to the improvement of lunar and planetary ephemerides, such as DE \\citep{standish2003}, EPM \\citep{pitjeva2005}, VSOP \\citep{bretagnon1988}, and INPOP \\citep{fienga2008}, and the determination of astronomical constants. Especially among these constants, the astronomical unit (of length) AU is one of the fundamental and important one which gives the relation of length units; [AU] of astronomical system of units and [m] of SI ones. AU is currently determined from the arrival time measurement data of planetary radar and spacecraft ranging within the accuracy of 0.1 [m] or 12 digits level as $1 ~[{\\rm AU}] = 1.495978706960 \\times 10^{11} \\pm 0.1 ~[{\\rm m}]$ \\cite{pitjeva2005}. Therefore from the point of view of the fundamental astronomy, the development of a accurate light propagation model is of great importance.  Until now, the light propagation formulae have been derived in terms of the post-Newtonian approximation, $g_{\\mu\\nu} = \\eta_{\\mu\\nu} + h_{\\mu\\nu}, ~|h_{\\mu\\nu}| \\ll 1$ in which the background metric $\\eta_{\\mu\\nu}$ is the static Minkowski one. However, it is interesting to examine the influence of cosmological expansion on the light/signal propagation  since the accuracy of astronomical observation is rapidly growing. Of course, it is presently hard to observe such a influence in the solar system experiments. Nevertheless, the progress of these or relating measurement techniques may enable to detect its trace in the future date.  In this paper, we will study the light propagation passing near the massive star as the Sun which is embedded in the cosmological expanding background. We will adopt the Robertson-McVittie spacetime introduced in section \\ref{rm-metric}, and concentrate on the gravitational time delay and derive the extra delay due to the effect of cosmological expansion with respect to the standard formula by Shapiro in section \\ref{time-delay}. Further in order to relate with the actual observation and its reduction process, we will take account of the time transformations in section \\ref{tt}; from coordinate time to proper one in subsection \\ref{ctp}, and conversely, from proper time to coordinate one in subsection \\ref{ptc}. As the application of results, we will consider the problem of increase of astronomical unit recently reported by \\cite{kb2004} in section \\ref{au}. ",
        "conclusions": "In this paper, we investigated the light propagation passing near the gravitating body as the Sun embedded in the cosmological spacetime by means of the Robertson-McVittie solution. We concentrated on the time delay and derived the correction terms with respect to the Shapiro's formula. We also took account of the time transformations; coordinate time to proper one, and conversely, proper time to coordinate one. These treatments are of importance when dealing with the observed round-trip time and its reduction process.  As the application of results, we considered the problem of increase of astronomical unit reported by Krasinsky and Brumberg (2004), since the AU is currently derived from the analysis of the round-trip time of light/signal. However, the effect of the cosmological expansion is too small to produce the observed $d{\\rm AU}/dt = 15$ [m/century] and then it dose not give the explanation of increase of AU.  Here it is worth to mention the interpretation of $d{\\rm AU}/dt$. Although we and Krasinsky and Brumberg (2004) investigated the secular increasing of astronomical unit in terms of the cosmological effect, another possibility is suggested, that is, increasing of AU is arisen due to the lack of the calibrations of internal delays of the radio signals within the spacecrafts and this may be the most plausible reason of observed $d{\\rm AU}/dt$. Nonetheless, till now, the origin of $d{\\rm AU}/dt$ is far from clear therefore this issue should be explored by means of the every possibility we can imagine. And the re-analysis of $d{\\rm AU}/dt$ using new data sets is also expected.  As we mentioned in the end of section \\ref{au}, so far, the dynamical terms due to cosmological expansion in the planetary dynamics has been studied actively, while the influence on light/signal propagation has hardly been examined. However, the accuracy of astronomical/astrophysical observations, especially the optical, radio and arrival time ones, rapidly increases in the solar system and they may make it necessary to incorporate the cosmological effect in the future date. For this purpose, our investigation in present paper is the first attempt to this direction. To consider this problem, the matching of BCRS and cosmological reference system discussed by \\cite{ks2004,kopeikin2007} is also important issue to be investigated.  (\\ref{base-metric}) expresses the completely static gravitational field, and is equivalent to Schwarzschild-de Sitter model. Nevertheless, as long as we focus on the solar system experiments, our results may be regarded as the dominant effects due to the cosmological expansion, see in (\\ref{hubble_0}). However in order to study the cosmological influence on the light/signal propagation in more detail and general cases, we must treat $H(t)$ and/or $a(t)$ as a function of time and adopt the time dependent spacetime model. This subject will be investigated at some future time.   We would like to express our gratitude to the referee for fruitful comments and suggestions. We acknowledge to G. A. Krasinsky for providing the information and comments about the AU issue. We also appreciate T. Fukushima, M. Kasai, H. Asada, Y. Itoh and M. Takada for fruitful discussions and comments. This work was partially supported by the Ministry of Education, Science, Sports and Culture, Grant-in-Aid, 18740165.  \\appendix"
    },
    "0808/0808.2978.txt": {
        "abstract": "We consider the problem of automatically (and robustly) isolating and extracting information about waves and oscillations observed in EUV image sequences of the solar corona with a view to near real-time application to data from the Atmospheric Imaging Array (AIA) on the {\\em Solar Dynamics Observatory} (\\sdo). We find that a simple coherence / travel-time based approach detects and provides a wealth of information on transverse and longitudinal wave phenomena in the test sequences provided by the {\\em Transition Region and Coronal Explorer} (\\trace). The results of the search are \u00d2pruned\u00d3 (based on diagnostic errors) to minimize false-detections such that the remainder provides robust measurements of waves in the solar corona, with the calculated propagation speed allowing automated distinction between various wave modes. In this paper we discuss the technique, present results on the \\trace{} test sequences, and describe how our method can be used to automatically process the enormous flow of data ($\\approx$1Tb/day) that will be provided by \\sdo/AIA after launch in late 2008. ",
        "introduction": "%\\label{s:intro}  Since the launch of the {\\em Transition Region and Coronal Explorer} \\citep[\\trace;][]{1999SoPh..187..229H} the community has invested a great deal of effort on the identification and analysis of longitudinal and transverse wave phenomena in loop structures seen in its EUV images of the corona \\citep[see][for a good overview]{2005LRSP....2....3N}. The great interest in finding and characterizing coronal waves is driven by the promise of coronal seismology which could lead to the determination of otherwise inaccessible physical properties of the solar atmosphere by studying phase speeds, amplitudes, dissipation, {\\it etc.} of the observed waves \\citep[][]{2005LRSP....2....3N}.  There are many unresolved issues in coronal seismology: for example, it is unclear why only a subset of coronal loops show transverse oscillations in the wake of a flare or CME \\citep[{\\it e.g.},][]{2002SoPh..206...99A}, which physical processes dominate the damping of these transverse oscillations \\citep[{\\it e.g.},][]{2002ApJ...576L.153O, 2007AAP...463..333A, 2007SoPh..246..231T}, why only a subset of coronal loops show propagating slow-mode oscillations \\citep[{\\it e.g.},][]{2000AAP...355L..23D, 2007SoPh..246...53D}, how the propagation of slow-mode waves at different temperatures can be used to probe the thermal structure and loop scale heights \\citep[][]{2003AAP...404L...1K}, or how lower-atmospheric oscillations generally leak into the corona \\citep[{\\it e.g.},][]{2005ApJ...624L..61D, 2006ApJ...648L.151J, 2008ApJ...676L..85K}?  Many of these issues require more extensive study with much larger statistical samples. It is surprising that the number of wave and oscillation detections in the above papers is only of order several dozen for both longitudinal and transverse oscillations. This is a very small number compared to the volume of \\trace{} data given that flares occur often and ``wave-leakage'' from the solar interior to the lower portions of the outer solar atmosphere is apparently quite abundant \\citep[][]{2005ApJ...624L..61D,2006ApJ...648L.151J,2008ApJ...676L..85K}. The discrepancy between the plethora of observed wave driven phenomena in the chromosphere and transition region, and the relative paucity of waves seen in the corona partly motivates the effort discussed herein. Is the small number of coronal wave observations caused by the fact that all of the \\trace{} observations of waves have been found as a result of a ``manual'' search, {\\it i.e.}, limited by human patience and detection skills? Or are there actually very few locations in the corona where waves can be observed?  One of our aims is to develop an automated wave-detection algorithm to help expand the number of known oscillations, and build up significant statistics on how frequent coronal oscillations occur. Such larger statistical samples are crucial to address many of the unresolved issues in coronal seismology, and will be within reach with the Atmospheric Imaging Assembly (AIA) after launch of the {\\em Solar Dynamics Observatory} (\\sdo) by 2009. This instrument will provide full-disk coronal imaging at significantly increased signal-to-noise in eight coronal wavelengths at ten-second cadence continuously, producing an enormous data flow of $\\approx$1Tb day$^{-1}$. The potential for significantly increased detection of coronal oscillations and waves with such a data rate and quality are very high. However, AIA's enormous data rate renders business-as-usual approaches that involve manually looking through data sets and individually flagging events unfeasible. Since it is unpractical to transfer the full volume of AIA data to remote users, it is critical that automated algorithms be developed that can search for wave-like phenomena and automatically flag temporal and spatial regions of interest for later, more detailed analysis. There have been a few papers recently  that have started addressing some of the issues that need to be resolved to enable full optimal use of the AIA data, such as automated detection of locations with significant oscillatory power \\citep[][]{2007SoPh..241..397N}, or semi-automated detection of propagating and standing waves \\citep[][]{2007SoPh..tmp..102S}. The approach presented in this paper may be the first that can reliably and automatically distinguish between longitudinal and transverse oscillations at the same time as rejecting most false positive detections of oscillations. Our approach is based on the technique used to analyze and detect coronal Alfv\\'{e}n waves with the CoMP instrument \\citep[][]{2007Sci...317.1192T}.  In the following section we discuss the method developed, illustrating its application on example \\trace{} datasets (see {\\it Sect.}~\\pref{s:data}) which are known to show a host of transverse and longitudinal wave phenomena and are therefore good for testing our approach. In {\\it Sect.}~\\pref{s:results} we explore the results of the analysis on the test data and discuss the method employed to minimize false detections. In {\\it Sect.}~\\pref{s:discuss} we compare some of our results with previous analyses of the \\trace{} datasets we have used, and discuss how our method could be used in a practical manner for \\sdo/AIA data. ",
        "conclusions": "\\label{s:conclude} % We have considered the problem of automatically (and robustly) isolating and extracting information about waves and oscillations observed in EUV image sequences of the solar corona with a view to near real-time application to data from the Atmospheric Imaging Array (AIA) on the {\\em Solar Dynamics Observatory} (\\sdo). We have found that a simple coherence / travel-time based approach detects and provides a wealth of information on transverse and longitudinal wave phenomena in the test sequences provided by the {\\em Transition Region and Coronal Explorer} (\\trace). The results of the search can be robustly ``pruned'' (based on diagnostic errors) to minimize false-detections such that the remainder provides reliable measurements of waves in the solar corona, with the calculated propagation speed allowing automated distinction between various wave modes and can be automatically applied to the enormous flow of data ($\\approx$1Tb day$^{-1}$) that will be provided by \\sdo/AIA. In addition to demonstrating the applicability of this approach we have found signatures of some interesting low frequency (1.5~mHz) oscillatory phenomena in the \\trace{} test data that motivate further study.  %% Figure % \\begin{figure} \\centerline{\\includegraphics[width=0.5\\textwidth,clip=true]{mcintoshf1a.eps} \\includegraphics[width=0.5\\textwidth,clip=true]{mcintoshf1b.eps}} \\caption{Examples of the coherence, weighted cross-power, phase difference and phase travel-time for for the 1 July 1998 (panels A-D) and 14 July 1998 (panels E -- H) datasets for 5.5~mHz and 1.5~mHz Fourier filter samples respectively. In each case the reference pixel is at 0,0, the solid black contour outlines regions of highly coherent signal and the fitted angle (solid white lines in panels C and G) to central coherence ``island'' is shown as a solid white line in panels C (14.6$^\\circ$) and G (-30.4$^\\circ$).}\\label{fig1} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.5\\textwidth,clip=true]{mcintoshf2a.eps} \\includegraphics[width=0.5\\textwidth,clip=true]{mcintoshf2b.eps}} \\caption{Examples of phase speed determination for the cases shown in Figure~\\pref{fig1}. In each case the green triangles mark the positions from the reference pixel and measured phase travel-time at that pixel. The gradient of the least-squares linear fit (red dot-dashed line) yields the phase-speed of the propagating signal.}\\label{fig2} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf3.eps}} \\caption{Example pixel-to-pixel travel-time analysis in the temporal domain. In panels A and B we show the \\trace{} 171\\AA{} lightcurve (black solid line) and its 5.5~mHz filtered counterpart (blue solid line) between the reference pixel [181, 97] and pixel [179, 97] for the 1 July 1998 dataset ({\\it cf}. Figure~\\pref{fig1}). In panel C we show the derived cross-correlation function (black solid line), envelope maxima (red triangles) and fit to the central peak (green solid line). The vertical dashed lines in panel C show the positions of the estimated group travel-time (red; -4.17 minutes) and phase travel-time (green; -0.527 minutes, {\\it i.e.}, consistent with the values in Figures~\\pref{fig1} and propagating left to right which is consistent with the movie of Figure~\\pref{fig4}) respectively.}\\label{fig3} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.5\\textwidth,clip=true]{mcintoshf4a.eps} \\includegraphics[width=0.5\\textwidth,clip=true]{mcintoshf4b.eps}} \\caption{Illustrations of the two \\trace{} 171\\AA{} timeseries used in the presented analysis; 1 July 1998 (left) and 14 July 1998 (right). In each panel we show (clockwise from the top left) snapshots of the \\trace{} 171\\AA{} intensity, 1.5~mHz, 5.5~mHz, and 3.5~mHz filtered timeseries respectively. The boxes marked on the panels correspond to the regions used to construct Figure~\\pref{fig1} and the subsequent analyses. See the online edition of the Journal for animations of these figures.}\\label{fig4} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf5a.eps}} \\caption{Results of the coronal wave detection algorithm for the 1.5~mHz filtered timeseries for the 1 July 1998 dataset. We show the average \\trace{} 171\\AA{} intensity image, weighted signal coherence, length, width, angle, phase speed, and relative error of the phase speed computed from the large coherence island.}\\label{fig5a} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf5b.eps}} \\caption{Results of the coronal wave detection algorithm for the 3.5~mHz filtered timeseries for the 1 July 1998 dataset. We show the average \\trace{} 171\\AA{} intensity image, weighted signal coherence, length, width, angle, phase speed, and relative error of the phase speed computed from the large coherence island.}\\label{fig5b} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf5c.eps}} \\caption{Results of the coronal wave detection algorithm for the 5.5~mHz filtered timeseries for the 1 July 1998 dataset. We show the average \\trace{} 171\\AA{} intensity image, weighted signal coherence, length, width, angle, phase speed, and relative error of the phase speed computed from the large coherence island.}\\label{fig5c} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf6a.eps}}9 \\caption{Results of the coronal wave detection algorithm for the 1.5~mHz filtered timeseries for the 1 July 1998 dataset. We show the average \\trace{} 171\\AA{} intensity image, weighted signal coherence, length, width, angle, phase speed, and relative error of the phase speed computed from the large coherence island.}\\label{fig6a} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf6b.eps}} \\caption{Results of the coronal wave detection algorithm for the 3.5~mHz filtered timeseries for the 14 July 1998 dataset. We show the average \\trace{} 171\\AA{} intensity image, weighted signal coherence, length, width, angle, phase speed, and relative error of the phase speed computed from the large coherence island.}\\label{fig6b} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf6c.eps}} \\caption{Results of the coronal wave detection algorithm for the 5.5~mHz filtered timeseries for the 14 July 1998 dataset. We show the average \\trace{} 171\\AA{} intensity image, weighted signal coherence, length, width, angle, phase speed, and relative error of the phase speed computed from the large coherence island.}\\label{fig6c} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf7a.eps}} \\caption{Final results of the coronal wave detection algorithm for the 1.5~mHz filtered timeseries for the 14 July 1998 dataset that have been ``pruned'' to reduce the appearance of false positive detections using the technique discussed in {\\it Sect.}~\\pref{s:rfp}. We show the average \\trace{} 171\\AA{} intensity image and the region averaged signal coherence, length, angle and phase speed, and the errors in the latter.}\\label{fig7a} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf7b.eps}} \\caption{Final results of the coronal wave detection algorithm for the 3.5~mHz filtered timeseries for the 1 July 1998 dataset that have been ``pruned'' to reduce the appearance of false positive detections using the technique discussed in {\\it Sect.}~\\pref{s:rfp}. We show the average \\trace{} 171\\AA{} intensity image and the region averaged signal coherence, length, angle and phase speed, and the errors in the latter.}\\label{fig7b} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf7c.eps}} \\caption{Final results of the coronal wave detection algorithm for the 5.5~mHz filtered timeseries for the 1 July 1998 dataset that have been ``pruned'' to reduce the appearance of false positive detections using the technique discussed in {\\it Sect.}~\\pref{s:rfp}. We show the average \\trace{} 171\\AA{} intensity image and the region averaged signal coherence, length, angle and phase speed, and the errors in the latter.}\\label{fig7c} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf8a.eps}} \\caption{Final results of the coronal wave detection algorithm for the 1.5~mHz filtered timeseries for the 14 July 1998 dataset that have been ``pruned'' to reduce the appearance of false positive detections using the technique discussed in {\\it Sect.}~\\pref{s:rfp}. We show the average \\trace{} 171\\AA{} intensity image and the region averaged signal coherence, length, angle and phase speed, and the errors in the latter.}\\label{fig8a} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf8b.eps}} \\caption{Final results of the coronal wave detection algorithm for the 3.5~mHz filtered timeseries for the 14 July 1998 dataset that have been ``pruned'' to reduce the appearance of false positive detections using the technique discussed in {\\it Sect.}~\\pref{s:rfp}. We show the average \\trace{} 171\\AA{} intensity image and the region averaged signal coherence, length, angle and phase speed, and the errors in the latter.}\\label{fig8b} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth,clip=true]{mcintoshf8c.eps}} \\caption{Final results of the coronal wave detection algorithm for the 5.5~mHz filtered timeseries for the 14 July 1998 dataset that have been ``pruned'' to reduce the appearance of false positive detections using the technique discussed in {\\it Sect.}~\\pref{s:rfp}. We show the average \\trace{} 171\\AA{} intensity image and the region averaged signal coherence, length, angle and phase speed, and the errors in the latter.}\\label{fig8c} \\end{figure}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %% Acknowledgements %"
    },
    "0808/0808.1646_arXiv.txt": {
        "abstract": "We study the evolution of linear density perturbations in the context of interacting scalar field-dark matter cosmologies, where the presence of the coupling acts as a stabilization mechanism for the runaway behavior of the scalar self-interaction potential as in the case of the Chameleon model. We show that in the ``adiabatic'' background regime of the system the rise of unstable growing modes of the perturbations is suppressed by the slow-roll dynamics of the field. Furthermore the coupled system behaves as an inhomogeneous adiabatic fluid. In contrast instabilities may develop for large values of the coupling constant, or along non-adiabatic solutions, characterized by a period of high-frequency dumped oscillations of the scalar field. In the latter case the dynamical instabilities of the field fluctuations, which are typical of oscillatory scalar field regimes, are amplified and transmitted by the coupling to dark matter perturbations. ",
        "introduction": "Cosmology has provided evidence of a dark physics sector which is necessary to account for about $95\\%$ of the cosmic matter content \\cite{CMB}. Despite the success of the $\\Lambda$CDM model to fit all cosmological observations, the existence of the dark energy phenomenon as well as its relation to the abundance and clustering of matter in the universe still pose puzzling questions.  Models of interacting dark energy-dark matter have been proposed to address such problems. In this scenario dark energy is a fundamental scalar field which directly couples to matter particles. This allows for a dynamical solution of the so called ``coincidence'' problem, since independently of the initial conditions the scalar interaction drives the dark energy-to-matter ratio toward a constant value (see e.g. \\cite{LucaDomenico,Chimento,Pietroni,JMAndre1}). These models are inspired by string and supergravity theories, where the compactification of extra-dimensions in the low-energy gives rise to massless scalars coupled to matter fields with gravitational strength. Therefore a distinct feature of this scenario is that matter particles experience a long-range scalar force and acquire a time dependent mass which cause violations of the Equivalence Principle (EP). The tight bounds imposed by EP tests are usually avoided as consequence of other possible mechanisms. As an example Damour and Polyakov have shown that in String Theory the couplings between the dilaton and different matter fields can be dynamically suppressed \\cite{Damour}. An interesting possibility has been proposed in the ``Chameleon'' model \\cite{Justin}, where the mass of the scalar field is assumed to depend on the local matter density. In such a case fifth-force effects can be strongly suppressed on Solar System scales, thus avoiding EP bounds. Another possibility has been explored in \\cite{JMAndre2} where the authors consider a dilatonic field to be differently coupled to various matter species such that the system can naturally evolve toward a late time attractor solution where General Relativity is recovered. Non-minimally coupled models can successfully describe the background expansion of the universe as probed by supernova type Ia luminosity distance or the position of the Doppler peaks in the Cosmic Microwave Background anisotropy power spectrum (see e.g. \\cite{Das,OlivaresAlimiFuzfa}). However testing the formation of structure in the universe more than standard cosmological tests may provide a key insight on this class of models. In fact the scalar coupling contributes to modifying the clustering properties of matter, implying that an accurate study of the evolution of density fluctuations both in the linear and non-linear phase of collapse can identifying unique signatures of dark sector interactions \\cite{PeeblesFarrar}. In the context of linear perturbation theory several interacting scalar field-dark matter models have been studied in the literature (see e.g. \\cite{DEDMintPert}). In some specific realizations it was found that the growth of linear density perturbations is spoiled by the presence of dangerous instabilities \\cite{Koivisto,Kaplinghat}, as in the case of ``Mass Varying Neutrino'' (MaVaN) models \\cite{Afshordi}. Recently a number of works have analysed the stability of perturbations in more general setups. For instance in \\cite{Bean} the authors have studied models with a background evolution characterized by an adiabatic regime, and shown that unstable growing modes of the perturbations exist for couplings much greater than gravitational strength. On the other hand the authors of \\cite{Valiviita,Abdalla} have considered the case of an interacting dark energy component with constant equation of state and found that for couplings proportional to the dark matter density the perturbations are unstable.  In this paper we provide a more detailed study of these instabilities, particularly in relation to the specificities of the background scalar field evolution. The paper is organized as follows: in Section~\\ref{DynEq} we introduce the interacting scalar field-matter model as well as the background and perturbation equations; in Section~\\ref{PertAnalysis} we present the results of our analysis; finally in Section~\\ref{Conclusions} we present our conclusions. ",
        "conclusions": "We have studied the evolution of linear perturbations in the case of an interacting scalar field with runaway potential directly coupled to dark matter particles. We have specifically analyzed the stability of perturbations during the adiabatic evolution of the field, and shown that as consequence of the slow-roll condition the onsets of instabilities is largely suppressed. This can be explained in terms of the adiabatic sound speed of the effective unified fluid. In fact during the adiabatic regime, despite being negative, it assumes negligibly small values and as consequence of this the growth of linear density perturbations remains stable. On the other hand instabilities may develop in strongly coupled adiabatic regimes, with a coupling constant much greater than gravitational strength. Interestingly during the adiabatic evolution of the field the coupled system behaves as a single adiabatic inhomogeneous fluid. We have also shown that large instabilities can spoil the growth of linear perturbations in the case of non-adiabatic solutions characterized by large scalar field oscillations. It is well known that scalar field fluctuations are unstable during oscillatory regimes, in such a case the scalar coupling amplifies and propagates such instabilities to the perturbations of the dark matter component.  Our analysis suggest that under minimal natural model assumptions Chameleon-like cosmologies are not affected by instabilities of the perturbations and can provide a viable period of structure formation more than previously believed."
    },
    "0808/0808.2645_arXiv.txt": {
        "abstract": "We present results from a numerical study of the multiphase interstellar medium in sub-Lyman-break galaxy protogalactic clumps. Such clumps are abundant at $z=3$ and are thought to be a major contributor to damped Ly$\\alpha$ absorption. We model the formation of winds from these clumps and show that during star formation episodes they feature outflows with {\\it neutral} gas velocity widths up to several hundred $\\kms$. Such outflows are consistent with the observed high-velocity dispersion in DLAs. In our models thermal energy feedback from winds and supernovae results in efficient outflows only when cold ($\\lsim300K$), dense ($\\gsim100\\msun\\pc^{-3}$) clouds are resolved at grid resolution of $12\\pc$. At lower $24\\pc$ resolution the first signs of the multiphase medium are spotted; however, at this low resolution thermal injection of feedback energy cannot yet create hot expanding bubbles around star-forming regions -- instead feedback tends to erase high-density peaks and suppress star formation. At $12\\pc$ resolution feedback compresses cold clouds, often without disrupting the ongoing star formation; at the same time a larger fraction of feedback energy is channeled into low-density bubbles and winds. These winds often entrain compact neutral clumps which produce multi-component metal absorption lines. ",
        "introduction": "Current numerical galaxy formation models can successfully reproduce some of the properties of damped Ly$\\alpha$ absorbers (DLAs), such as the lower end ($N_{\\rm HI}\\le10^{21.5}\\cm^{-2}$) of the column density distribution and the total incidence rate \\citep{pontzen........08,razoumov.....08}, the distribution of metals, and the slope of the relation between metallicity and low-ion velocity width which appears to originate in the mass-metallicity relation in the models \\citep{pontzen........08}. On the other hand, simulations tend to overpredict the number of DLAs with $N_{\\rm HI}\\gsim10^{21.5}\\cm^{-2}$ and systematically produce fewer high-velocity systems. Most such systems feature multiple components in their absorption line profiles, but unfortunately one cannot map these components from the velocity space to real space to identify the absorption regions and constrain the mechanism producing such high velocities.  In general, the velocity dispersion of neutral gas clouds can come either from the gravitational infall in the process of hierarchical buildup of galaxies, in the form of random velocities of protogalactic clumps \\citep{haehnelt..98}, or from feedback from stellar winds and supernovae (SNe) \\citep{schaye01}. In fairly massive $10^{12}\\msun$ halos at $z=3$ as much as $20-30\\%$ of gas by mass can be in the cold phase surviving the infall \\citep{razoumov.....08}. The corresponding $v_{\\rm circ}\\sim250\\kms$ can account for part of the observed neutral gas velocity dispersion. However, more massive halos are rare at $z=3$, and the fraction of cold gas drops sharply in $\\gt10^{12}\\msun$ halos, leaving us in search of other mechanisms to produce high velocities.  Galactic winds driven by the feedback energy from stellar winds and SNe are an obvious candidate \\citep{schaye01}. Star-forming Lyman-break galaxies (LBGs) at $z\\sim3$ show evidence for large-scale outflows with typical velocities of hundreds $\\kms$ \\citep{pettini.....98,pettini.......01}. In fact, with a simple semi-analytical model \\citet{mcdonald.99} showed that feedback at the rate $1.8\\times10^{50}\\erg\\yr^{-1}$ per $10^{12}\\msun$ of halo dark matter mass added to the velocity dispersion of neutral clouds inside virialized halos works out perfectly to explain the observed DLA kinematics. However, this energy transfer takes place on pc scales currently inaccessible to cosmological models. Moreover, the inability of numerical galaxy formation models to capture physics on such small scales has led to a number of predicaments, the most famous of which is the overcooling problem, accompanied by the excessive loss of angular momentum in simulated galactic disks.        This classical problem \\citep{katz92} has been somewhat alleviated in recent years \\citep{thacker.00,sommer-larsen..03} as it was realized that feedback from young stars can be very efficient at keeping gas in a diluted state preventing it from rapid collapse and conversion into stars. However, even galaxy models at a sub-kpc ($0.1-1\\kpc$) resolution cannot capture propagation of supernova blastwaves into the interstellar medium (ISM), as the injected thermal energy is radiated away very quickly before it can be converted into kinetic energy. The reason is very simple: the mass of a resolution element to which the feedback energy is supplied is usually several orders of magnitude larger than the typical mass of a SN ejecta. Therefore, the temperature and expansion velocity of the post-shock regions are greatly underestimated, and so is the cooling time which scales as $\\propto T^{1/2}$ above $10^7{\\rm\\,K}$ \\citep{dallaVecchia.08}.  Cosmological simulations must then turn to ad-hoc assumptions about the role of stellar feedback at scales below their resolution limit. Two types of solutions have been popular. The first one is suppressing radiative cooling in the feedback regions for the duration of the starburst \\citep{mori...97, thacker.00, sommer-larsen..03, stinson.....06}, to allow more efficient conversion of feedback energy into hydrodynamical expansion. This approach leads to more realistic simulated galaxies that correctly reproduce many of the observed properties of present-day disk galaxies and have only with a small deficiency in angular momentum. On the other hand due to lack of resolution the feedback energy is supplied to a fairly large mass, and the outflow velocities are usually underestimated. This method is known to produce ``puffy'' galaxies that cannot reproduce the high-end tail of the observed DLA velocity dispersion \\citep{razoumov.....08}. Since these galaxies extend to larger radii, their outer regions may pick part of the velocity dispersion of the local galaxy group, so that they have slightly less severe kinematics problem \\citep{pontzen........08} than similar-resolution models which do not suppress radiative cooling.  The second widespread approach is to use kinetic feedback instead of thermal feedback \\citep{navarro.93,springel.03,dallaVecchia.08}, usually implemented in particle-based simulations. Although there are several variations of this method, the basic idea is to give a velocity kick to a small fraction of gas particles near the star-forming regions, adjusting the mass loading and velocities to reproduce observations. Some of the recipes turn off hydrodynamical interaction of the wind particles with the gas to obtain large-scale outflows, while other stress the importance of such interaction to create hot bubbles in the ISM and develop galactic fountains \\citep{dallaVecchia.08}.  A popular method to circumvent some of the resolution problems that can be combined with either of the above two approaches is to use a sub-resolution multiphase model that describes analytically growth of cold clouds embedded in a hot intercloud medium, star formation (SF) in these clouds, feedback and cloud evaporation \\citep{yepes...97,springel.03}. In such models SF and feedback are self-regulating. However, different phases are not dynamically separated from each other, and therefore by itself such model cannot result in outflows.       At the other end of the resolution spectrum, detailed models of small patches of galactic disks, usually in the context of the Milky Way galaxy, provide sufficient resolution to study turbulent ISM stirred by SN explosions. Such models resolve hot bubbles driven by individual SNe, fragmentation of shells created by these bubbles, and the structure of cold dense clouds on pc scales \\citep[e.g.,][]{joung.06}. High SF rates in such models naturally lead to galactic outflows, galactic fountains rising to several kpc away from the midplane, and shell fragments raining back onto the disk as intermediate-velocity cold clouds \\citep{joung.06}. Such high-resolution simulations can in principle be used to develop subgrid models of stellar energy feedback for a given SF rate in cosmological simulations, although to the best of our knowledge currently there are no subgrid models in the literature that separate the dynamics of different components of the ISM.     In the past few years it has become possible to extend such high-resolution 3D models to entire galactic disks, albeit at a lower spatial resolution. \\citet{tasker.06a} used adaptive mesh refinement (AMR) models to study the multiphase ISM in a quiescent Milky Way-sized disk galaxy. They employ two SF prescriptions, one based on cosmological simulations, with low SF threshold ($0.02\\cm^{-3}$) and low efficiency, and the other one with a much higher threshold ($10^3\\cm^{-3}$) and a high efficiency. Their highest numerical resolution is $25\\pc$ which is a typical size of giant molecular clouds; they include cooling to $300{\\rm\\,K}$ and later add photoelectric heating \\citep{tasker.08}. Their models reproduce a multiphase ISM with most of the mass in cold, dense clouds, while SN feedback drives gas out of the plane of the galaxy, but most of it eventually falls back on the disk. All of their models reproduce the slope of the observed relation between the SF rate and the gas surface density, on both global and local scales, although the high-density threshold models tend to produce more intermittent outflows and occasionally triggered SF in the outer disk.  \\citet{saitoh.......08a} carried out SPH simulations of an isolated gas disk with $10^6-10^7$ particles to study the effect of various SF prescriptions on the structure of the ISM. Similar to \\citet{tasker.06a,tasker.08}, they test both a cosmological ($0.1\\cm^{-3}$) and a high-density ($100\\cm^{-3}$) SF thresholds, but also vary the SF efficiency $\\esf$. Only the high-density threshold models could reproduce the complex multiphase structure of the gas disk, regardless of the value of $\\esf$. In these runs the SF rates depend on the modeled global gas flow from intermediate densities to the actual sites of SF rather then the actual prescribed SF efficiency. On the other hand, the low-density threshold models produce thicker and smoother disks with SF rates highly sensitive to the chosen value of $\\esf$. Therefore, they conclude, the use of a high SF threshold will avoid uncertainties in the SF models.        \\citet{ceverino.08} developed a realistic prescription for modeling feedback formulating conditions under which simulations would resolve the formation of hot bubbles in the multiphase ISM. Such bubble can only be created if simulations resolve cold dense clouds in which feedback occurs -- only in these clouds heating can exceed cooling so that a large fraction of feedback energy is converted into hydrodynamical expansion, creating the hot phase and driving winds. \\citet{ceverino.08} also added heating by massive binary systems which are ejected from molecular clouds when one of the components becomes a SN. These ``runaway stars'' carry energy very efficiently away from high-density regions, eventually exploding as SNe in lower density environments and thus facilitating the formation of the hot gas component even in low-resolution cosmological models.  In this paper the approach of \\citet{ceverino.08} is used to study the formation of winds in high-redshift protogalactic clumps responsible for damped Ly$\\alpha$ absorption. We show that a brief episode of SF in a sub-LBG galaxy that creates a multiphase medium can also drive winds with {\\it neutral} gas velocity dispersions up to several hundred $\\kms$. If a substantial fraction of $z=3$ protogalactic clumps form such winds at any given time, these winds can explain the observed DLA kinematics \\citep{schaye01}. Our models resolve the effect of massive stars in protogalactic clumps. Although the spatial resolution of this study ($12\\pc$) is not sufficient to follow the details of shell fragmentation or turbulent interactions, it is argued that this resolution is adequate for modeling the multiphase medium for the purposes of computing galaxy formation and launching galactic winds in the cosmological context. ",
        "conclusions": "We use high-resolution hydrodynamic simulations of isolated $z=3$ protogalactic clumps to show for the first time that the high-end tail of the DLA neutral gas velocity width distribution can be naturally produced if a substantial fraction of clumps in $10^{10}-{\\rm few}\\times10^{11}\\msun$ halos experience a starburst event at the rate $\\sim5\\msun\\yr^{-1}$ for a sustained period of few tens of Myrs. Such starbursts would produce a multiphase ISM in which the volume fraction of the hot gas exceeds $50\\%$ in the midplane of the clump. Hot bubbles created by feedback expand both in the horizontal direction to fill the interclump material, and vertically to form chimney-like structures which lead to high-velocity winds. Shells and neutral gas fragments embedded in such winds $1-3\\kpc$ above the disk give rise to neutral gas absorption with the median velocity $\\vmed\\approx90\\kms$ lasting roughly for the entire duration of the starburst.  Similar to other multiphase disk models, the cold component of the ISM is distributed in a complex network of filamentary structures confined by hot bubbles and voids. Inside these filaments dense clouds form as gravitational instabilities the growth of which is assisted by the external pressure. Once their growth begins, the clouds proceed to form stars very quickly. Only a small fraction of each cloud ($\\sim1-2\\%$ by mass) is converted into stars, as the clouds are being continuously destroyed by stellar winds from inside, interaction with shells pushed by rapidly expanding nearby hot bubbles, and collisions with other clouds. In our simulation clouds are transient objects constantly exchanging mass and energy with the cold shells and filaments, as well as with the hot gas, and rarely surviving as discrete entities for longer than a fraction of the rotational period.   In this paper we chose the typical clump masses and initial conditions representative of protogalactic environments at $z=3$. We expect similar winds to arise in grid-based cosmological simulations at $\\sim10\\pc$ spatial resolution. Although none of the current cosmological models have such resolution, we argue that using AMR techniques to zoom in only on those protogalactic clumps that have a high gas infall rate will allow to obtain galactic winds from thermal feedback in the cosmological context, without suppressing cooling or using kinetic feedback."
    },
    "0808/0808.0195_arXiv.txt": {
        "abstract": "We study the central dark matter (DM) cusp evolution in cosmologically grown galactic halos. Numerical models with and without baryons (baryons$+$DM, hereafter BDM model, and pure DM, PDM model, respectively) are advanced from identical initial conditions, obtained using the Constrained Realization method. The DM cusp properties are contrasted by a direct comparison of pure DM and baryonic models. We find a divergent evolution between the PDM  and BDM models within the inner few$\\times 10$~kpc region. The PDM model forms a $R^{-1}$ cusp as expected, while the DM in the BDM model forms a larger isothermal cusp $R^{-2}$ instead. The isothermal cusp is stable until $z\\sim 1$ when it gradually levels off. This leveling proceeds from inside out and the final density slope is shallower than $-1$ within the central 3~kpc (i.e., expected size of the $R^{-1}$ cusp), tending to a flat core within $\\sim 2$~kpc. This effect cannot be explained by a finite resolution of our code which produces only a 5\\% difference between the gravitationally softened force and the exact Newtonian force of point masses at 1~kpc from the center. Neither is it related to the energy feedback from stellar evolution or angular momentum transfer from the bar. Instead it can be associated with the action of DM$+$baryon subhalos heating up the cusp region via dynamical friction and forcing the DM in the cusp to flow out and to `cool' down. The process described here is not limited to low $z$ and can be efficient at intermediate and even high $z$. ",
        "introduction": "\\label{sec:intro}  Cosmological numerical simulations of pure dark matter (DM) halo formation have led to ``universal\" density profile which can be approximated by $\\rho(r) = 4 \\rho_{\\rm s}/(r/R_{\\rm s})(1+r/R_{\\rm s})^2$ (Navarro et al. 1997, NFW), with \\rs\\ being a characteristic ``inner\" radius where logarithmic density slope is $-2$ and \\rhos\\ is the density at \\rs. This profile steepens outside \\rs\\ and tends to a cusp $\\sim R^{-1}$ toward the center, at $R_{\\rm cusp}\\sim 0.1 R_{\\rm s}$. While this profile remains invariant under mergers (El-Zant 2008), its origin is yet to be explained (e.g., Syer \\& White 1998; Nusser \\& Sheth 1999; Shapiro et al. 2004; Ascasibar et al. 2007)  The NFW density profile is a source of an ongoing controversy --- its universality in the numerical pure DM models is in apparent contradiction with at least some of the observations of disk galaxies and galaxy clusters, which exhibit rather flat density cores (e.g., Flores \\& Primack 1994; Kravtsov et al. 1998; Salucci \\& Burkert 2000; Sand et al. 2002; de Blok et al. 2003; de Blok 2005, 2007; Simon et al. 2003; but see Rhee et al. 2004; answered by Gentile et al. 2004, 2005; de Naray et al. 2008). While a number of `exotic' explanations involving less conventional physics have been proposed, attempts to resolve this discrepancy within the CDM cosmology have been made as well. The DM cusp leveling off was attributed to the DM-baryon interactions, such as a dynamical friction of DM/baryon inhomogeneities (i.e., substructure) against the DM halo background (El-Zant et al. 2001, 2004; Tonini et al. 2006), stellar bar--DM interaction (Weinberg \\& Katz 2002; Holley-Bockelmann et al. 2005; but see Sellwood 2003; McMillan \\& Dehnen 2005; Dubinski et al. 2008), and baryon energy feedback (e.g., Mashchenko et al. 2006; Peirani et al. 2008). In this Letter, we test the first option (El-Zant et al.) --- the effect of the baryon$+$DM substructure in the fully self-consistent numerical simulations of halo formation in the $\\Lambda$CDM cosmology with WMAP3 paramaters (more details in Romano-D\\'{\\i}az et al. 2008).  Early arguments about change of the central density profile claimed that baryons drag DM in the so-called adiabatic contraction, steepening the DM density slope (e.g., Blumenthal et al. 1986). However, they have neglected the clumpy nature of the accreting material in the assembling halo, which can give rise to clump-background dynamical friction and energy transfer from the clumps to the background. The ability of baryons to radiate their internal energy increases the binding energy and acts as a `glue' on the DM substructure (El-Zant et al. 2001). El-Zant et al. have used the Monte-Carlo technique to calculate the leveling of the NFW cusp in the presence of such clumps and found this process to be efficient over $\\sim 1-2$~Gyr. The initial conditions presumed the existence of a NFW cusp, a spherically-symmetric DM halo and indestructible clumps.  \\begin{figure*}[ht!!!!!!!!!!!!!!!!!!!] \\begin{center} \\includegraphics[angle=0,scale=0.28]{fig1a.ps} \\includegraphics[angle=0,scale=0.28]{fig1b.ps} \\includegraphics[angle=0,scale=0.28]{fig1c.ps} \\end{center} \\caption{Redshift evolution of DM density profiles, $\\rho(R)$ (left), $\\rho(R) R$ (middle) and $\\rho(R) R^2$ (right) in PDM and BDM models: $z=3.55$ (solid), 2.12 (dotted), 1.0 (dashed), 0.61 (dot-dashed), 025 (dot-dash-dotted) and 0 (long dashed). The PDM and BDM curves are displaced vertically for clarity. The inner 40~kpc of halos are shown. The vertical coordinate units are logarithmic and arbitrary. For the PDM model, the density is well fitted by the NFW profile over a large range in $z$, and $R_{\\rm s}\\sim 28$~kpc at $z=0$. For the BDM model, the NFW fit is worse and $R_{\\rm iso}\\sim 15$~kpc at the end. The insert provides $\\rho$ within 200~kpc range for a comparison. } \\end{figure*}   The NFW cusp is characterized by the ``temperature\" (i.e., DM dispersion velocities) inversion which makes it thermodynamically improbable but dynamically stable in the absence of energy transfer mechanism (El-Zant et al. 2001). Dynamical friction of sufficiently bound clumps can trigger such energy flux into the cusp, increasing its dispersion velocities and washing it out. El-Zant et al. found that the necessary condition for the dynamical friction to have an effect is the aggregation of mass into the clumps more massive than $\\sim 10^{-4}$ of the DM halo --- result confirmed for the galaxy clusters as well (El-Zant et al. 2004).  In fully self-consistent numerical simulations with baryons, the question is whether the NFW cusp forms in the first place. Such simulations have been either limited to high redshifts ($z=3.3$, Gnedin et al. 2004) or focused on the adiabatic contraction and had an insufficient resolution (Gustafsson et al. 2006). They have verified that baryons lead to a steeper DM density profile than the NFW cusp, but left open its subsequent evolution. This issue is addressed here. ",
        "conclusions": "We have performed a direct comparison between the DM halo evolution in the pure DM and baryonic models, from identical cosmological initial conditions. Our goal is to understand the effect of the baryons on the DM density profile. We, therefore, focus on the DM evolution within the central region, quantify the cuspiness of the density distribution there, and attempt to relate its evolution to that of the DM substructure. Our results show that the DM in the PDM model has established a NFW cusp early, and the only subsequent changes we observe are related to sudden increases in \\rs\\ at the time of the major mergers (e.g., Romano-D\\'{\\i}az et al. 2006, 2007). On the other hand, the baryonic (BDM) model goes through a two-step process: first, in the central $\\sim 10$~kpc, the DM is dragged in by the baryons in what fits the adiabatic contraction (e.g., Gnedin et al. 2004; Gustafsson et al. 2006). This stage is concurrent with the formation and growth of a galactic disk, but we did not verify whether the disk is solely responsible for the DM contraction. For an extended period of time, after $z\\sim 4$, an isothermal cusp forms with $\\rho_{\\rm DM}\\sim R^{-2}$ and persists. In the next step, the DM cusp is being gradually washed out from inside out, mainly after $z\\sim 1$. By $z=0$, the isothermal cusp is largely erased, and the density slope within the central 10~kpc is less steep than $-2$. In fact within the central $\\sim 3$~kpc, it is less steep than the NFW slope of $-1$. We note, that the {\\it DM in the BDM model does not form the NFW cusp in the first place, but rather by-passes it in favor of an isothermal DM cusp which is subsequently washed out.} For example, the `temperature' inversion characteristic of the NFW cusp is never observed in the BDM model (Fig.~2).   We first argue, that this leveling of the DM cusp is not a numerical artifact, i.e., is not caused by a finite resolution of the code which is $2\\epsilon_{\\rm grav}=1$~kpc. At this distance, softened and exact Newtonian forces between point masses differ by 5\\%  --- this cannot account for the observed changes in density profile further out. We also note that the DM density peak in the BDM model is occasionally offset from the baryonic peak by a few kpc, around $z\\sim 1$ and especially 0.4, with the subhalos influx into the center, in what can be treated as $m=1$ instability. Therefore, we have fitted the DM halo profile based on the position of the DM density peak at each frame. The cusp will be `smeared' if this procedure is not followed (McMillan \\& Dehnen 2005). Also, the prime halo moves with respect to the center of mass of the computational sphere, this is related to the residual large-scale streaming motions in the DM (Romano-D\\'{\\i}az et al. 2008).  \\begin{figure}[ht!!!!!!!!!!!!!!!!!!!!!!!!!!!!!] \\begin{center} \\includegraphics[angle=0,scale=0.56]{fig4.ps} \\end{center} \\caption{Evolution of the subhalo number $N_{\\rm sbh}$ within central 30~kpc of the prime halo with redshift --- PDM (lower blue) and BDM (upper red). Note the sharp increase in $N_{\\rm sbh}$ for the BDM model after $z\\sim 1$, 0.5 and 0.2. } \\end{figure}  Next, we note that most of the cusp leveling happens after $z\\sim 1$ (Fig.~1). This correlates with a decrease in the slope of the central dispersion velocities (Fig.~2). The slope change in $\\beta$ (Fig.~3) comes from the decrease in the central dispersion velocities around few kpc, while $\\sigma_{\\rm DM}$ at $R\\sim 10$~kpc remains unchanged. What is the reason for this apparent `cooling' of the isothermal DM cusp?  Any decrease in central $\\sigma_{\\rm DM}$ of a dissipationless self-gravitationg entity in a virial equilibrium is counter-intuitive. It clearly has nothing to do with the major mergers, as this epoch ends at $z\\sim 1.5$ in our models. Any subsequent {\\it smooth} accretion of dissipative baryons will only increase $\\sigma_{\\rm DM}$ as a result of the ongoing adiabatic contraction. However, if one considers accretion of {\\it clumpy} baryons, which act as a `glue' on the DM subhalos where they reside, additional process of `heating' the DM by means of a dynamical friction becomes important. The heating rate is skewed to the density peak and will cause the DM to stream out of the cusp's gravitational well, eventually decreasing the `temperature' gradient at the center. Fig.~3 displays this behavior of $\\beta$ after it reaches its peak at $z\\sim 0.7$. For a comparison, the PDM halo shows a flat $\\beta$ after the epoch of major mergers, for nearly 10~Gyr.  Central region crossing by subhalos (as well as accretion and tidal disruption) continues to the present time. Why do changes in the cusp become significant after $z\\sim 1$? Fig.~4 displays the evolution of the subhalo population within the central 30~kpc of the BDM halo. We observe that at least in two instances, $z\\sim 1$ and 0.4, there is a clear excess in the BDM subhalo population, corresponding to the splash in the subhalo influx rates. The number of these subhalos also exceeds that of the PDM by $\\sim 2$. The BDM subhalos appear in waves within the central region and these waves coinside with the DM streaming out of the center and `cooling' down. Hence, the process of washing out the DM cusp correlates with the influx of subhalos into the innermost region of the BDM halo. We show elsewhere that waves of subhalos crossing the central region of the prime halo originate in the filament, cluster there and enter the prime halo before they merge among themselves.  We find that the mass distribution of subhalos (in the computational sphere) with masses $M_{\\rm sbh}$ evolves with redshift --- at high $z\\gtorder 4$ it can be approximated by a power law, $N_{\\rm sbh}\\sim M_{\\rm sbh}^{-1}$, both in PDM and BDM models (Romano-Diaz et al. 2008). But in the central virialized 100~kpc of the halo it evolves differently from the field subhalos. Specifically, it is heavily skewed toward more massive ones after $z\\sim 1$, while that of the PDM subhalos is only weakly so. This is significant, because the efficiency of dynamical friction increases substantially with the clump-to-prime halo virial mass ratio (El-Zant et al. 2001). At $z\\sim 4$, this ratio for the most massive subhalo residing within the virialized halo is $\\sim 0.06$. By $z\\sim 1$ this ratio is $\\sim 10^{-3}$, and at $z=0$ it is $\\sim 8\\times 10^{-4}$ --- still higher than the minimum required by El-Zant et al.  The size of the obtained core in the DM distribution of the BDM model can be compared with observationally inferred cores. Following Salucci et al. (2007), for the virial mass of our halo at $z=0$, \\mvir $\\sim 4\\times 10^{12}~\\msun$, the expected core is $\\sim 1.6$~kpc. Spano et al. (2008) lists a large range of core sizes, starting from $\\sim 1$~kpc. Both are compatible with the value obtained here.  In summary, we find that the DM density cusp corresponding to $R_{\\rm cusp}$ in the PDM model is leveled off by the action of subhalos in the presence of baryons in the BDM model. The energy feedback from stellar evolution has decayed by this time without any effect on the cusp. We do not detect the angular momentum ($J$) transfer from the bar to the cusp --- the latter $J$ stays constant in time. We also find that the BDM model forms a steeper isothermal rather than NFW-type cusp for extended time period. It goes through a two step process which involves gravitational contraction, followed by the influx of subhalos which heat up and dissolve the cusp.  Our models, although containing about $1.6\\times 10^6$ particles per halo, are still insufficient to fully resolve the formation of substructure around the prime halo. They should be viewed rather as a lower limit on the efficiency of the process reported here. Neither is the process described here limited to low $z$ --- one can easily envision scenario(s) when the subhalos act at intermediate and even high $z$. Overall, this work stands in agreement with Gnedin et al. (2004) who stopped the simulation at redshift $z=3.3$, well before the leveling of the density cusp. However, Gustafsson et al. (2006) barely lacked resolution to detect the cusp behavior observed here. We have recently learned that J.P. Ostriker (priv. comm.) obtained similar results on the cusp flattening."
    },
    "0808/0808.0706_arXiv.txt": {
        "abstract": " ",
        "introduction": "An important class of inflationary models \\cite{Inflation},  chaotic inflation \\cite{Linde:1983gd}, involves an inflaton field excursion that is large compared to the Planck scale $M_P$ \\cite{Lyth}.  These models have a GUT-scale inflaton potential, and are accessible to observational tests via a B-mode polarization signature in the CMB \\cite{Bmodes, Bmodeobs}.  The Planckian or super-Planckian field excursions required for high-scale inflation may be formally protected by an approximate shift symmetry in effective field theory. A canonical class of examples with a field excursion $\\Delta\\Phi\\simeq M_P$, known as Natural Inflation, employs a pseudo-Nambu-Goldstone boson mode (an axion) as the inflaton \\cite{Freese:1990rb,OtherNatural}.  Because inflation is sensitive to Planck-suppressed operators, however, it is still of significant interest to go beyond effective field theory and realize inflation in string theory, a candidate ultraviolet completion of gravity. Conversely, CMB observations which discriminate among different inflationary mechanisms provide an opportunity to probe some basic features of the ultraviolet completion of gravity.  The lightest scalar fields in string compactifications roughly divide into radial and angular moduli. Radial moduli, such as the dilaton and the compactification volume, have an unbounded field range as they go toward weak-coupling limits.  In these limits their contributions to the potential are typically very steep, not sourcing large-field inflation in any example yet studied. Angular moduli, such as axions, have potentials that are classically protected by shift symmetries. However, in the case of axions it has been argued that the field range contained within a single period is generally sub-Planckian in string theory \\cite{Banks:2003sx}, leading to proposals to extend the field range by combining many axions \\cite{Nflation}.  In the present work, we show that in the presence of suitable wrapped branes, the potential energy is no longer a periodic function of the axion.  When this {\\it monodromy} in the moduli space is taken into account, a single axion develops a kinematically unbounded field range with a potential energy growing {\\it linearly} with the canonically normalized inflaton field.  This implements the monodromy mechanism introduced in \\cite{Silverstein:2008sg}\\ in a wide class of string compactifications.  Because the basic idea is very simple, let us indicate it here. Axions arise in string compactifications from integrating gauge potentials over nontrivial cycles.  For example, in type IIB string theory, there are axions $b_I=\\int_{\\Sigma_I^{(2)}}B$ arising from integrating the  Neveu-Schwarz (NS) two-form potential $B_{MN}$ over two-cycles $\\Sigma_I^{(2)}$, and similarly axions $c_I=\\int_{\\Sigma_I^{(2)}}C$ arise from the Ramond-Ramond (RR) two-form $C_{MN}$. In the absence of additional ingredients such as fluxes and space-filling wrapped branes, the potential for these axions is classically flat, and develops a periodic contribution from instanton effects.  A Dp-brane wrapping $\\Sigma_I^{(2)}$, on the other hand, carries a potential energy that is {\\it not} a periodic function of the axion: in fact, this energy increases without bound as $b_I$ increases.  The effective action for such a wrapped brane is the DBI action, given in terms of the embedding coordinates $X^M(\\xi)$ as s% \\be S_{DBI}=-\\int {d^{p+1}\\xi\\over{(2\\pi)^p}}\\alpha'^{-(p+1)/2}e^{-\\Phi}\\sqrt{\\det{(G_{MN}+B_{MN})\\, \\del_\\alpha X^M\\del_\\beta X^N}} \\label{DBIgen} \\ee where we have omitted the corresponding Chern-Simons term, which will be unimportant for our considerations. A key example is a D5-brane wrapped on a two-cycle $\\Sigma^{(2)}$ of size $\\ell\\sqrt{\\alpha'}$, which yields a potential \\be V(b) = {\\epsilon\\over{g_s(2\\pi)^5\\alpha'^2}}\\sqrt{\\ell^4+b^2} \\label{Bpot} \\ee that is linear in the axion field $b$ at large $b$. (Here we have included a factor $\\epsilon$ to represent the effects of warping, which we describe more carefully below.) Similarly, an NS5-brane wrapped on $\\Sigma_I^{(2)}$ introduces a monodromy in the $c_I$ direction.  Monodromy is a common phenomenon in string compactifications.  In the past, it has been studied extensively in the context of particle states in field theory \\cite{Witten:1979ey}\\ and the corresponding non-space-filling wrapped branes of string theory \\cite{wrappedmonod}. The present case of monodromy in the {\\it potential energy} arises when a would-be periodic direction $\\gamma$ is ``unwrapped'' by the inclusion of an additional space-filling ingredient whose potential energy grows as one moves in the $\\gamma$ direction, extending the kinematic range of the corresponding scalar field. Because the wrapped branes are space-filling, their charge must be cancelled within the compactification.  We will do so with an antibrane wrapped on a distant, homologous two-cycle as depicted in Fig. 2 in \\S4\\ below.  In the bulk of this paper, we analyze the conditions under which this yields controlled large-field inflation in string theory. We find a reasonably natural class of viable models. As is usually the case in inflationary model building from string theory, much of the challenge is to gain systematic control of Planck-suppressed corrections to the effective action.  After ensuring that our candidate inflaton potential does not destabilize the compactification moduli, and that fluxes do not affect the structure of our candidate inflaton potential, we establish that instanton effects, which produce sinusoidal contributions to the axion potential, can be naturally suppressed. We assess these conditions for both perturbative and nonperturbative stabilization mechanisms, drawing examples based both on Calabi-Yau compactifications and on more general compactifications that break supersymmetry at the Kaluza-Klein scale.  In the case of nonperturbative stabilization mechanisms in type IIB string theory, we find a controlled set of models for the RR two-form axions $c_I$, while perturbative stabilization mechanisms suggest opportunities for inflating in the $b_I$ as well as in the $c_I$ directions. These varied implementations of our axion monodromy mechanism give identical predictions for the overall tilt and tensor to scalar ratio in the CMB, as they are all well-described by a linear potential for a canonically-normalized inflaton.\\footnote{There may also be novel signatures from finer details of the power spectrum originating in the repeated circuits of the fundamental axion period, as we discuss further below.} Our prediction for these quantities lies well within the exclusion contours from present data \\cite{Observations}, and is ultimately distinguishable from the predictions of other canonical models via planned CMB experiments \\cite{Bmodeobs,Planck} (see Fig.~3).  Our mechanism relies on specific additional ingredients -- branes -- intrinsic in the ultraviolet completion of gravity afforded by string theory.  Although string theory restricts the range of the original axion period in the first place, it then recycles a single axion via monodromy, providing a simple generalization of \\cite{Linde:1983gd,Freese:1990rb}\\ with its own distinctive predictions.  The subject of axion inflation has thus almost come full circle.\\footnote{Though we hope to have added something to the subject this time around.} ",
        "conclusions": "Monodromy is a generic phenomenon in string compactifications.  We have shown that the axion monodromy introduced by space-filling wrapped fivebranes leads to a linear potential, over a super-Planckian distance, for the canonically normalized axion field.  Axion monodromy therefore provides a mechanism for realizing chaotic inflation, with a linear potential, in string theory.   We have shown that this mechanism is compatible with various methods of moduli stabilization, including nonperturbative stabilization of type IIB string theory on warped Calabi-Yau manifolds, as well as perturbative stabilization on more general Ricci-curved spaces. This produces a clear signature in the CMB of $r\\approx 0.07$, with model-dependent opportunities for further, novel, signatures arising from oscillating corrections to the slow-roll parameters.   Our mechanism is reasonably robust and natural because of the presence of perturbative axionic shift symmetries. In our examples, the inflaton potential itself is the leading effect that breaks the shift symmetry, with instanton corrections naturally exponentially suppressed. We would like to remark that a related symmetry structure is plausibly also present in configurations with more general monodromies not involving axions. Monodromy in the potential energy arises when a would-be circle in a direction $\\gamma$ in the approximate moduli space is lifted by an additional ingredient whose potential energy grows as one moves in the $\\gamma$ direction.  This unwraps the circle direction and extends the kinematic range of the corresponding field.  Then, symmetries translating around the original circle do much to control the structure of the potential along the eventually-unwrapped direction. Thus, monodromy-extended directions are not just long; they also generically profit from approximate symmetries.  Monodromy-extended directions can be used for large-field inflation if the underlying moduli potential depends sufficiently weakly on $\\gamma$ and if all corrections to the slow-roll parameters are sufficiently suppressed; the classes of compactifications analyzed here and in \\cite{Silverstein:2008sg}\\ provide two particular realizations of this effect.\\footnote{Recently, monodromy has been used as a method to model Chain Inflation \\cite{ChainInflation}\\ in string theory \\cite{chainmonod}.}  There is much more to be done at the level of model building. The examples we have provided in this work are useful as proofs of principle, and to that end we have focused on demonstrating parametric suppression of corrections to the inflaton potential, and in particular on enumerating a wide array of mechanisms, such as warping, axionic symmetries, extended local supersymmetry, etc., that serve to control such contributions.  We have not yet attempted to construct a minimal realization of linear axion inflation that uses the smallest possible subset of these control mechanisms.  This is an interesting problem for future work, as methods for analyzing string compactifications and string-theoretic instantons improve.  A further lesson of this work, as of \\cite{Silverstein:2008sg}, is that in large-field models based on monodromy, a degree of suppression of otherwise problematic contributions to the potential that suffices for inflation is also sufficient to make firm predictions for the tilt of the scalar power spectrum.  This is in sharp contrast to typical small-field models, where fine-tuning the inflaton potential to be flat enough for inflation is not a strong enough restriction to be predictive: slight variations in the fine-tuned contributions can noticeably change the tilt.  The difference in our case is that the problematic terms arise as periodic modulations of the potential; requiring that inflation occurs at all implies that the amplitude of these modulations is small compared to the scale of changes in the inflaton potential itself.  In turn, this implies that the {\\it average} tilt is not affected at a detectable level by these modulations.  On the other hand, it would be very interesting if the oscillations in the detailed power spectrum produced by a modulated linear potential had characteristic features accessible to future observations.  In any case, this class of models is falsifiable on the basis of its gravity wave signature."
    },
    "0808/0808.3125_arXiv.txt": {
        "abstract": "Dark energy models with a single scalar field cannot cross the equation of state divide set by a cosmological constant. More general models that allow crossing require additional degrees of freedom  to ensure gravitational stability. We show that a parameterized post-Friedmann description of cosmic acceleration provides a simple but accurate description of multiple scalar field crossing models.  Moreover the prescription provides a well controlled approximation for a wide range of ``smooth\" dark energy models.  It conserves energy and momentum and is exact in the metric evolution on scales well above and below the transition scale to relative smoothness.  Standard linear perturbation tools have been altered to include this description and made publicly available for studies of the dark energy involving cosmological structure out to the horizon scale. ",
        "introduction": "\\label{sec:intro}  Observational constraints on the acceleration of the expansion have continued to close in on a dark energy equation of state of a cosmological constant $w_e=-1$ that delineates the phantom divide. Testing the small deviations from that value in the future requires a description of the dark energy that allows the equation of state to evolve across the phantom divide possibly multiple times.  It is well known that single scalar fields are gravitationally unstable to such a crossing of the phantom divide \\cite{Vik04,Hu04c,CalDor05}. Dark energy that is minimally coupled to the matter requires additional degrees of freedom to cross the divide stably.  While specific models with multiple fields can be constructed \\cite{Hu04c,Feng:2004ad} they are cumbersome or impossible to implement in a general analysis of the dark energy.  The usual approach in the literature for finessing such cases is to artificially turn off the dark energy perturbations explicitly or implicitly by limiting the range of observables.  Doing so violates energy-momentum conservation whenever $w_e \\ne -1$ and leads to inconsistencies between the Einstein equations for the evolution of the metric due to the Bianchi identities which can persist even on small scales. Though the impact of perturbations tend  to be small near $w_e=-1$, cosmological constraints often require the exploration of a large swath of parameter space around the maximum likelihood.  Excising the instability around the transition provides another, albeit rather ad hoc approach \\cite{Hue05}.    In this {\\it Brief Report}, we show that the so-called parameterized post-Friedmann (PPF) approach to describing linear metric evolution in a Friedmann-Robertson-Walker (FRW) universe  provides a simple solution to this dilemma.  The PPF framework was ostensibly introduced for describing modified gravity theories under a metric framework with strict local conservation of energy and momentum \\cite{HuSaw07b}.  As such it also applies to dark energy models \\cite{Hu08a} and in particular the class of models which have  a well-defined Jeans scale under which the dark energy is smooth compared to the dark matter. This framework also has the benefit of being an exact description for the metric evolution well above and well below this scale and hence provides a very well-controlled approximation that is simple to implement in an Einstein-Boltzmann linear perturbation code. ",
        "conclusions": "We have shown that the PPF prescription for describing the evolution of metric perturbations in an FRW universe is sufficiently general to encompass multiple scalar field models whose joint equation of state evolves across the phantom divide at $w_e=-1$. This description is accurate to well below the cosmic variance limit as long as the transition scale to relative smoothness is comparable to the horizon.  Moreover it  is in fact exact for the metric evolution well above and well below the transition scale.  As such it provides a well-controlled approximation for any model where the energy and momentum of the dark energy is separately conserved and features a transition of this type.  This prescription is useful for the joint analysis of growth and distance measures of the dark energy, especially those involving horizon scale perturbations like the integrated Sachs-Wolfe effect in the CMB.  The CAMB Einstein-Boltzmann package has been altered to include PPF  \\cite{Fangetal08} and a version for the dark energy has been made publically available \\footnote{\\url{http://camb.info/ppf/}}. Potential future uses include principal component approaches to dark energy constraints where $w_e$ is allowed to cross the phantom divide multiple times (e.g.~\\cite{HutSta03}).  Here explicit matching to multiple scalar fields is cumbersome if not  impossible.  The PPF prescription provides a simple but general approach that explicitly enforces conservation of energy and momentum and all of the Einstein equations removing potential ambiguities to the meaning of a ``smooth\" dark energy component.   \\smallskip  WF is supported by the Initiatives in Science and Engineering program at Columbia University and by DOE contract DE-FG02-92ER-40699; WH by DOE contract DE-FG02-90ER-40560, the David and Lucile Packard Foundation and the KICP under NSF PHY-0114422; AL by an STFC Advanced Fellowship.  \\vfill"
    },
    "0808/0808.3580_arXiv.txt": {
        "abstract": "{We present 1D numerical simulations aimed at studying the hot-flasher scenario for the formation of He-rich subdwarf stars.  Sequences were calculated for a wide range of metallicities and physical assumptions, such as the stellar mass at the moment of the helium core flash. This allows us to study the two previously proposed flavors of the hot-flasher scenario (``deep'' and ``shallow'' mixing cases) and to identify a third transition type.  Our sequences are calculated by solving simultaneously the mixing and burning equations within a diffusive convection picture, and in the context of standard mixing length theory.  We are able to follow chemical evolution during deep-mixing events in which hydrogen is burned violently, and therefore able to present a homogeneous set of abundances for different metallicities and varieties of hot-flashers.  We extend the scope of our work by analyzing the effects of non-standard assumptions, such as the effect of chemical gradients, extra-mixing at convective boundaries, possible reduction in convective velocities, or the interplay between difussion and mass loss. Particular emphasis is placed on the predicted surface properties of the models.   We find that the hot-flasher scenario is a viable explanation for the formation and surface properties of He-sdO stars. Our results also show that, during the early He-core burning stage, element diffusion may produce the transformation of (post hot-flasher) He-rich atmospheres into He-deficient ones. If this is so, then we find that He-sdO stars should be the progenitors of some of the hottest sdB stars. } ",
        "introduction": "Hot subluminous stars are an important population of faint blue stars that can roughly be grouped into the cooler sdB stars, whose spectra display no or only weak helium lines, and the hotter sdO stars, which have higher helium abundances and can even be dominated by helium. Subluminous O, B stars were identified with stars populating the hot end of the horizontal branch (HB) of some globular and open clusters (de Boer et al. 1997). While the sdB stars form a homogeneous spectroscopic class, a large variety of spectra is observed among sdO stars (Lemke et al.  1997). This diversity of spectra and the helium-enhanced surface abundances observed in many of these stars pose a challenge to our understanding of stellar evolution. As a consequence, some non-canonical evolutionary scenarios were proposed for their formation.  Among them, the merger of two white dwarfs (Saio \\& Jeffery 2000) and a core helium flash after departure from the red giant branch (i.e. ``hot-flasher scenario'') offer the most promising explanations of their formation (Str\\\"oer et al. 2007, Moehler et al. 2007, and Hirsch et al. 2008). The second possibility is the subject of the present article. The hot-flasher scenario was proposed to explain the existence and characteristics of blue hook stars (such as helium or carbon enhancement) of some globular clusters e.g. $\\omega$Cen and NGC 2808 (D'Cruz et al. 2000, Moehler et al. 2002, Moehler et al. 2004, and Moehler et al. 2007). The coexistence of extreme horizontal branch (EHB) stars and helium core white dwarfs in globular clusters (Calamida et al. 2008) also agrees with a natural prediction of the hot-flasher scenario (see Castellani et al.  2006). Although in this scenario mass loss must be increased artificially, from values usually adopted in the modeling of low mass AGB stars, the driving mechanism behind RGB mass loss is not well understood (Espey \\& Crowley 2008). Differences in metallicity, initial He abundances or rotational velocities may, in principle, produce enhanced mass loss rates (the last two by an increase in the RGB-tip luminosity). The later is particularly interesting in view of the fact that some He-sdO stars are known to be relatively fast rotators (in comparison with sdB stars; Hirsch et al. 2008).  The existence of high mass-loss rates in old metal-rich populations is also an important ingredient of one of the most successful hypothesis for modeling the UV-upturn of elliptical galaxies (Yi \\& Yoon 2004).  If this is so, it is very probable that a fraction of stars in these populations will undergo the helium core flash (HeCF) after departing from the first red giant branch (RGB) --- i.e. will be ``hot-flashers'' as coined by D'Cruz et al. (1996).  Finally, in high stellar density enviroments, it is possible that enhanced mass loss might result from star-to-star encounters.   In general, low mass stars undergo the HeCF at the tip of the RGB. However, Castellani \\& Castellani (1993) showed that, if sufficient mass is lost on the RGB, the star can depart from the RGB and experience the HeCF while descending the hot white dwarf cooling track. Afterwards D'Cruz et al. (1996) analyzed the frequency of these HeCFs at high temperatures (``Hot Flashers'') by assuming different values of the mass-loss efficiency and suggested that this scenario might provide a path for populating the hot end of the HB.  Sweigart (1997) noticed the similarity between these hot-flashers and the late helium shell flashes that produce ``born-again AGB stars'' (Iben 1984) and showed that --- as happens in the born-again scenario --- the convective zone, generated by the extremely high energy liberation rate ($L_{\\rm He}\\sim 10^{10}L_\\odot$) during the flash, reaches the H-rich envelope and engulfs it. As a consequence, much or almost all H is burned leading to a He-enriched surface abundance. Brown et al. (2001) explored this scenario to explain anomalies in the hot HB of NGC 2808. However, it was not until the work of Cassisi et al. (2003) that simulations including the violent mixing and burning of H were performed and yields for the abundances of helium, carbon and nitrogen were calculated consistently.  All of these works were focused mainly on globular clusters and therefore on low metallicity progenitors. However, counterparts of these hot-flashers are also expected in the Galactic disk. This was considered by Lanz et al. (2004) who assessed the possibility that (field) He-sdB stars could be formed by the hot-flasher scenario. In this work, the authors classified three different clases of hot-flashers:(1) early hot-flashers, which experience the HeCF during the evolution at constant luminosity in the HR diagram (after departing from the RGB-tip) and become hot subdwarf stars with standard H/He envelopes; (2) late hot-flashers with shallow-mixing (SM), which become He-enriched hot subdwarfs due to convective dilution of the envelope (at the low-luminosity, low temperature region of the HR diagram) into deeper regions of the star; and (3) late hot flashers with deep-mixing (DM) in which the H-rich envelope is engulfed and burned in the convective zone generated by the primary HeCF.  Due to the numerical difficulties related to the calculation of the violent H burning that occurs during a deep-mixing event (see Cassisi et al. 2003), Lanz et al. (2004) failed to provide surface abundances in this case. Finally, Str\\\"oer et al. (2005, 2007) and Hirsch et al. (2008) analyzed the possibility that field He-sdO could be generated by this mechanism and concluded that the idea could not be rejected, although some serious discrepancies exist between models and observations. However, the lack of a quantitative study of hot-flasher events and the chemical composition that arise from them for a wide range of metallicities, ensures that comparison with observations remains difficult. We note that only one previous work (Cassisi et al. 2003) has calculated the abundances produced by the deep-mixing event but only for one remnant mass and metallicity (Z=0.0015). On the other hand Lanz et al. (2004) calculated several sequences for solar-like metallicities but no calculation of the final surface abundances for the deep-mixing events were presented.  The amount of observational data for these stars is increasing rapidly as a byproduct of surveys of faint blue objects such as SDSS and SPY. As concluded by Cassisi et al. (2003) to assess the viability of the hot-flasher scenario, a study of the dependence of the hot-flasher predictions ($T_{\\rm eff}$, $g$ and abundances) on the physical details is needed.  The core feature of this article is to present a homogeneous set of simulations of the hot-flasher scenario for a wide range of cases in which the chemical abundances of the models are consistently calculated.  In particular we present simulations for 4 different metallicities and different post-RGB remnant masses (which determines the luminosity of the model at the moment of the HeCF). The work is organized as follows. In the following section, we discuss the main ingredients of the simulations performed in this work. We then, in Sect. 3, discuss different flavors of hot-flasher and He-enrichment and describe our simulations results under standard assumptions. In Sect. 4, we analyze deviations from the standard assumptions such as possible extra-mixing at convective boundaries, the effect of chemical gradients, or a possible reduction in convective velocities. We also address the possible effects of element diffusion in the outer layers. In Sect. 5, we discuss the results of our standard sequences and compare them with observations. Finally, in Sect. 6, we present some concluding remarks. ",
        "conclusions": "In Sect. 3.3, we showed that standard sequences lead to final H abundances (see Table 3) immediately after a DM event that are much lower than inferred for He-sdO stars by Str\\\"oer et al. (2007). Only a minority of standard sequences (those which undergo SM/SM$^*$ events) display abundances similar to He-sdO stars immediately after the flash.  Also, as we have already mentioned, the predicted location in the $g-T_{\\rm eff}$ would differ from that observed. In particular, since the EHB evolution is more than 40 times longer than the evolution from the primary HeCF to the ZAHB, most of the stars should be clustered around the EHB, which is not observed. However, there are some viable scenarios in which our results are consistent with observations. These scenarios depend on the existence or not of weak homogeneous winds in He-rich subdwarf stars as we now discuss.  As mentioned in Sect. 4.4, in a very recent study Unglaub (2008) showed that weak homogeneous winds are impossible in hot subdwarfs at high gravities. In this case, as shown in Fig. \\ref{fig:chem_t}, the low amount of H remaining in our standard sequences after a DM event is sufficient to convert a He-rich to a He-deficient surface composition in about $10^6$ yr. Within this picture, after the DM event the star first evolves towards the ZAHB star as a He-rich subdwarf. Then, when gravity has increased sufficiently, the star crosses the wind limit causing homogeneous winds to halt, then diffusion leads to an increase in the H abundance as the star evolves towards the ZAHB (both processes occur on timescales of the order of 1 Myr). We note in particular (Fig. \\ref{fig:chem_t}) that  in only a few thousand years the H abundance rises from its post-DM value to the values usually measured in He-sdO stars (Str\\\"oer et al. 2007) and thus the correct H-abundances are predicted. When the HB has been reached the evolutionary timescale becomes more than a factor of 40 longer (65 to 90 Myr, see tab. 2) and the star turns into a H-dominated sdB star. This possible evolutionary channel is sketched in Fig. \\ref{fig:channel}. In this case no clustering of He-sdO stars should be expected close to the EHB, in good agreement with the results of Str\\\"oer et al. (2007). It is interesting to note that the effective temperatures expected for these stars, when on the HB (and with a H-dominated atmosphere), agree with those of the hotter sdB stars (which also have higher He abundances than their cooler counterparts). The conversion of surface abundances with C$>$N into surfaces with N$>$C as H diffuses outwards, might also be related (see Figs. \\ref{fig:chem_t} and  \\ref{fig:channel}) to the lack of C-rich He-sdO stars with lower $N_{\\rm He}/N_{\\rm H}$ values in the study of Str\\\"oer et al. (2007). Within this scenario and in the absence of any other channel that creates hot-subdwarfs, the number ratio of He-sdO stars to hot ($T_{\\rm eff}\\gtrsim 33\\,500$) sdB stars should be of the order of 1/40 in a complete sample (although higher ratios can still be possible if some process that disturbs diffusion is present).  Once the star finishes the He-core burning stage (or EHB), it moves above the wind limit again (see Fig. \\ref{fig:channel}) and homogeneous winds again become possible. If this is so, since the H-rich layer is very thin, a new conversion into a He-sdO star might happen (now at higher surface gravities and temperatures). Finally, it is interesting to note in Fig. \\ref{fig:chem_t} that the C-abundance decrease rapidly by about a factor of 25 and the star changes its surface composition from C$>$N to C$<$N. This occurs while the star is still strongly H-deficient. We note that, at this point (between t=1000 and 10000 yr in Fig. \\ref{fig:chem_t}), the star displays surface abundances qualitatively similar (H$\\sim 2.5\\times 10^{-3}$, N$\\sim 2\\times 10^{-3}$ and C$<$N by mass fraction), although quantitatively different, to those inferred by Lanz et al. (2004) for the peculiar star LB 1766.  In the case of post-SM/SM$^*$ stars our simulation indicates that the atmosphere should turn into a H-dominated one in about 1000 yr from the moment in which homogeneous winds cease. Then, if homogeneous winds do not exist at high gravities, post-SM/SM$^*$ stars would only be He-rich while they are above the wind limit (see Fig. \\ref{fig:channel}). This is clearly not the case for most He-sdO stars in the sample of Str\\\"oer et al. (2007) but it does describe the observed surface properties of JL87 ---$T_{\\rm eff}= 29\\,000$ K and log$g=5.5$, Lanz et al. (2004); $T_{\\rm eff}= 26\\,000$ K and log$g=4.8$, Ahmad et al. (2007).   We emphasize that the possible conversion of He-sdO stars into hot sdB stars in the HB discussed in the previous paragraphs is  interesting, not only because it predicts the correct distribution in the log $g-T_{\\rm eff}$ diagram but also in view of the significantly higher He abundances of sdB stars with $T_{\\rm eff}\\gtrsim 33\\,500$K compared with their cooler counterparts, which might reflect differences in their previous evolution.     In contrast, if the effects of element diffusion were of minor importance for all He-rich subdwarfs, then only the abundances resulting from SM/SM$^*$ events would be consistent with the surface properties measured for He-sdO stars (Str\\\"oer et al. 2007), unless DM abundances differ strongly from those predicted by the standard MLT theory (see Sect. 4.1).  We note that the fraction of stars undergoing SM/SM$^*$ events corresponds to only a small fraction of hot-flashers in standard sequences. In this context, our results in Sect. 4.3 indicate that the role of chemical gradients should be analyzed to determine whether the fraction of shallow-mixing episodes might be increased by the effect of chemical gradients that tend to prevent the penetration of convection into the H-rich envelope.  The hot-flasher scenario therefore can qualitatively predict the correct surface properties and distribution of He-sdO stars. It is then interesting to note that the lack of close binaries (with periods of the order of days as observed for sdB stars) in He-sdO stars is also expected naturally within the hot-flasher scenario. Very close binaries with final periods of the order of days will indeed fill their Roche lobes before reaching the RGB-tip and never undergo the HeCF. Also, it is not clear if it would be possible to have hot-flashers within a common envelope system. In any case, in order to have a hot-flasher within a close binary system, a very fine tuning of the initial masses and periods seems to be needed. Consequently, post hot-flasher H-deficient stars should not be usual in these systems. The lack of close binary systems with He-sdO components should then not be taken as an argument against the hot-flasher and in favor of the merger scenario. In contrast, wide systems would be a much better way of distinguishing between both possible evolutionary origins for He-sdO stars, because they should be more probable within the hot-flasher picture."
    },
    "0808/0808.3255_arXiv.txt": {
        "abstract": "\\noindent We study the impact of cosmological parameters' uncertainties on estimates of the primordial NG parameter $\\fnl$ in local and equilateral models of non-Gaussianity. We show that propagating these errors increases the $\\fnl$ $1\\sigma$ uncertainty by a term ${\\delta \\fnlloc / \\fnlloc} \\simeq 16 \\%$ for WMAP and ${\\delta \\fnlloc / \\fnlloc} \\simeq 5 \\%$ for Planck in the local case, whereas for equilateral configurations the correction term are  ${\\delta \\fnleq / \\fnleq} \\simeq 14 \\%$ and ${\\delta \\fnleq / \\fnleq } \\simeq 4\\%$, respectively. If we assume for $\\fnlloc$ a central value $\\langle \\fnlloc \\rangle \\simeq 60$, according to recent WMAP 5-years estimates, we obtain for Planck a final correction $\\Delta \\fnlloc \\simeq 3$. Although not dramatic, this correction is at the level of the expected estimator uncertainty for Planck, and should then be taken into account when quoting the significance of an eventual future detection. In current estimates of $\\fnl$ the cosmological parameters are held fixed at their best-fit values. We finally note that the impact of uncertainties in the cosmological parameters on the final $\\fnl$ error bar would become totally negligible if the parameters were allowed to vary in the analysis and then marginalized over. ",
        "introduction": "\\noindent Research of primordial non-Gaussianity in the Cosmic Microwave Background (CMB) is an important field of cosmology today. In a recent work, Yadav and Wandelt \\cite{YadavWandelt} claim a $3 \\sigma$ detection of a large NG signal in the WMAP 3-years data. The following WMAP 5-years analysis produced similar results but showed a reduction of statistical significance from about $3$ to $2 \\sigma$ \\cite{WMAP5}. Yadav and Wandelt estimate, using WMAP 3-years data, is \\beq 27 \\leq \\fnlloc \\leq 147  \\;\\;  (95\\% \\, {\\rm c.l.}) \\; . \\eeq The WMAP 5-years estimate is: \\beq -9 \\leq \\fnlloc \\leq 111  \\;\\;  (95\\% \\, {\\rm c.l.}) \\; , \\eeq where $\\fnl$ is a dimensionless parameter defining the strength of primordial non-Gaussianity (NG) \\cite{NGreview}, and the superscript ''local'' indicates that we are considering a primordial NG curvature perturbation of the form \\beq \\fix = \\fiLx + \\fnlloc \\left(\\fiLxsquare - \\langle \\fiLxsquare \\rangle \\right) \\; , \\eeq where $\\fix$ is the total gravitational potential, $\\fiLx$ is the gravitational potential computed at the linear level and $\\langle\\cdots \\rangle$ stands for the ensemble average. The reason for calling this a local NG model relies in the fact that the NG part of the primordial curvature perturbation is a local functional of the Gaussian part. The local shape of NG arises from standard single-field slow-roll inflation \\cite{standard} as well as from alternative inflationary scenarios for the generation of primordial perturbations, like the curvaton \\cite{curvaton} or inhomogenous (pre)reheating models \\cite{gamma1}, or even from alternatives to inflation, such as ekpyrotic and cyclic models \\cite{ek}. Other models, such as DBI inflation \\cite{dbi} and ghost inflation \\cite{ghost}, predict a different kind of primordial NG, called \"equilateral\", because the three point function for this kind of NG is peaked on equilateral configurations, in which the lengths of the three wavevectors forming a triangle in Fourier space are equal \\cite{shape}. This second form of NG is characterised by the parameter $\\fnleq$, which defines the amplitude of the equilateral triangles. In the following we will sometimes write the amplitude of NG simply as $\\fnl$, without any superscript. When we do so, we mean that our conclusions apply to both the local and the equilateral case, with no need for distinction.  Standard single field inflation predicts $\\fnlloc \\sim 10^{-2}$ at the end of inflation \\cite{standard}  (and therefore a final value $\\fnlloc \\sim $ unity  after general relativistic second-order perturbation effetcs are taken into account \\cite{second}). It is thus clear that large central values of $\\fnlloc$, like those obtained in the above mentioned analyses, are going to rule out the simplest scenarios of inflation as viable models of the Early Universe. On the other hand, the low statistical significance of the final WMAP 5-years result makes any conclusion premature at this stage. With its high angular resolution and sensitivity Planck will allow to significantly improve the statistical estimate of $\\fnlloc$, reducing the final error bars from the present $\\Delta \\fnlloc \\simeq 30$ to a final value of $\\Delta \\fnlloc \\simeq 5$ \\cite{KomatsuSpergel} thus allowing a many $\\sigma$ detection of non-Gaussianity if the present large central values of $\\fnlloc$ were to be confirmed. An eventual detection of a large $\\fnl$ by Planck would not  however automatically imply that the observed non-Gaussianity is primordial in origin. A number of effects can produce a {\\em spurious} NG signal that can bias the final estimate. The most relevant examples of NG contaminants are probably given by diffuse foregrounds emission, unresolved point sources contamination, NG noise. Both the analyses by Yadav and Wandelt and by the WMAP team consider all these effects and conclude that they do not significantly affect the $\\fnlloc$ measurement from WMAP data. The picture gets however more complicated when we consider future Planck data.  It has been recently shown that several effects that are not important in the analysis of WMAP data become no longer negligible with Planck. For example, Serra and Cooray \\cite{SerraCooray} studied NG contributions arising from several secondary sources, and concluded that effects such as the cross-correlations SZ-lensing and ISW-lensing produce a bias in the estimate of $\\fnlloc$ which is at the level of the expected estimator variance at Planck angular resolution. Analogous conclusions have been reached by Babich and Pierpaoli \\cite{BabichPierpaoli} for the cross correlations of density and lensing magnification of radio and SZ point sources with the ISW effect. Note that all these effects are unimportant for present analyses both because they involve higher multipoles than those reached at the WMAP angular resolution, and because they produce a bias $\\Delta \\fnlloc \\sim 1$, thus much smaller than the WMAP sensitivity to $\\fnlloc$, which is $\\Delta \\fnlloc \\sim 30$. However, the much higher angular resolution achieved by Planck and an expected predicted sensitivity on $\\fnlloc$ given by $\\Delta \\fnlloc \\sim 5$ (and $\\Delta \\fnlloc \\sim 3$ if polarisation data are included in the analysis) will make the above mentioned sources of NG contamination no longer negligible in the future. In other words, the same nice properties that make Planck more sensitive to the detection of a primordial NG signal (i.e. high angular resolution and sensitivity) make it in fact also much more sensitive to the observation of NG contaminants and require a very careful investigation of all the potential sources of bias in the estimate of $\\fnl$.  In this paper we will consider another  potential source of uncertainty in the detection of NG, namely the propagation of the uncertainties in the cosmological parameters on the measured value of $\\fnl$. This effect can be summarised as follows: the estimator usually employed to measure $\\fnl$ assumes a given underlying cosmological model obtained by fixing the cosmological parameters at their best-fit values (obtained from the two-point CMB likelihood analysis of the experiment under exam). However the cosmological parameter estimates are characterised by uncertainties that should be propagated into the $\\fnl$ estimate in order to accurately quote the final error bars. The uncertainties in the cosmological parameters can be safely neglected as long as they are much smaller than the variance of the NG estimator. While this works well for WMAP, it is a priori unclear whether it is still a good approximation for Planck.  The paper is structured as follow: in section \\ref{sec:bias} we will describe the $\\fnl$ estimator commonly employed in the analysis and study in detail the effect of cosmological parameters uncertainties on this estimator. After deriving the error propagation formulae (\\ref{eqn:bias}) and (\\ref{eqn:propagation}) we will apply them to obtain analytical and numerical estimates of the expected $\\fnl$ uncertainty for WMAP and Planck, both in the local and equilateral case. In section \\ref{sec:Fisher} we will then consider the possibility of applying a more complex analysis in which $\\fnl$ is estimated by firstly allowing the cosmological parameters to vary and then by marginalising over them, rather than by fixing the cosmological model. In this case we adopt a Fisher matrix approach to propagate the parameter uncertainties on the final predicted $\\Delta \\fnl$. We will finally discuss our results and draw our conclusion in section \\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions} \\noindent In this paper we considered the effect of propagating cosmological parameters uncertainties on the estimate of the primordial NG parameter $\\fnl$. We firstly show that, accounting for the large central value of $\\fnl$ presently measured \\cite{YadavWandelt,WMAP5}, the final correction from parameters uncertainties is of order $30 \\%$ of the quoted $\\fnlloc$ error bar for WMAP and at about the same level of the predicted $\\fnlloc$ error bars for Planck. If a large $\\fnl$ will be observed by Planck, the effect of these uncertainties will then be not big enough to change the conclusion that a large level of primordial non-Gaussianity is present in the data. However the effect is important enough to change the significance of the detection and should be taken into account when quoting the error bars. We finally show that the effect of cosmological parameters uncertainties becomes totally negligible if we do not fix the cosmological parameters in the analysis,  but we treat them as nuisance parameters and marginalise over their distribution in order to obtain the final $\\fnl$ estimate. Even if optimal, this last approach is nevertheless probably still inconvenient. A joint-likelihood evaluation would require a large amount of time and the final gain in the error bar would be significant only for large values of $\\fnl$, but those would produce a significant detection even in the sub-optimal approach."
    },
    "0808/0808.2253_arXiv.txt": {
        "abstract": "Imaging Atmospheric Cherenkov Telescopes (IACTs) have resulted in a break-through in very-high energy (VHE) gamma-ray astrophysics. While early IACT installations faced the problem of detecting any sources at all, current instruments are able to see many sources, often over more than two orders of magnitude in energy. As instruments and analysis methods have matured, the requirements for calibration and modelling of physical and instrumental effects have increased. In this article, a set of Monte Carlo simulation tools is described that attempts to include all relevant effects for IACTs in great detail but aims to achieve this in an efficient and flexible way. These tools were originally developed for the HEGRA IACT system and later adapted for the H.E.S.S. experiment. Their inherent flexibility to describe quite arbitrary IACT systems makes them also an ideal tool for evaluating the potential of future installations. It is in use for design studies of CTA and other projects. ",
        "introduction": "Any simulation of the IACT technique consists of two major components: the development of the extensive air showers and emission of Cherenkov light by the shower particles as one part and the detection of this light and recording of signals by the instrument as the other part. While earlier simulation tools used to come with a rather primitive description of the instrumental response included with the shower simulation, the simulation tools described here separates these parts into distinct programmes. This allows for enhanced flexibility in the implementation and sharing of components in the community.  The first major component is based on the CORSIKA program {\\cite{CORSIKA-physics}}. A first implementation of Cherenkov light in CORSIKA, by M.~Rozanska, F.~Arqueros, and S.~Martinez, implemented a rectangular grid of detectors -- for the non-imaging HEGRA AIROBICC instrument {\\cite{1995NIMPA.357..567M}}. The current implementation, by the author of this paper and with contributions by others, unified the originally separate functions for Cherenkov emission by electrons/positrons and by other particles, improved its efficiency and accuracy, and added a great level of flexibility. The CORSIKA code base includes different output formats now, some more generic and some dedicated to specific experiments. Only the most generic of these, the IACT option also developed by us, will be covered here. With the IACT option, a telescope or other Cherenkov detector in CORSIKA is defined by a fiducial sphere containing the reflector or detector.  \\begin{figure*}[tb] \\ifwithtwocolumns \\hc{\\resizebox{0.8\\textwidth}{!}{\\includegraphics{gamma+p+iron_n}}} \\else \\resizebox{13.8cm}{!}{\\includegraphics{gamma+p+iron_n}} \\fi \\caption{Cherenkov light production in CORSIKA for different primary particle examples in simulations for a site at 2200~m altitude. Darkness of the particle tracks shown increases with increasing emission of Cherenkov light.\\label{fig:shower-primaries}} \\end{figure*}  The second major component, termed \\tmtexttt{sim\\_tel\\-array} or -- in its current incarnation for the H.E.S.S. experiment \\citep{HESS-LOI} -- \\tmtexttt{sim\\_hess\\-array}, implements all the details of the detectors. This includes optical ray-tracing of the photons, their registration by photomultiplier tubes, switching of discriminators or comparators at the pixel level and at the telescope trigger level, as well as digitization of the resulting signals. Since \\tmtexttt{sim\\_telarray} is highly efficient and, in a typical configuration, requires only a small fraction of the CPU time needed for the shower simulation, it is usually run several times in parallel, reading directly from a common output pipe of the CORSIKA IACT option. Each instance might correspond to a different viewing direction, a different atmospheric transmission, or even a completely different mirror and camera definition. ",
        "conclusions": "The CORSIKA program with the IACT/ATMO extensions aims to provide the most detailed simulation of air showers and the production of atmospheric Cherenkov light, both for imaging and non-imaging detectors. For efficiency reasons not every possible detail is normally used, like the wavelength dependence of the index of refraction, but is available on demand. A variety of options are at hand to adapt the level of detail and performance as needed for specific applications. Parts of the photon detection probability can be applied within CORSIKA or can be provided by the second simulation stage, the detector simulation. The machine- and compiler-independent {\\tmem{eventio}} data format allows for a variety of schemes. The IACT extension, normally but not necessarily using this flexible data format, allows treatment of the largest telescope or detector arrays envisaged for coming generations of atmospheric Cherenkov experiments or observatories. It is part of the standard CORSIKA distribution for several years now.  For the second stage of the simulation process, a dedicated program for imaging telescopes with tessellated primary mirrors, termed \\tmtexttt{sim\\_telarray}, aims to provide the greatest level of detail in the simulation of all optical and electronics components involved in the measurement process. At the same time, \\tmtexttt{sim\\_telarray} is highly efficient. The {\\tmem{eventio}} data from the CORSIKA IACT option can be piped concurrently into several \\tmtexttt{sim\\_telarray} processes, each with a different configuration, while avoiding I/O bottlenecks on large computer clusters.  The simulation results from \\tmtexttt{sim\\_telarray} can be either written directly into an experiment-specific format, as initially done with the HEGRA IACT system, or into a more generic, {\\tmem{eventio}}-based format. The resulting data can be analysed directly from the latter format or converted into experiment-specific formats, e.g.\\ the H.E.S.S. data format, and then analysed with the same software as measured data.  A wide range of cross checks with measured data -- mainly by H.E.S.S. -- confirms that CORSIKA and \\tmtexttt{sim\\_telarray} can reproduce the data extremely well, for newly commissioned telescopes just based on pre-production lab measurements of mirror qualities, quantum efficiencies and so on. Examples of that include the optical PSF {\\citep{HESS-optics-2}}, the intensity, ring radius, and ring width of muon ring images {\\citep{Bolz-PhD}}, image shapes of gamma and background showers as well as the reconstructed shower PSF for point-like gamma-ray sources {\\citep{HESS-Crab,HESS-PKS2155}}, as well as the trigger rates {\\citep{HESS-trigger}}. After commissioning, H.E.S.S. simulations only had to be adapted to the slowly degrading optical throughput of the telescopes.  Future developments in \\tmtexttt{sim\\_telarray} may include non-spherical mirror tiles or secondary mirrors, and extensions to the photon detection code, as detectors with novel properties come into operation. Since the aim is to provide data like a real instrument, more capabilities of future read-out systems, like advanced post-processing in FPGAs which may include peak finding and shape analysis, are among the most recent additions.  The CORSIKA program together with the IACT/ATMO package is available from Forschungszentrum Karlsruhe. For details see \\\\ \\texttt{http://www-ik.fzk.de/corsika/}.  The \\tmtexttt{sim\\_telarray} program is available on request from the author. A release under an open source licence is foreseen."
    },
    "0808/0808.0583_arXiv.txt": {
        "abstract": "We aim at applying the ammonia method, recently proposed to explore the electron-to-proton mass ratio $\\mu = m_{\\rm e}/m_{\\rm p}$, to Galactic sources. Emission lines due to ammonia are originated in the numerous interstellar molecular clouds allowing to probe $\\mu$ at different galactocentric distances. High quality radio-astronomical observations of molecular cores in lines of NH$_3$ $(J,K) = (1,1)$, CCS $J_N = 2_1-1_0$, HC$_3$N $J = 5-4$, and N$_2$H$^+$ $J = 1-0$ are used to measure the relative radial velocity offsets, $\\Delta V$, between the NH$_3$ (1,1) and other molecular transitions. Robust statistical analysis is applied to the $n = 207$ individual measurements in the Perseus molecular cloud (PC), the Pipe Nebula (PN), and the infrared dark clouds (IRDCs) to deduce the mean value of $\\Delta V$ and its uncertainty. The measured values of $\\Delta V$ from the PC, PN, and IRDCs\\ show statistically significant positive velocity shifts between the line centers of NH$_3$ and other molecules. The most accurate estimates are obtained from carefully selected subsamples of the NH$_3$/CCS pairs observed in the PC ($n = 21$) and the PN ($n = 8$), $\\Delta V = 36\\pm7_{\\rm stat}\\pm13.5_{\\rm sys}$ \\ms\\ and $53\\pm11_{\\rm stat}\\pm13.5_{\\rm sys}$ \\ms, respectively, and from the $n = 36$ NH$_3$/N$_2$H$^+$ and $n = 27$ NH$_3$/HC$_3$N pairs observed in the IRDCs, $\\Delta V = 148\\pm32_{\\rm stat}\\pm13.6_{\\rm sys}$ \\ms\\ and $115\\pm37_{\\rm stat}\\pm31_{\\rm sys}$ \\ms, respectively. Being interpreted in terms of the electron-to-proton mass ratio variation, this gives \\dmm $\\sim (4-14)\\times10^{-8}$, which is an order of magnitude more sensitive than previous astronomical constraints on this quantity. If \\dmm\\ follows the gradient of the local gravitational potential, as suggested in some scalar field models, then the obtained result is in sharp contradiction ($\\sim 5$ orders of magnitude) with laboratory atomic clock experiments carried out at different points in the gravitational potential of the Sun. However, the measured signal might be consistent with chameleon-type scalar field models which predict a strong dependence of masses and coupling constants on the ambient matter density. The confirmation of a spatial variation of \\dmm\\ will help to distinguish between many theoretical proposals suggested to explain the nature of dark energy. New control measurements of the velocity offsets involving other molecules and a wider range of objects along with verification of rest frequencies are needed to investigate different instrument systematics and to come to a more certain conclusion. ",
        "introduction": "\\label{Sect1}  A fundamental characteristic of the universe is that decelerating expansion changed to accelerating at redshift $z \\sim 1$ \\cite{RFC98,RKS99,PAG99}. This late time acceleration is usually attributed to some negative-pressure component of the bulk of energy density, referred to as `dark energy'. The nature of dark energy remains unknown in spite of a good deal of theoretical and experimental efforts \\cite{CST06}. The most popular scenario is a dynamically evolving scalar field $\\varphi$, named `quintessence' \\cite{CDS98}, with the energy density subdominant at a matter dominated (decelerating) epoch ($1 \\la z \\la 1000$), and dominant at lower redshifts \\cite{PR03}. Scalar field models allow to alleviate the so-called `cosmic coincidence' problem~-- the fact that matter and vacuum energy contribute by comparable amounts to the energy density at the present cosmological epoch \\cite{BBM01}. The scalar field(s) coupling to ordinary matter leads unavoidably to long-range forces and variability of the physical constants violating the Equivalence Principle. Thus, any laboratory or space-based experiments as well as astronomical studies aimed at probing the variations of the physical constants, are the most important tools to test new theories against observations.  Present experimental facilities allow us to probe the variability of the fine-structure constant $\\alpha = e^2/(\\hbar c)$ and the electron-to-proton mass ratio $\\mu = m_{\\rm e}/m_{\\rm p}$, or different combinations of the proton gyromagnetic ratio $g_{\\rm p}$ with $\\alpha$ and $\\mu$  \\cite{FK07}.  The most accurate laboratory constraint on the temporal variation of $\\alpha$ of $\\dot{\\alpha}/\\alpha = (-1.6\\pm2.3)\\times10^{-17}$ yr$^{-1}$ is obtained from atomic clocks experiments \\cite{RHS08}. Being linearly extrapolated to redshift $z \\sim 2$, or $t \\sim 10^{10}$ yr, this laboratory bound leads to \\daa~= $(-1.6\\pm2.3)\\times10^{-7}$ which is well below a conservative upper limit on $|\\Delta\\alpha/\\alpha| < 10^{-5}$ stemmed from the analysis of quasar absorption-line systems \\cite{MWF08,SCP08,MLM08}. Here $\\Delta \\alpha/\\alpha = (\\alpha' - \\alpha)/\\alpha$, with $\\alpha, \\alpha'$ denoting the values of the fine-structure constant in the laboratory and the specific absorption/emission line system of a galactic or extragalactic object (the same definition is applied to \\dmm). However, laboratory experiments and quasar absorption spectra investigate very different time scales and different regions of the universe, and the connection between them is somewhat model dependent \\cite{MB04a,MB04b}.  The direct and model-free estimate of time variation of $\\mu$ was deduced from comparison of the frequency of a rovibrational transition in SF$_6$ molecule with the fundamental hyperfine transition in Cs: $\\dot{\\mu}/\\mu = (3.8\\pm5.6)\\times10^{-14}$ yr$^{-1}$ \\cite{SBC08}. Astrophysical estimates of $\\mu$ at high redshifs are controversial in the claims of nonzero and zero changes of $\\mu$: $\\Delta \\mu/\\mu = (-2.4\\pm0.6)\\times10^{-5}$ was inferred from the analysis of the H$_2$-bearing clouds at $z = 2.6$ and $3.0$ \\cite{RBH06}, but this value was not confirmed later \\cite{WR08,KWM08}. A better accuracy bound was obtained at lower redshift $z = 0.68$ from the analysis of radio-frequency transitions in NH$_3$, CO, HCO$^+$, and  HCN: $\\Delta \\mu/\\mu = (-0.6\\pm1.9)\\times10^{-6}$ \\cite{FK07b}, which favors a non-variability of $\\mu$ at the level of $\\sim3\\times10^{-6}$.  In these studies it was assumed that the rate of time variations dominates over the rate of possible spatial variations. However, this may not be the case if scalar fields trace the gravitational field inhomogeneities \\cite{B05,BM05}. The spatial variation of constants in a changing gravitational potential of the Sun were studied in laboratory measurements of the atomic clock transition $^1$S$_0$--$^3$P$_0$ in neutral $^{87}$Sr relative to the Cs standard and of the radio-frequency $E1$ transition between nearly degenerate, opposite-parity levels of atomic dysprosium (Dy). The values for the couplings of $\\alpha$ and $\\mu$ to the gravitational field were reported to be $k_\\alpha = (2.3\\pm3.1)\\times10^{-6}$, $k_\\mu = (-1.1\\pm1.7)\\times10^{-5}$ \\cite{BLC08}, and $k_\\alpha = (-8.7\\pm6.6)\\times10^{-6}$ \\cite{FCL07}, respectively.  Thus, the current laboratory and astrophysical estimates of the temporal and/or spatial changes of the coupling constants and fermion masses do not exhibit any meaningful features which could be interpreted in terms of scalar field effects. On the other hand, we do observe the late time acceleration of the universe driven most probably by ultra-light scalar fields. To reconcile these results chameleon-type models were suggested \\cite{KW04a,KW04b,GK04,BvB04,FN06,UGK06,OP08} which allow scalar fields to evolve on cosmological time scales today and to have simultaneous strong couplings of order unity to matter, as expected from string theory. The key assumption in this scenario is that the mass of the scalar field depends on the local matter density. This explains why cosmological scalar fields, such as quintessence, are not detectable in local tests of the Equivalence Principle since we live in a dense environment where the mass of the field can be sufficiently large and $\\varphi$-mediated interactions are short-ranged. On the other hand, the cosmological matter density is about $10^{30}$ times smaller and thus the mass of the scalar field can be very low, of order of $H_0$, the present Hubble constant. This allows the field to evolve cosmologically today. Laboratory experiments aimed at detections of chameleon particles are nowadays carried out at Fermilab \\cite{CWB08}.  In this paper we explore the astrophysical implications of the scalar field effect on the electron-to-proton mass ratio, $\\mu$, following suggestions on experimental tests of \\dmm\\ in our Galaxy as formulated in \\cite{OP08}. We consider the constraints on the spatial variations \\dmm\\ caused by the relative Doppler shifts of molecular transitions in cold molecular cores observed in the Milky Way disk at different galactocentric distances.  Our assumptions and basic physical properties of dense molecular cores which will be needed in the sequel are discussed in Sect.~\\ref{Sect2}. The statistical analysis of the line position measurements in the Perseus molecular complex, the Pipe Nebula, and the infrared dark clouds are presented in Sect.~\\ref{Sect3}. The obtained results are discussed in Sect.~\\ref{Sect4}, and conclusions are given in Sect.~\\ref{Sect5}.   \\begin{table}[t!] \\caption{Molecular transitions and $V_{\\rm LSR}$ uncertainties, $\\varepsilon_v$, used in calculations of \\dmm. The numbers in parentheses correspond to $1\\sigma$ errors (see Sect.~\\ref{Sect4-1} for more details). } \\begin{ruledtabular} \\label{tbl-1} \\begin{tabular}{llcc} \\noalign{\\smallskip} \\multicolumn{1}{c}{Transition} & \\multicolumn{1}{c}{$\\nu_{\\rm rest}$,} & $\\lambda_{\\rm rest},$ & \\multicolumn{1}{c}{$\\varepsilon_v$,} \\\\ & \\multicolumn{1}{c}{GHz} & cm &  \\multicolumn{1}{c}{\\ms} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\medskip} CCS  $J_N = 2_1-1_0$ & 22.344033(1)\\footnotemark[1] & 1.34 & 13.4 \\\\ NH$_3$   $(J,K) = (1,1)$ & 23.6944955(1)\\footnotemark[1] & 1.27 & 1.3  \\\\ HC$_3$N   $J = 5-4$ & 45.4903102(3)\\footnotemark[2] & 0.66 & 2.8  \\\\ N$_2$H$^+$   $J = 1-0$ & 93.173777(4)\\footnotemark[2] & 0.32 & 13.5  \\\\ \\end{tabular} \\end{ruledtabular} \\footnotetext[1]{Rest frequency from Ref.~\\onlinecite{RPF08}.} \\footnotetext[2]{Rest frequency from Ref.~\\onlinecite{SSK08}.} \\end{table} ",
        "conclusions": "\\label{Sect5}  In this paper we propose to use the ammonia inversion transitions in conjunction with low-lying rotational transitions of other molecules to probe the spatial changes of the electron-to-proton mass ratio $\\mu$ from observations of molecular clouds located in the disk of the Milky Way at different galactocentric distances.  The reported estimates of \\dmm\\ are obtained for the Perseus molecular cores \\cite{RPF08}, the Pipe Nebula \\cite{RLM08}, and the infrared dark clouds \\cite{SSK08}. We found evidence for the spatial variability of $\\mu$ at the level of $\\Delta\\mu/\\mu \\sim (4-14)\\times10^{-8}$ which is the most accurate astrophysical value to date. The statistical reliability of this result seems to be high enough but some instrument systematics could be possible. The revealed positive velocity offset between the NH$_3$ inversion transition and rotational transitions of other molecules should be confirmed by independent observations at different radio telescopes.  The result obtained can be compared with other tests of spatial variations of physical constants in the solar system which are based on atomic clock laboratory measurements. Thus, if \\dmm\\ follows the gradient of the local gravitational potential, as suggested in some scalar field models, then our estimate of \\dmm\\ contradicts significantly ($\\sim 5$ orders of magnitude) the atomic clock constraints obtained at different points in the gravitational potential of the Sun on the Earth's orbit. However, the measured signal of \\dmm $\\sim (4-14)\\times10^{-8}$ is in agreement with chameleon-type scalar field models which predict a strong dependence of masses and coupling constants on the ambient matter density.  To be completely confident that the derived velocity shift is not due to kinematic effects in the clouds but is the appearance of the spatial variation of \\dmm, new high precision radio-astronomical observations are needed for a wider range of objects. Such observations should include essentially optically thin lines of molecules co-spatially distributed with ammonia in order to reduce the Doppler noise. It is also desirable to increase the accuracy of the rest frequencies of the CCS ($2_1-1_0$) and N$_2$H$^+$ ($1-0$) transitions since their present uncertainties translate into the radial velocity error of $\\sim 13$ \\ms\\  whereas the error of the velocity shift between, for instance, CCS and NH$_3$ in the Perseus cloud can be as small as $5$ \\ms\\ (Table~\\ref{tbl-2}).  In addition, the search for spatial variations of the fine-structure constant $\\alpha$ in the Milky Way using mid- and far-infrared fine-structure transitions in atoms and ions \\cite{KPL08}, or the search for spatial variations of the combination of $\\alpha^2/\\mu$ using the [{C}\\,{\\sc ii}] $\\lambda158$ $\\mu$m line and CO rotational lines \\cite{LRK08} would be of great importance for cross-checking the results."
    },
    "0808/0808.0310_arXiv.txt": {
        "abstract": "{Rapid variability of the TeV emission in several blazars implies a central black hole mass $M_{BH}<10^8M_\\odot$, appreciably smaller than the values estimated from the $M_{BH}-L_{bulg}$ relation, and Doppler factors for the $\\gamma$-ray emitting fluid much larger than those associated with radio patterns.  We discuss the conditions in the central engine required to account for the short timescales and large luminosities observed, and propose some explanations for the inferred kinematics of the source on various scales.}  \\FullConference{Workshop on Blazar Variability across the Electromagnetic Spectrum\\\\ April 22-25 2008\\\\ Palaiseau, France}  \\begin{document} ",
        "introduction": "Recent observations of TeV AGNs raise new questions regarding the parameters of the central engine, and the location and kinematics of the TeV emission zone.   In particular, the rapid flares reported for Mrk 501 and PKS 2155-304 appear to indicate a relatively small black hole mass, which is difficult to reconcile with the black-hole/bulge relation, and accretion rates much larger than previously thought.  The systematic differences between the Doppler factor inferred from TeV observations and that associated with radio knots, and the rapid variations of the X-ray flux emitted from the HST-1 knot in M87, a region located at a distance of about $10^6 (2GM/c^2)$ from the central black hole, are also issues of interest and may be related. These observations motivate reconsideration of the standard view, according to which dissipation of the bulk energy on sub-parsec scales is accomplished predominantly through formation of internal shocks in colliding fluid shells.  Below we review these issues in greater detail and speculate about possible explanations for some of the open questions. ",
        "conclusions": ""
    },
    "0808/0808.3140_arXiv.txt": {
        "abstract": "Numerical simulation of  magnetohydrodynamic (MHD) turbulence makes it possible to study accretion dynamics in detail.  However, special effort is required to connect inflow dynamics (dependent largely on angular momentum transport) to radiation (dependent largely on thermodynamics and photon diffusion).  To this end we extend the flux-conservative, general relativistic MHD code \\texttt{HARM} from axisymmetry to full 3D.  The use of an energy conserving algorithm allows the energy dissipated in the course of relativistic accretion to be captured as heat.  The inclusion of a simple optically thin cooling function permits explicit control of the simulated disk's geometric thickness as well as a direct calculation of both the amplitude and location of the radiative cooling associated with the accretion stresses. Fully relativistic ray-tracing is used to compute the luminosity received by distant observers.  For a disk with aspect ratio $H/r \\simeq 0.1$ accreting onto a black hole with spin parameter $a/M = 0.9$, we find that there is significant dissipation beyond that predicted by the classical Novikov-Thorne model.  However, much of it occurs deep in the potential, where photon capture and gravitational redshifting can strongly limit the net photon energy escaping to infinity.  In addition, with these parameters and this radiation model, significant thermal and magnetic energy remains with the gas and is accreted by the black hole. In our model, the net luminosity reaching infinity is $6\\%$ greater than the Novikov-Thorne prediction.  If the accreted thermal energy were wholly radiated, the total luminosity of the accretion flow would be $\\simeq 20\\%$ greater than the Novikov-Thorne value. ",
        "introduction": "For the past thirty-five years, it has been the standard view in the astrophysical community that the total amount of energy per unit mass dissipated in the course of accretion onto a black hole is exactly equal to the binding energy of the innermost stable circular orbit \\citep{\\NT}; consequently, it depends only on the black hole spin parameter $a/M$.  The argument leading to this result depended on a number of assumptions: that the flow is time-steady and axisymmetric, that any heat dissipated is promptly radiated, and that the $r$-$\\phi$ component of stress goes to zero at the ISCO.  Although the first several assumptions appear to be relatively innocuous, the last was regarded as questionable almost from the beginning \\citep{T74} and has been subject to renewed questioning in more recent years \\citep{K99,G99}.  If accretion flows were fundamentally hydrodynamic, the heuristic argument for this boundary condition (that the inertia of matter in the plunging region should always be much smaller than the inertia in the stable-orbit portion of the disk) would be cogent; the point raised by all its critics is that if magnetic fields play an important role, their stress would not necessarily be diminished even in regions of small matter density.  Resolution of this point is important because continued forces at and inside the ISCO would permit continued dissipation, possibly substantially increasing the total.  Although the significance of magnetic forces was no more than a speculation when the zero-stress boundary condition was first criticized, in recent years it has been recognized that they are, in fact, essential to accretion \\citep{BH98}.  Stimulated by this recognition, the past decade has seen many numerical simulations of global disk dynamics incorporating magnetic forces under the assumption of ideal magnetodhyrodynamics (MHD) \\citep{HK01,HK02,AR01a,AR01b,AR03,MM03,\\dVHK,KHH05,GSM04}.  Initially these simulations assumed Newtonian dynamics in a pseudo-Newtonian potential; in the middle of this effort, new codes were developed that permit simulations in general relativity \\citep{\\dVH,\\GMT}.  So far, while often treating angular momentum transport quite accurately, all of these simulations have handled thermodynamics and energy transport comparatively crudely: In {\\tt GRMHD}, the code developed by De Villiers and Hawley, only an internal energy equation is solved, in which the gas is assumed to behave adiabatically except in shocks; in this code, therefore, there are sizable (and uncontrolled) numerical losses of energy whenever magnetic field or kinetic energy is lost on the gridscale. By contrast, in \\texttt{HARM}, the code developed by Gammie et~al., a total energy equation is solved (so no energy is lost), but there are also no radiative losses.  The best effort that could be made to estimate actual radiative efficiency was therefore through plausible, but {\\it ad hoc} models, usually keyed to the magnetic stress \\citep{2008arXiv0801.2974B}.  In an effort to remedy this situation, we have altered the \\texttt{HARM} code in two significant ways.  First, we have extended it from 2D (axisymmetric) to 3D.  This extension has two major consequences: we can study nonaxisymmetric fluctuations, and are free from 2D artifacts like the ``channel solution\"; and we are not limited by the anti-dynamo theorem to short duration simulations.  Second, we have introduced a toy-model optically thin cooling function.  By this means, we can track how much radiation might be produced (and where) in order to compute the radiative efficiency explicitly.  We can also use this cooling function to regulate the geometric thickness of the accretion flow.  A detailed description of the new code (which we call \\texttt{HARM3D}) can be found in \\S~2.  For our first use of this code, we chose to run a simulation that would illustrate how MHD turbulence influences the global energetics of accretion onto a black hole.  Its results can be compared directly to those of NT: Time-averages of its data can be matched against the classical model's steady state.  Quantities integrated over constant-radius shells can be compared with the corresponding vertically-integrated ones derived assuming axisymmetric and ``razor-thin\" disks.  The cooling function can be designed to (almost) reproduce the prompt radiation assumption. However, we have no need to impose any guessed boundary condition on the stress because in this numerical calculation we are able to use the the real physical boundary condition to accretion dynamics: the black hole's event horizon.  Thus, the ratio of the energy radiated in this simulation to the mass accreted in it provides a direct test of how much the zero-stress boundary condition affects the radiative efficiency.  In addition, of course, we will also be able to examine the interesting effects of non-stationary flow, non-axisymmetry, and so on.  Because we recognize that quantitative results may well depend on a number of parameters (magnetic field configuration and disk thickness, most notably) and because our radiation model does not fully represent any particular physical situation, we emphasize that the numbers we present here are only preliminary samples.  When we discuss these results, we will explain more specifically the degree to which they are model-dependent.  We intend to explore more fully in future work both how to model this process more realistically and how external parameters such as magnetic configuration and accretion rate couple with black hole spin to control the radiative output of accretion onto black holes. ",
        "conclusions": "Global disk simulations have for many years focused on dynamical effects, i.e., angular momentum transport leading to inflow.   To link them to observations, however, requires including considerations of thermodynamics, for the energy to radiate photons is drawn from the thermal energy of the gas (whether or not the particle distribution functions are in fact near those of thermal equilibrium).  By combining an energy conserving algorithm with an explicit cooling function in a new simulation code, \\texttt{HARM3D}, we are able to begin the first steps toward drawing that connection.  In this first application of our new technique, we have found that a disk with $H/r \\simeq 0.1$ accreting onto a black hole with spin parameter $a/M = 0.9$ carries thermal and magnetic energy past the ISCO at a rate $\\simeq 0.05$ per unit rest-mass, while producing radiation that reaches infinity at a rate $\\simeq 0.15$ per unit rest-mass.  These numbers contrast with those of the classical NT model, in which the flow carries no thermal or magnetic energy, and for $a/M=0.9$ radiates $\\simeq 0.14$ per unit rest-mass to infinity.  Determining the observed luminous efficiency of a more realistic accretion disk, as opposed to the ideal NT model, depends on the careful assessment of several potentially offsetting effects.  First, additional thermal, magnetic, and radiated energy can be drawn from the orbital energy by magnetic stresses that can persist through the location of the ISCO and all the way down to the horizon. However, only a fraction of that energy need be radiated, with much of the remainder retained as heat and magnetic field captured by the hole. Next, even if there is enhanced photon production near and inside the ISCO, for this particular spin, the combination of comparatively high capture probability and gravitational redshift means that little radiation from inside the ISCO reaches infinity.  For lower spin holes the ISCO is further from the horizon and the plunging region can be more effectively represented in the luminosity at infinity \\citep{2008arXiv0801.2974B}.  These results have implications for the spectral shape of the emitted radiation.  Generically, the effect of the continuing stresses is to move the radius of peak emission inward and raise the fluid-frame effective temperature at that location.   For example, in this instance the maximum in $dL/dr$ (after allowing for photon capture and all Doppler shifts) moves from the NT prediction of $r \\simeq 4.3M$ to $r \\simeq 3.5M$.  Similarly, the fluid-frame flux at the peak of $dL/dr$ is about $30\\%$ greater ($7\\%$ higher effective temperature) in the simulation data than in the NT model.  In terms of the radiation edge terminology introduced by \\cite{KH02} and \\cite{2008arXiv0801.2974B}, we find that $95\\%$ of the radiation reaching infinity is produced outside $r=2.75M$, in contrast to $3.6M$ in the NT model.  In a previous study, \\cite{2008arXiv0801.2974B} used the stress distributions observed in an ensemble of disk simulations to estimate the dissipation that might be associated with those stresses, and from this the accretion efficiency and maximum temperature in the spectrum reaching infinity.  After accounting for photon capture and Doppler-shifting effects, they found that, depending on the particular simulation examined and the topology of the initial magnetic field, the luminosity reaching infinity could be anywhere from $20\\%$--$100\\%$ greater than NT when $a/M = 0.9$.  The low end of this range was produced by an accretion flow whose initial field was entirely toroidal, the high end by an accretion flow whose initial field was a large dipolar loop, as in the present simulation. Thus, there is a sizable gap between their estimate of the radiative efficiency and ours.  Applying Beckwith et~al.'s expression to our data leads to a prediction for the dissipation rate very similar to theirs\\footnote{In fact, the accretion rate histories of the two simulations are remarkably similar, suggesting that the underlying physics imposes a long-term order despite the significant difference in computational algorithms.}.  The fact that our radiation rate is considerably less than this prediction suggests that the simple {\\it ansatz} used by Beckwith et~al. to directly compute dissipation from stress and equate dissipation with radiation is a simplification that likely overestimates the net emission.  For example, because our cooling function's radiation rate is at most comparable to the inflow rate near and inside the ISCO, not all the heat dissipated in that region can be radiated. However, even if all the heat generated in this simulation were radiated, the increase in efficiency relative to NT would be only $\\simeq 20\\%$. In addition, not all the work done by the stress necessarily goes into a form that is dissipated.  Kinetic and magnetic energies can be advected with the accretion flow into the black hole, producing no effective increase in\\ overall efficiency.  This point is closely related to the issue of advected energy discussed in \\S~\\ref{sec:extrap}.  Framed in the context of predictions for real accretion flows in Nature, these questions emphasize the importance of realistic dissipation and radiation physics for obtaining more accurate accounts of radiation associated with accretion.  In the vicinity of the ISCO, where the energy available for release is largest, one cannot say with confidence that in general the dissipation and cooling times are shorter than the inflow time.  Moreover, both processes are likely to depend on the detailed circumstances pertaining to any particular accreting black hole, so that there may not be a single efficiency number applicable to all black holes of a given spin.  In sum, we have shown that by use of a toy-model optically-thin cooling function, it is possible both to control the thickness of the accretion flow and to tally (approximately) the rate at which radiation can be produced by dissipation in the flow.  At relatively large radii, where the inflow time is long compared to the cooling time, our {\\it ansatz} of substituting gridscale dissipation for genuine microphysics and radiating the heat so generated at an arbitrarily chosen rate is capable of capturing the global energetics of accretion reasonably well. However, at smaller radii (particularly near and inside the ISCO), where the inflow time can be comparable to the cooling time, use of realistic dissipation and radiation rates can be more important.  Having demonstrated the technical feasibility of this approach, we will next employ it to explore more fully how accretion onto black holes depends on disk thickness and on black hole rotation rate.  In this context, we point out that although there is a standard notation for describing black hole rotation (the spin parameter $a/M$), there are several extant definitions of the scale-height, differing from one another by factors of order unity. We use the vertical density moment; standardization of this definition would be of benefit so that different calculations can be compared quantitatively without confusion.  Lastly, we remark that in this paper we have set aside the fact that photons are not the only form in which energy can be sent to infinity from the vicinity of black holes.  Accreting black holes are also capable of driving mass motions, often relativistic, that can carry significant power in Poynting flux.  Simulational work exploring the associated luminosity has already begun \\citep{2004ApJ...611..977M, 2006ApJ...641..103H,2008ApJ...678.1180B}.  In future work, we will use the new simulation code introduced here to relate the energetics of those outflows more closely to the accretion energy budget."
    },
    "0808/0808.1145_arXiv.txt": {
        "abstract": "The bulk viscosity due to the non-leptonic process involving hyperons in $K^-$ condensed matter is discussed here. We find  that the bulk viscosity is modified in a superconducting phase. Further, we demonstrate how the exotic bulk viscosity coefficient influences $r$-modes of neutron stars which might be sources of detectable gravitational waves. ",
        "introduction": "Neutron stars have different classes of quasi normal modes (QNMs) \\cite{Kok} depending on restoring forces acting on a perturbed fluid element. Here we are interested in Coriolis restored fluid modes known as the inertial $r$-modes. The $r$-modes become unstable due to gravitational radiation. It might be responsible for regulating the spin of newly born neutron stars as well as old, accreting neutron stars. It was shown that the leptonic and non-leptonic bulk viscosities might effectively damp the instability in different temperature regimes. The bulk viscosity coefficient due to non-leptonic weak processes in neutron star matter involving hyperons \\cite{Jon1,Jon2,Lin02,Han,Dal,Nar,DR1,DR2} , antikaon condensate \\cite{DR3,DR4} and  quark matter \\cite{Mad92,Mad00,Alf} was extensively calculated by many authors.  In this article, we discuss the hyperon bulk viscosity due to the non-leptonic process $n + p \\rightleftharpoons p + \\Lambda~$ in $K^-$ condensed matter and its role on damping the $r$-mode instability. In section 2, we describe the composition and equation of state (EoS) of neutron star matter. Results of hyperon bulk viscosity coefficients are explained in section 3. Section 4 gives the summary. ",
        "conclusions": "We have studied hyperon bulk viscosity in the presence of $K^-$ condensate. We find the inversion of the temperature dependence of hyperon bulk viscosity in our calculation. Though the hyperon bulk viscosity is suppressed in $K^-$ condensed matter than that in the hadronic phase, it still effectively damps the $r$-mode instability."
    },
    "0808/0808.1659_arXiv.txt": {
        "abstract": "s{ The construction of the southern site of the Pierre Auger Observatory is almost completed. Three independent  measurements  of the flux of the cosmic rays with energies larger than \\unit[$10^{18}$]{eV} have been  performed during the construction phase. The surface detector data collected until August 2007 have been used to establish  a flux suppression at the highest energies with a 6$\\sigma$ significance. The observations of cosmic rays by the  fluorescence detector allowed the extension of the energy spectrum to lower energies, where the efficiency of the surface detector is less than 100\\% and a change in the spectral index  is expected. } ",
        "introduction": "Cosmic rays are particles that travel through the galactic and intergalactic space, arriving on Earth in a broad energy range up to \\unit[100]{EeV}. Their flux drops steeply with energy, from  a few particles per second per  m$^2$ up to one particle per km$^2$ per century at the highest energies. The shape of the flux depends on the evolution of the sources, the mechanisms that accelerate particles up to highest energies and the energy losses during the propagation of the particles from the sources to earth. The transition from galactic to extragalactic source is expected to be in the \\unit[0.1-5]{EeV}  energy range.~\\cite{bib:bere,bib:allard}  In the highest energy range, above \\unit[60]{EeV}, a flux suppression is expected due to the  Greisen-Zatsepin-Kuzmin effect (GZK)~\\cite{bib:Greisen,bib:Zatsepin} and/or due to the maximum energy that a cosmic ray accelerator can reach. The GZK suppression is a propagation effect: protons interact with the cosmic microwave background (CMB) radiation losing about 15\\% of their energy at each encounter with these CMB photons. Heavier elements are dissociated through photo-disintegration.   The previous experiments AGASA~\\cite{bib:AGASA} and HiRes~\\cite{bib:HiRes} have given contradictory results regarding the ultrahigh energy end of the energy spectrum where only indirect observations of cosmic rays are possible by the detection of extensive air showers. The previous measurements of the comic rays flux were dominated by statistical or systematic uncertainties. In order to decrease the uncertainties the Pierre Auger Observatory was built as a hybrid detector, combining the two techniques employed by the forerunner experiments: a surface detector array  and a fluorescence detectors. Due to the hybrid technique the nearly calorimetric estimation of the energy of the primary  particle as obtained from the fluorescence technique can be transfered to the large number  of events  recorded by the surface detector.  Recently the flux suppression  has been seen both by the Pierre Auger  Observatory~\\cite{bib:ICRC07Roth} and by the  HiRes collaboration~\\cite{bib:HiRes}.    The observatory is  described in the second section. The surface detector data are used to deduce the energy spectrum above \\unit[3]{EeV} where the trigger efficiency is 100\\%,  as is described in the third section. Another complementary data set is delivered by the fluorescence detector itself. It can be used to extend the energy range down to \\unit[1]{EeV}. This measurement is described below in the last section together with the method of combining it with the surface detector flux estimate. ",
        "conclusions": "\\label{sec:conc}  The energy spectrum has been measured at the Pierre Auger Observatory with three independent data sets and the agreement is better than 4\\%.  A flux suppression at the highest energies has been established with a significance of 6 standard deviations.  Combined with the recent result that the highest energy cosmic rays are anisotropic and a good correlation only with relatively nearby sources has been found~\\cite{bib:Science}, it hints towards a GZK effect.  A change of the spectral index occurs at about  \\unit[4]{EeV} which might indicate either the transition from galactic to extragalactic origin of cosmic rays or a propagation effect. To distinguish between the models a determination of the mass of the cosmic rays~\\cite{bib:ICRC07Unger}  over the entire energy range is necessary. The exact location of the spectral index change and an accurate measurement of the spectrum below \\unit[3]{EeV} will be possible in the near future with AMIGA and HEAT extensions of the observatory.  The exact shape of the energy spectrum in the highest energy range, which can be used to distinguish between acceleration scenarios~\\cite{Yamamoto:2007xj}, will be determined with greatly improved precision over the next 10 years of data taking."
    },
    "0808/0808.1832_arXiv.txt": {
        "abstract": "% High-energy gamma-ray emission is theoretically expected to arise in tight binary star systems (with high mass loss and high velocity winds), although the evidence of this relationship has proven to be elusive so far. Here we present the first bounds on this putative emission from isolated Wolf-Rayet (WR) star binaries, WR\\,147 and WR\\,146, obtained from observations with the MAGIC telescope. ",
        "introduction": "%  WR stars represent an evolved stage of hot ($T_\\textrm{eff} > 20000$\\,K), massive ($M_\\textrm{ZAMS}>25 M_\\odot$) stars, and display some of the strongest sustained winds among galactic objects: their terminal velocities may reach the range $v_\\infty >1000-5000$\\,km/s. Perhaps with the exception of the short lived luminous blue variables (LBVs), they have the highest known mass loss rate $\\dot M \\sim 10^{-4}...10^{-5}~M_\\odot/$\\,yr of any stellar type.  Thus, colliding winds of massive star binary systems are considered as potential sites of non-thermal high-energy photon production, via leptonic and/or hadronic process after acceleration of primary particles in the collision shock \\citep[e.g.,][]{Eichler93}.  This possibility is substantiated by the detection of (non-thermal) synchrotron radio-emission from the expected colliding wind location in some binaries (see below). Many models have been proposed to predict GeV to TeV emission from these binaries, with different levels of detail \\citep[e.g., among recent works see][]{Benaglia01,Benaglia03,Pittard06, Reimer06}. Conceptually, the process would mimic the cases of LS\\,5039 \\citep{Aha05b}, PSR\\,B1259-63 \\citep{Aha06} or LS\\,I\\,+61\\,303 \\citep{Alb06,Alb08d,Alb08e}, particularly if they result in pulsar-driven $\\gamma$-ray binaries in which $\\gamma$-rays may arise from a shock region produced by the interaction of the winds of the two components.\\footnote{We recall that these binaries have orbital modulated TeV emission, and are formed by a massive star with a strong wind and a compact object, which only in the case of PSR B 1250-63 is known to be a pulsar.}  In a recent paper, the High Energy Stereoscopic array (H.E.S.S.) collaboration reported the discovery of very high energy (VHE) $\\gamma$-ray emission coincident with the young stellar cluster Westerlund\\,2 \\citep{Aha07}.  The High Energy Gamma Ray Astronomy (HEGRA) and the Major Atmospheric Gamma Imaging Cherenkov (MAGIC) telescope detected the source TeV\\,J2032+4130 and suggested a possible connection with the Cygnus OB2 cluster \\citep{Aha02,Aha05a,Alb08a}. Theoretically, the relationship between stellar associations and high-energy emission has been put forward by many authors, since early-type stellar associations have long been proposed as cosmic-ray acceleration sites and also as providers of target material for cosmic-ray interactions \\citep[see for instance the recent papers by][]{ Parizot04,Torres04,Bednarek05,Eva06} and references therein.  In the case of Cygnus\\,OB2, the VHE emission is supposed to occur at a region displaced from the center of the association, where detailed multiwavelength studies revealed an overdensity of hot OB stars, although not WRs \\citep{Butt03,Butt06}. In the case of Westerlund\\,2, the stellar cluster contains at least a dozen early-type O-stars, and two remarkable WR stars, WR\\,20a and WR\\,20b. In particular, WR\\,20a was recently established to be a binary \\citep{Rauw04,Bonanos04}. Based on the orbital period, the minimum masses were found to be around $(83 \\pm 5)$\\,$M_{\\odot}$ and $(82 \\pm 5)$\\,$M_{\\odot}$ for the binary components \\citep{Rauw05}, what certainly qualifies it among the most massive binary systems in our Galaxy.  No significant flux variability or orbital periodicity was found in the corresponding data sets, neither for the Cygnus\\,OB2 nor for the Westerlund\\,2 regions. Unless such an orbital period is detected in future datasets, it would be very difficult or impossible for the current generation of instruments to distinguish whether the radiation observed from these associations is coming from a isolated binary system or rather is generated as a collective effect of the whole cluster. Thus, a direct measurement of single WR binary systems, to explore whether they are able to produce high-energy $\\gamma$-ray emission in an isolated condition, is worth pursuing. After briefly motivating the scenario for $\\gamma$-ray emission from isolated binary systems, we present the results of MAGIC observations of two such systems. ",
        "conclusions": "%  Our search for VHE $\\gamma$-ray emission from two archetypical cases of WR binaries produced the first bounds on such systems. These bounds constrain theoretical models (or, assuming correctness of the models, their internal parameters, such as the -unknown- orbital phases of the systems). The establishment of WR binaries as VHE $\\gamma$-ray sources is yet pending.  The case for WR\\,147 as a potential $\\gamma$-ray source for MAGIC was theoretically established before, as shown in the corresponding curves of Figure~\\ref{fig:wr147} from \\citet{Reimer06}, for most of the orbital phases. The validity of this model is baselined on an assumed ensemble of orbital parameters, which are still unknown for this system. For instance, ignorance of its current phase as well as of its eccentricity and inclination makes a direct ruling out of this model impossible, although that the presented scenario could nominally survive only for phases close to 0, defined where the line of sight encounters first with the WN8 and then the B0.5V star. The MAGIC observations show that irrespective of phase, GLAST should see a flux cutoff well within its range of detectability in the tens of GeV regime, if it is able to detect the stars at all.    {\\it We acknowledge A. Reimer for providing us the theoretical curves depicted in Figure 1. We also would like to thank the Instituto de Astrofisica de Canarias for the excellent working conditions at the Observatorio del Roque de los Muchachos in La Palma. The support of the German BMBF and MPG, the Italian INFN and Spanish CICYT is gratefully acknowledged. This work was also supported by ETH Research Grant TH 34/043, by the Polish MNiSzW Grant N N203 390834, and by the YIP of the Helmholtz Gemeinschaft.}"
    },
    "0808/0808.3006_arXiv.txt": {
        "abstract": "This paper describes an application of the Mathisson-Papapetrou-Dixon (MPD) equations in analytic perturbation form to the case of circular motion around a radially accreting or radiating black hole described by the Vaidya metric. Based on the formalism presented earlier, this paper explores the effects of mass accretion or loss of the central body on the overall dynamics of the orbiting spinning particle. This includes changes to its squared mass and spin magnitude due to the classical analog of radiative corrections from spin-curvature coupling. Various quantitative consequences are explored when considering orbital motion near the black hole's event horizon. An analysis on the orbital stability properties due to spin-curvature interactions is examined briefly, with conclusions in general agreement with previous work performed for the case of circular motion around a Kerr black hole. ",
        "introduction": "\\label{sec:1}  The Mathisson-Papapetrou-Dixon (MPD) equations \\cite{Mathisson,Papapetrou,Dixon1,Dixon2} represent a well-known description of classical spinning particle motion in the presence of a curved space-time background. They comprise the ``pole-dipole approximation'' for the dynamics of extended bodies with spin angular momentum in the vicinity of black holes, neutron stars, or other sources of space-time curvature where a strong gravitational field is generated. This includes sources which themselves are time-varying for the duration of a spinning particle's motion along its worldline. As a consequence, the spin-curvature coupling term in the MPD equations, which generates an external force and torque to act on the spinning particle, also becomes time-varying, leading to potentially very interesting dynamical effects experienced by the particle.  One particularly interesting space-time background with an explicit time dependence is known as the Vaidya metric, which describes space-time curvature due to a spherically symmetric compact source that either radially accretes surrounding radiation, or radiates away its central mass. Given that most astrophysical sources have at least some orbital or spin angular momentum during their formation, it is unlikely to find candidate sources in the night sky that carry the properties exactly described by the Vaidya metric. However, because of its relative simplicity compared to the Kerr metric to describe rotating black holes, while also having a time-dependent central mass, the Vaidya metric nonetheless provides an ideal testing ground for understanding subtle properties of the MPD equations for an orbiting spinning particle that is sensitive to a time-varying gravitational field. A recent paper \\cite{Singh1} presents an extensive numerical investigation of the MPD equations in a Vaidya background, modelling a point dipole in circular orbit around a much heavier non-rotating black hole described by a monotonically increasing central mass function in terms of known functions. This paper also shows, using only the quadrupole moment formula, that the dynamical background due to a growing central mass can influence the shape and frequency of gravitational waveforms generated by the spinning particle for a sufficiently large mass accretion rate, with potentially useful implications for low-frequency gravitational wave astronomy via the space-based LISA observatory \\cite{LISA}.  Although a numerical treatment of the MPD equations in a Vaidya background is undoubtedly a useful exercise, an analytical exploration of the same problem is definitely beneficial in many respects. For example, knowing the explicit time-dependence of the mass function within an analytical expression of the MPD equations allows for the study of conditions where instabilities in the dynamical system most likely will occur. It can also potentially give useful insight for knowing when a mass increase or loss will lead to macroscopic changes in the particle's orbit for a predetermined mass accretion or loss rate. Furthermore, the general results obtained from such a study can provide clues for how a spinning particle may respond due to a more realistic time-dependent source than one described by the Vaidya metric, such as a pulsating star, particularly on determining the most dominant contribution to its response.  A recent development on the study of the MPD equations involves a linear perturbative approach first introduced by Chicone, Mashhoon, and Punsly (CMP) \\cite{Chicone}, with an application by Mashhoon and Singh \\cite{Mashhoon2} for determining a first-order perturbation of a circular orbit around a Kerr black hole due to spin-curvature coupling. This first approach was more recently generalized by Singh \\cite{Singh0} to accommodate for higher-order contributions in powers of $s/(m \\, r) \\ll 1$, where $\\rho = s/m$ is the M{\\o}ller radius \\cite{Mashhoon2,Moller} in terms of the particle's spin magnitude $s$ and mass $m$, and $r$ is the radial distance from the background mass source to the particle's location. This generalization can be applied to formally infinite order in the perturbation expansion parameter and makes no reference to any particular space-time metric or symmetries therein. A detailed application of the generalized CMP approximation to the MPD equations was just presented for the case of circular motion around a Kerr black hole. For future reference, this recent paper is now identified as ``Paper~I'' \\cite{Singh2}. It would be very interesting to perform the same investigation as found in Paper~I, but this time applied to the Vaidya metric.  The purpose of this paper is to apply the generalized CMP approximation of the MPD equations to describe circular motion around a static compact object described by the Vaidya metric, and incorporate both mass accretion from null radiation and outgoing radiation within the formalism. This paper begins with a brief review of the MPD equations and the generalized CMP approximation \\cite{Mashhoon2,Singh0,Singh2}, found in Sec.~\\ref{sec:2}. An introduction to the Vaidya metric \\cite{Singh1} and its application to the generalized CMP approximation is then presented in Sec.~\\ref{sec:3}. Following this, Sec.~\\ref{sec:4} describes the main results for the case of circular motion around the central body to second order in the perturbation expansion parameter, including the ``radiative corrections'' of the squared mass and spin magnitudes predicted within the underlying formalism \\cite{Singh0,Singh2}. Afterwards, a discussion of the obtained results is given in Sec.~\\ref{sec:5}, followed by a brief conclusion. Consistent with Paper~I, the Riemann and Ricci tensors follow the conventions of MTW \\cite{MTW} with signature $+2$, and assuming geometric units of $G = c = 1$. ",
        "conclusions": "\\label{sec:6}  This paper is an application of the generalized CMP approximation approach to the Mathisson-Papapetrou-Dixon equations of motion of a spinning point particle in orbit around a spherical black hole in the presence of radially inflowing and outflowing radiation, as described by the Vaidya metric. When compared to a similar analysis performed for orbital motion around a Kerr black hole, as described in Paper I \\cite{Singh2}, all relevant computations, including the ``radiative corrections'' to the particle's squared mass and spin magnitudes, have nontrivial properties due to the explicit time-dependence of the space-time background. It is somewhat ironic that, while the Vaidya metric is arguably much simpler in form than the Kerr metric, its time-dependence leads to much more complicated mathematical structure in the generalized CMP approximation than displayed in the previous application. As with Paper I, some numerical analysis is performed to illustrate conditions for the emergence of instabilities in the particle's orbit, with the suggestion that the M{\\o}ller radius needs to remain positive-valued in order to avoid the transition away from stable motion.  The next step in this exploration is to obtain the perturbed orbit from the results determined with the generalized CMP approximation, while also incorporating the effects of gravitational radiation within the process. As noted in Paper I, this generalization introduces conceptual and technical challenges that are still not clearly understood at present. This consideration will be deferred to a future publication once these challenges are overcome. Another possibility is to explore a many-body interaction with spin incorporated through the generalized CMP approximation. To do this requires understanding the expected tidal and spin-spin interactions to be found when dealing with such a problem, which have a separate set of conceptual and technical challenges to consider. Nonetheless, both the work presented here and in Paper I illustrate the potential that comes from this line of research."
    },
    "0808/0808.0233_arXiv.txt": {
        "abstract": "We present various properties of nuclear and compact-star matter, comparing the predictions from two kinds of phenomenological approaches: relativistic models (both with constant and density-dependent couplings) and non-relativistic Skyrme-type interactions. We mainly focus on the liquid-gas instabilities that occur at sub-saturation densities, leading to the decomposition of the homogeneous matter into a clusterized phase. Such study is related to the description of neutron-star crust (at zero temperature) and of supernova dynamics (at finite temperature). ",
        "introduction": "The knowledge of the equation of state (EOS) of nuclear matter under exotic conditions is essential for our understanding of the nuclear force and for astrophysical applications. This implies high isospin asymmetries, finite temperatures, and a wide density range (both for subsaturation and suprasaturation densities). The next generation of observational and experimental data is expected to bring new constraints in order to refine the theoretical models: for instance, the forthcoming radioactive-ion-beam facilities (such as FAIR@GSI and SPIRAL2@GANIL) will allow the investigation of the isospin degree of freedom in nuclear structure and dynamics.  The present work is dedicated to the predictions of different effective nuclear models. It is mainly focused on the liquid-gas instabilities present in nuclear and stellar matter at sub-saturation density. These instabilities are directly related to the bulk EOS. They are used to explain the multifragmentation phenomenon occurring in collisions around the Fermi energy~\\cite{multi}: in the spinodal decomposition scenario, fragment formation is induced by the fast development of spinodal instabilities in the low-density expanding matter formed just after the collision~\\cite{chomaz-PhysRep}. Finite-size liquid-gas instabilities are also important for compact-star physics: matter non-homogeneities in the (hot) core of type II supernovae is expected to affect the dynamics of the explosion, and the crust of (cold) neutron stars contains a non-homogeneous phase commonly named {\\em pasta phase}~\\cite{pethick,horo,maruyama05}. It should be noticed that the study of liquid-gas instabilities is complementary to the equilibrium approaches which are also used to describe the clusterized stellar matter: namely (at very low density) the virial equation of state ~\\cite{hor06}, and (at higher density) the calculation of the pasta phases as the ground state shaped by the competition between Coulomb repulsion and surface tension. Although nuclear equilibrium is expected to be reached in most stellar conditions, the spinodal-instability properties should help to understand the physics of  compact stars in the following ways: (i) giving an estimation of clusterized-matter properties, such as cluster size and composition; (ii) showing the minimal region where the equilibrated matter must be formed of clusters; (iii) possibly, playing a direct role in cluster formation for specific situations involving very short time scales (as may happen during a supernova explosion).   In this paper, we compare predictions from two kinds of models based on phenomenological density functionals: relativistic and non-relativistic. Both are commonly used to describe asymmetric nuclear matter, in the framework of exotic nuclei as well as compact stars. However, it is well-known that, although all give a quite good description of stable nuclei (consistently with the constraints included in the fitting procedures), they present different behaviours as soon as exotic conditions are reached, especially in the isovector channel. Our scope is to explore the impact of these different behaviors on quantities of interest for compact-star physics, such as the clusterization properties and the matter composition at $\\beta$-equilibrium.  As a non-relativistic approach, we use the effective density dependent Skyrme-type interaction~\\cite{original,brink,jirina07}. The simple form of the Skyrme functional makes it an attractive model for the description of both nuclei and compact-star matter. It was originally intended to describe nuclear properties through the mass table, and the older parametrizations only include in their fits constraints from magic-nucleus properties along the stability line. Trying to give a reliable description of exotic nuclei and stellar matter, the modern Skyrme parametrizations also include in their fitting procedures results from microscopic calculations of neutron-rich matter. The Skyrme-Lyon (SLy) forces for instance have been used in studies of neutron-star crust~\\cite{Douchin-PLB485, Douchin-NPA665}. Such parametrizations are among the 27 forces which were not ruled out for unfit neutron-star properties in the extensive study by J.R.~Stone et al.~\\cite{stone}, where 87 Skyrme parametrizations were checked.  In contrast with the Skyrme approach, the relativistic nuclear models are, by construction, causal and can thus be applied to a wider region of the compact stars (as long as matter is supposed to be in an hadronic phase). Relativistic Mean-Field models (RMF) have been used to describe the EOS of compact stars~\\cite{prak97,glen00,compact1}, both cold and warm. The Density Dependent Hadronic models (DDH) are an alternative approach for the description of nuclear matter and finite nuclei~\\cite{fuchs}. In DDH, the non-linear self-interactions of the mesons occurring in constant coupling models are substituted by density-dependent meson-nucleon coupling parameters, motivated by Dirac-Brueckner calculations of nuclear matter. Such models are found to behave more closely to the non-relativistic ones.   Relativistic and Skyrme approaches have been compared from the formal point of view in a recent work~\\cite{bcmp07}, where a low-density expansion of the RMF and DDH models have been used in order to directly compare the different density functionals. oth kinds of models have been used separately in several previous works for the study of spinodal instabilities in nuclear and compact-star matter, at zero and finite temperature: see for instance Refs.~\\cite{inst04, inst06, umodes06,umodes06a,umodes08} for relativistic models, and Refs.~\\cite{JMPC-PRC67, CD-A2, CD-A3} for Skyrme models. The same qualitative features are reproduced (general shape of the instability regions, isospin-distillation property of the phase separation). The scope of the present paper is then to have a direct look at the quantitative differences between relativistic and Skyrme-model predictions. We wish to investigate the extent to which the different temperature and  isospin dependences of the nuclear EOS can affect the neutron-star properties and determine the sensitive features that have to be constrained.     In section II we briefly review the relativistic and Skyrme models used in the present work. In section III, we present the Vlasov formalism that we use to address the dynamic instabilities, in both frameworks. Nuclear-matter properties are discussed for the different models in section IV, where we present the nuclear EOS (in isoscalar and isovector channel) as well as the spinodal instabilities (in both thermodynamical and dynamical frameworks). Properties of stellar matter, including homogeneous $\\beta$-equilibrium matter and instabilities against clusterization, are discussed in section V. In the last section we draw the main conclusions from our work. ",
        "conclusions": "The present work provides a direct comparison of Skyrme and relativistic models predictions for nuclear and compact-star matter properties, involving the bulk equation of state and the finite-size liquid-gas instabilities. In our comparison, many similarities are found, and some differences are pointed out.  As expected, the largest differences are obtained for the isovector properties. The relativistic models with constant couplings have a very hard symmetry energy. On the opposite, the older Skyrme forces predict a decrease of the symmetry energy at high density. Between these two extremes, modern Skyrme forces and DDH models present similar behaviors. However, we can notice the specificities of DDH$\\delta$ and SLy230a in the high density region: both present a stiffer $a_{\\s}$ evolution, and atypic rates $\\mu_3/\\mu_3^{\\rm{para}}$ (for DDH$\\delta$, this is linked to the inclusion of the $\\delta$ meson).  Concerning the thermodynamic spinodal region, we have verified that its isoscalar-density extension $\\rho_{\\s}$ is correlated to the saturation density $\\rho_0$. At zero temperature, the spinodal contour reaches very high asymmetry for all models; using the $(\\rho,\\mu_3)$ representation, we have obtained different $\\mu_3$ extensions, reflecting the low density values of the symmetry energy. Relativistic models tend to yield smaller spinodal regions, both in $\\rho$ and $\\mu_3$ directions. The thermodynamic instability direction leads to the usual isospin distillation for all models; however  for larger densities this effect becomes stronger for RMF models with constant couplings and presents a reduction for Skyrme and DDH models. Some very recent DBH results seem to confirm this trend, although with a smaller reduction~\\cite{isaac08}.  The dynamic finite-size instabilities leading to matter clusterization have been addressed in the Vlasov formalism. The wavelength associated with the most unstable mode gives an estimation of the cluster size issuing from a spinodal decomposition: it was shown that Skyrme parametrizations show larger growth rates and favor fluctuations of shorter wavelengths. The favored wavelengths reflect the finite range behaviour of the force. In the Skyrme parametrizations, only $q^2$ terms are included. With relativistic models, the finite range is described by the exchange of mesons. For all models, the dynamic distillation effect is reduced with respect to the bulk one; this reduction appears stronger for relativistic models including the $\\delta$ meson.  For the study of compact-star matter properties, we have considered $\\beta$-equilibrium under two different conditions: for neutrino-free matter, and for matter with trapped neutrinos according to a fixed lepton fraction. For neutrino-free matter, the proton fraction at $\\beta$-equilibrium is very sensitive to the symmetry energy $a_{\\s}(\\rho)$. The RMF models with constant couplings thus predict very high proton fractions, allowing the direct URCA process already at quite low densities. However, it was observed that the proton-fraction is not uniquely determined by $a_{\\s}(\\rho)$, but also depends on the parabolic behavior of the isovector EOS: see SLy230a and DDH$\\delta$, leading to  proton fractions lower than  other models with similar symmetry energy. The situation is different for matter with trapped neutrinos: fixing a lepton fraction $Y_l=0.4$, all models (except SIII) predict a nearly constant proton fraction $\\sim 0.3-0.35$.  We have finally discussed clusterization of stellar matter in two different contexts: at zero temperature, where it is related to the extension of the neutron-star crust, and at finite temperature, where it should influence supernova dynamics. In the first case, we give a lower estimation ($\\rho_{\\rm{cross,in}}$) of the transition density at the inner edge of the crust. Our results are in reasonable agreement with values obtained within pasta-phase calculations. Modern Skyrme parametrizations and DDH models give similar results, $\\rho_{\\rm{cross,in}}\\sim 0.75$~fm$^{-3}$.  Stellar matter at finite temperature is addressed only with Skyrme models and NL3. Finite-$T$ calculations still have to be performed for the relativistic models with density-dependent couplings, but we do not expect those results will affect our present conclusions. For all the models we show, neutrino trapping is needed to reach the instability region at $T=10$~MeV. The required trapping rate is higher for models with low symmetry energy at subsaturation density, such as NL3.   Globally, we can say that the isovector EOS shows quite large quantitative differences even between modern forces: new data are still needed to better constrain the neutron-rich matter properties. We have shown the consequences of the stiff symmetry energy of the RMF models:stronger distillation, smaller $\\mu_3$-extension of the spinodal, and larger proton-fractions at high density, allowing the URCA process. The finite-range of the force is also important: its different behavior between Skyrme and relativistic models causes the difference in the predictions of typical cluster size.  A similar work comparing phenomenological models with results from other approaches such as Brueckner-Hartree-Fock should be performed."
    },
    "0808/0808.2693_arXiv.txt": {
        "abstract": "We present $10\\arcmin \\times 50\\arcmin$ scan maps around an M supergiant $\\alpha$ Ori in at 65, 90, 140 and 160$\\mu$m obtained with the {\\sl AKARI Infrared Astronomy Satellite}. Higher spatial resolution data with the exact analytic solution permit us to fit the de-projected shape of the stellar wind bow shock around $\\alpha$ Ori to have the stand-off distance of $4\\farcm8$, position angle of $55^{\\circ}$ and inclination angle of $56^{\\circ}$. The shape of the bow shock suggests that the peculiar velocity of $\\alpha$ Ori with respect to the local medium is $v_* = 40~n_{\\rm H}^{-1/2}$, where $n_{\\rm H}$ is the hydrogen nucleus density at $\\alpha$ Ori. We find that the local medium is of $n_{\\rm H} = 1.5$ to 1.9 cm$^{-3}$ and the velocity of the local flow is at 11 km s$^{-1}$ by using the most recent astrometric solutions for $\\alpha$ Ori under the assumption that the local medium is moving away from the Orion OB 1 association. {\\sl AKARI} images may also reveal a vortex ring due to instabilities on the surface of the bow shock as demonstrated by numerical models. This research exemplifies the potential of {\\sl AKARI} All-Sky data as well as follow-up observations with {\\sl Herschel Space Telescope} and {\\sl Stratospheric Observatory for Infrared Astronomy} for this avenue of research in revealing the nature of interaction between the stellar wind and interstellar medium. ",
        "introduction": "Alpha Orionis (Betelgeuse, HD 39801; hereafter $\\alpha$ Ori) is a supergiant of spectral type M2 Iab:\\ and is a very bright far-infrared (far-IR) source. Some extended structure in the far-IR was recognized around $\\alpha$ Ori first by \\citet{stencel88} in the All-Sky Survey data of the {\\sl Infrared Astronomical Satellite} ({\\sl IRAS}\\/). Using an improved image reconstruction technique, \\citet{nc97} produced images of better quality for $\\alpha$ Ori and identified the extended structure around the star as an arc-shaped structure at the interface between the stellar wind and the local interstellar medium (ISM).  Such structures can be formed around mass-losing stars at which the ram pressure of the ambient ISM balances with that of the stellar wind. These stellar wind bow shock arcs had already been found around hotter, more luminous OB stars in the {\\sl IRAS} All-Sky maps \\citep{vbm88}. Most recently, \\citet{ueta06} discovered a stellar wind bow shock arc around R Hya, an asymptotic giant branch star (an evolved star of lower mass and lower rate of mass loss), using the {\\sl Spitzer Space Telescope}.  Far-IR emission from these arcs is probably due mainly to thermal emission of cold dust components of the stellar wind bow shock whose temperature peaks at far-IR. There may be contribution from low-excitation atomic lines such as [O\\emissiontype{I}] 63 $\\mu$m, [O\\emissiontype{I}] 145 $\\mu$m and [C\\emissiontype{II}] 158 $\\mu$m at these wavelengths. However, the exact emission mechanism of these far-IR bow shocks remains unclear, while an attempt to identify their spectroscopic nature is currently on-going with the {\\sl Spitzer Space Telescope}. Meanwhile, theoretical studies have been done both analytically \\citep{wilkin96, wilkin00} and numerically (e.g.\\ \\cite{ml91,dgani96,bk98,wareing07}) and the structure of the stellar wind bow shocks has been understood reasonably well.  The {\\sl AKARI Infrared Astronomy Satellite} ({\\sl AKARI}\\/; \\cite{murakami07}) makes the All-Sky Survey in the far-IR for the first time in 25 years since {\\sl IRAS} at much finer spatial resolution.\\footnote{% The effective beam size of {\\sl AKARI} ($0\\farcm5$ to $0\\farcm9$) is much smaller than that of {\\sl IRAS} ($2\\farcm0$ to $5\\farcm0$).} The original {\\sl IRAS} images of $\\alpha$ Ori allowed only rough identification of the shape of the extended far-IR emission \\citep{stencel88}. The enhanced {\\sl IRAS} images of $\\alpha$ Ori yielded just the mean radius and width of the arc with a rough estimate for the position angle of the apex \\citep{nc97}. In this paper, we characterize the structure of the far-IR stellar wind bow shock around $\\alpha$ Ori, using higher spatial-resolution {\\sl AKARI} images, and investigate the kinematics of the star, bow shock and ISM in the vicinity of the star by adopting the analytic solution of the bow shock \\citep{wilkin96} and the most recent astrometric solutions for the star \\citep{harper08}. ",
        "conclusions": "AKARI/FIS scan maps around $\\alpha$ Ori ($10\\arcmin \\times 50\\arcmin$) at 65, 90, 140 and $160\\mu$m are presented. These images show the extended emission core and most of the circumstellar arc plus the northern bar structure at much higher spatial scale than in the previously obtained {\\sl IRAS} maps. Spatial resolution of the data is good enough to define the structure of the arc to be fit with the exact analytic solution for stellar wind bow shocks, while the scan did not cover the arc in its entirety and there are some anomalies by pixel cross-talk in the SW band and by ghosting in the $160\\mu$m band due to the bright central star.  Rather circular appearance of the arc (eccentricity 0.02) suggests that the stellar wind bow shock cone of $4\\farcm8$ de-projected stand-off distance is inclined at $56^{\\circ}$ with respect to the plane of the sky and oriented at $55^{\\circ}$ position angle (east of north). Adopting the distance of 197 pc, rate of mass loss at $3.1 \\times 10^{-6}$ M$_{\\odot}$ yr$^{-1}$, and wind velocity of $17$km s$^{-1}$, the peculiar velocity of the star with respect to the ISM at $\\alpha$ Ori is found to be $v_{*} = 40~n_{\\rm H}^{-1/2}$km s$^{-1}$.  By comparing the Galactic space-velocity components of the peculiar velocity of the star (based on the orientation of the bow shock cone) and that of the apparent motion of the star (based on the astrometric solutions), we derived the space-velocity components of the ISM flow around the star. Assuming that the ISM flow in the vicinity of $\\alpha$ Ori emanates from the Orion OB 1 association, we find the particle number density per hydrogen nucleus ($n_{\\rm H}$) in the ISM around $\\alpha$ Ori to be 1.5 to 1.9 cm$^{-3}$, which translates to a 11 km s$^{-1}$ flow.  Owing to higher spatial resolution of the data, we may be witnessing the development of a vortex ring along the surface of the bow due to instabilities. This research demonstrates that far-IR images of stellar wind bow shocks are excellent diagnostic tools to investigate the kinematics of the bow and the characteristics of the ISM in the vicinity of the star. These stellar wind bow shocks have been found not only around M supergiants but also around OB stars \\citep{vbm88} and an AGB star \\citep{ueta06}. Therefore, this avenue of research is going to flourish with the coming of the AKARI All-Sky Survey \\citep{murakami07} followed by new opportunities with {\\sl Herschel Space Telescope} and {\\sl Stratospheric Observatory for Infrared Astronomy}, because there will be a wealth of new far-IR data on the stellar wind bow shocks around variety of mass-losing stars.         \\bigskip This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France. Ueta is grateful to Drs.\\ R.\\ E.\\ Stencel and G.\\ M.\\ Harper for illuminating discussions. Ueta also acknowledges contribution by a summer student, A.\\ Karska, in developing some of the custom data reduction algorithms."
    },
    "0808/0808.0143_arXiv.txt": {
        "abstract": "Based on the first principles (i.e. (i) by balancing the magnetic field advection with the term containing electron pressure tensor non-gyrotropic components in the generalised Ohm's law; (ii) using the conservation of mass; and (iii) assuming that the weak magnetic field region width, where electron meandering motion supports electron pressure tensor off-diagonal (non-gyrotropic) components, is of the order of electron Larmor radius;) a simple model of magnetic reconnection in a collisionless regime is formulated. The model is general, resembling its collisional Sweet-Parker analogue in that it is not specific to any initial configuration e.g. Harris type tearing unstable current sheet, X-point collapse or otherwise. In addition to its importance from the fundamental point of view, the collisionless reconnection model offers  a much faster  reconnection rate ($M_{c'less}=( c / \\omega_{pe})^2 / (r_{L,e} L)$) than Sweet-Parker's classical one ($M_{sp}=S^{-1/2}$). The width of the diffusion region (current sheet) in the collisionless regime is found to be $\\delta_{c'less}=( c / \\omega_{pe})^2 /r_{L,e}$, which is independent of global reconnection scale $L$ and is only prescribed by micro-physics (electron inertial length, $c / \\omega_{pe}$, and electron  Larmor radius, $r_{L,e}$). Amongst other issues, the fastness of the reconnection rate alleviates e.g. the problem of interpretation of solar flares by means of reconnection, as for the typical solar coronal parameters the obtained collisionless reconnection time can be a few minutes, as opposed to Sweet-Parker's equivalent value of $<$ a day. The new theoretical reconnection rate is compared to the MRX (Magnetic Reconnection Experiment) device experimental data by \\citet{yamada,ji} and a good agreement is obtained. ",
        "introduction": "Magnetic reconnection is the process that enables to convert magnetic energy into heat and kinetic energy of accelerated particles \\cite{pf00,biskamp,bp07}. In many astrophysical and laboratory plasma situations plasma beta (the ratio of thermal and magnetic pressures) is much smaller than unity. Thus, magnetic reconnection has attracted a considerable attention as the plasma heating and charged particle acceleration mechanism. It is believed that magnetic reconnection is responsible for solar \\cite{1997Natur.386..811I,2008AdSpR..42..895W} and stellar flares \\cite{2008A&A...481..799S,2008ApJ...676L..69C}. It has been also observed in the Earth geomagnetic tail \\cite{2008JGRA..11307S27R}.  This work has several motivations:  (i) The longstanding problems with resistive MHD (Magnetohydrodynamic), i.e. {\\it collisional} description of the magnetic reconnection has been its too slow rate if a consistent, the-first-principles approach is used  \\cite{sweet,parker}; or a fast rate \\cite{petschek}, but a lack of a proper motivation (\\cite{biskamp}, p. 79) (one can still get the fast reconnection if the resistivity is enhanced in the diffusion region, as one expects in physical applications). In the {\\it collisionless} regime there is no known simple, analytical reconnection rate (e.g. analogous to Sweet-Parker, or Petschek rates). What does exist, however, is a bulk of mostly numerical simulation work \\cite{hesse99,birn01,pritchett01,th07,th08}. It should be mentioned that a significant progress (including analytical) has been made in study of the details of collisionless reconnection such as the diffusion region structure, outflows, reconnection rates \\cite{kuz98,hesse99,birn07,klimas08}. However these were either specific to a narrow class of collisionless reconnection models, namely tearing unstable Harris current sheet; or too slow reconnection rates were obtained \\cite{vp95}. Some numerical simulation models \\cite{kuz98,hesse99,birn01,pritchett01,birn07,th07,th08,klimas08} can produce as fast reconnection rates as $M>0.5$ where $M$ is the inflow Alfv\\'enic Mach number. The lack of a simple, analytical model of collisionless reconnection ({\\it that is not specific to any initial configuration}) has been somewhat hampering the progress in understanding.  (ii) Quite often, particularly in space plasmas, there is a good reason to go beyond resistive, single-fluid MHD. A simple line of reasoning can be outlined on an example of e.g. solar corona. Fixing coronal temperature at $1.0 \\times 10^6$ K, Coulomb logarithm at 18.1, the Lundquist number (using Spitzer resistivity) is $3.7 \\times 10^{11}$. Here $L=10^5$ m was used. One of the arguments for going beyond resistive MHD is comparing typical width of a Sweet-Parker current sheet $\\delta_{sp}=S^{-1/2} L=0.16$ m to the ion inertial length. Typical scale associated with the Hall term in the generalised Ohm's law at which deviation from electron-ion coupled dynamics is observed is, $c/\\omega_{pi} =  7.2$ m. Here particle density of $n=1.0 \\times 10^{15}$ m$^{-3}$ is used. Hence, the fact that $ c/(\\omega_{pi} \\delta_{sp})=44 \\gg 1 $ warrants going {\\it beyond single fluid resistive MHD approximation} (a similar conclusion is reached by \\citet{yamada} in their Figure 12 based on laboratory plasma experiment known as MRX (Magnetic Reconnection Experiment)).   (iii) Previous results on collisionless reconnection both in tearing unstable Harris current sheet \\cite{kuz98,hesse99,birn01,pritchett01} and stressed X-point collapse \\cite{th07,th08} has shown that magnetic field is frozen into electron fluid and the term in the generalised Ohm's law that is responsible for breaking the frozen-in condition is electron pressure tensor off-diagonal (non-gyrotropic) components. Thus, a need for inclusion of the {\\it electron pressure tensor non-gyrotropic components} in a model of collisionless reconnection has become clear.  (iv) There is a growing amount of work \\cite{bp07,th07,th08} that suggests that in the collisionless regime, on the scales less than $c/\\omega_{pi}$ magnetic field is frozen into the electron fluid rather than bulk of plasma. One can write in general $\\vec V_e=\\vec V_i- \\vec j / (e n)$. This relation clearly shows that in collisional regime (when the number density $n$ is large), the difference between electron and ions speeds diminishes $V_e=V_i=V$. However, as one enters collisionless regime (when the number density $n$ is small) the deviation between electron and ion speeds starts to show. In \\citet{th08} we  proposed a possible explanation why the reconnection is fast when the Hall term is included. Inclusion of the latter means that in the reconnection inflow magnetic field is frozen into {\\it electron} fluid. As it was previously shown in \\citet{th07} (see their Figs.(7) and (11)) speed of electrons, during the reconnection peak time, is at least 4-5 times greater than that of ions. This means that electrons can bring in / take out the magnetic field attached to them into / away from the diffusion region much faster than in the case of single fluid MHD which does not distinguish between electron-ion dynamics. Thus, a need for inclusion {\\it magnetic field transport by electrons} in a model of collisionless reconnection has become clear.  (v) A pioneering work of \\citet{yamada} has demonstrated a clear transition from collisional to collisionless reconnection regimes by varying the plasma density, and established that in the collisionless regime the reconnection rate is much greater than the Sweet-Parker rate. Despite a well motivated explanation of their experimental results, one was surprised to see only experimental data points and numerical simulation results on their Figure 12. Thus, we set out with a task of formulating a simple model of magnetic reconnection in a collisionless regime. In what follows we shall be using well-motivated arguments of Sweet-Parker model, but shall be applying it to the collisionless case. ",
        "conclusions": "A new model of collisionless reconnection is formulated that is based on simple conservation laws. The obtained collisionless reconnection rate, $M_{c'less}=( c / \\omega_{pe})^2 / (r_{L,e} L)$, naturally does not depend on Lundquist number, $S$, (because it is collisionless) and is much faster than the Sweet-Parker rate, but yet somewhat slower than the Petschek rate (at least for solar coronal plasma parameters). In particular, for the same set of parameters e.g. for solar coronal plasma the rates are $M_{petschek}=0.04$, $M_{ c'less}=1.3 \\times 10^{-4}$, $M_{sp}= 1.6 \\times 10^{-6}$. The main implication for solar flares is that if this collisionless rate is used, flare time can be as short as a few minutes, that is commensurate with the observations. Note that the formulated collisionless reconnection model is general and is not specific to any initial configuration e.g. Harris type tearing unstable current sheet, X-point collapse or otherwise. In differ to previous results it relates the reconnection rate to simple, generic spatial scales such as electron inertial length, Larmor radius and global reconnection length. Therefore it is easily applicable to different space or laboratory plasma situations."
    },
    "0808/0808.1881_arXiv.txt": {
        "abstract": "B-Pol is a medium-class space mission aimed at detecting the primordial gravitational waves generated during inflation through high accuracy measurements of the Cosmic Microwave Background (CMB) polarization. We discuss the scientific background, feasibility of the experiment, and implementation developed in response to the ESA Cosmic Vision 2015-2025 Call for Proposals. ",
        "introduction": "The quest to understand the origin of the tiny fluctuations about a perfectly homogeneous and isotropic universe lies at the heart of both modern cosmology and high-energy physics. Inflationary theory offers the most satisfying and plausible explanation for the initial conditions of the universe. Inflation is a phase of superluminal expansion of space itself, within $10^{-35}s$ of the Big Bang, during which quantum fluctuations are stretched to cosmological scales (see e.g., % \\cite{star82,mukh81,guth82,hawk82,ruba82,lind83,bst84,abbo84,kolb90}). Results from Cosmic Microwave Background (CMB) experiments have established that the universe is almost spatially flat, with a nearly Gaussian, scale-invariant spectrum of primordial adiabatic perturbations (see e.g., \\cite{debe00,stom01,sper03,sper06}). These features are all consistent with the simplest models of inflation. By providing accurate measurements of the E-mode (gradient component) polarization of the CMB, the ESA mission Planck will offer more stringent tests of the inflationary paradigm \\cite{Plan05}. Nevertheless, even with such an accurate characterization of the scalar perturbations, a decisive confirmation of inflation will be lacking and large uncertainties in the allowed inflationary potentials will persist. Inflation predicts the existence of primordial gravitational waves on cosmological scales. Their detection would firmly establish the existence of a period of inflationary expansion in the early universe, and confirm the quantum origin of cosmological fluctuations that led to the large scale structure observed today. The search for primordial B-mode (curl component) polarization of the CMB provides the only opportunity to detect in the foreseeable future the imprint of these gravitational waves. Measuring the amplitude of these tensor perturbations at one length scale would fix the energy scale of inflation and its potential. Measuring their amplitude at more than one length scale would provide a powerful consistency check for a broad class of inflationary models.  If as suggested by recent CMB and large scale galaxy surveys, the power spectrum of primordial perturbations is not exactly scale invariant, then in a wide class of inflationary models the level of gravitational waves will be within the range accessible to a properly designed mission, as shown in Fig.~\\ref{Fig:UnderlyingAnisotropies}. The bulk of the statistical weight for detecting inflationary B modes is concentrated at two angular scales on the sky (see Fig.~\\ref{Fig:UnderlyingAnisotropies}): firstly, at the reionization bump at multipoles $\\ell = 2-10$ and secondly at the multipole region from about 20 to 100 (corresponding to angles larger than $\\approx 1^\\circ $ on the sky). Given that most of the signal lies on large angular scales, a full-sky survey with exquisite stability and control of systematic errors of both instrumental and astrophysical origin is required, hence the need to go to space. The B mode polarization is a clean probe of gravitational waves, since primordial scalar perturbations do not contribute to B-modes, and the effects due to intervening gravitational lenses are calculable and of order $5 \\mu K \\cdot {\\rm arcmin}$. This sets the sensitivity target for the mission. In fact, as shown by the blue dotted line in Fig.~\\ref{fig:science}, the lensing contribution to the B-mode is below the primordial B-mode for tensor to scalar perturbation ratios, $r=T/S$, above few~$\\times 10^{-2}$ at least at multipoles $\\ell \\lsim 100$, while, for example, a lensing subtraction at $\\sim 10$\\% accuracy level (in terms of angular power spectrum) is adequate to identify the primordial B-mode for $T/S$ above few~$\\times 10^{-3}$.  \\def\\gtorder{\\mathrel{\\raise.3ex\\hbox{$>$}\\mkern-14mu \\lower0.6ex\\hbox{$\\sim$}}} \\def\\ltorder{\\mathrel{\\raise.3ex\\hbox{$<$}\\mkern-14mu \\lower0.6ex\\hbox{$\\sim$}}} \\newcommand{\\muK}{\\mu  {\\rm K}} \\newcommand{\\muKarcmin}{\\mu  {\\rm K\\cdot  arcmin}} \\newcommand{\\kmbysbyMpc}{\\,{\\rm km~s^{-1}~Mpc^{-1}}} \\begin{figure}[ht] \\begin{center} \\vskip 2.cm { \\setlength{\\unitlength}{0.6cm} \\begin{picture}(14, 6)(0,0) \\put(11.6,7.2){TT,scalar} \\put(11.6,6.3){TE,scalar} \\put(11.6,5.8){EE,scalar} \\put(11.6,4.3){BB$\\leftarrow $ EE, scalar lensed} \\put(11.6,3.9){{\\tiny TT, tensor $(T/S)=10^{-1}$ }} \\put(11.6,3.0){{\\tiny TE, tensor $(T/S)=10^{-1}$ }} \\put(11.6,2.75){{\\tiny EE, tensor $(T/S)=10^{-1}$ }} \\put(11.6,2.4){BB, $(T/S)=10^{-1}$} \\put(11.6,1.7){BB, $(T/S)=10^{-2}$} \\put(11.6,1){BB, $(T/S)=10^{-3}$} \\put(5.,0.2){Multipole number ($\\ell $)} \\put(0.6,3){\\begin{rotate}{90} $\\ell (\\ell +1)C_{\\ell }/2\\pi ~[\\mu K^2]$ \\end{rotate}} \\includegraphics[width=7.2cm]{underlying_anisotropies_28juin.ps} \\end{picture} } \\end{center} \\vskip -0.5cm \\caption{\\small \\baselineskip=8pt {{\\bf Inflationary Prediction for the CMB Temperature and Polarization Anisotropies, for the Scalar and Tensor Modes.} The horizontal axis indicates the multipole number $\\ell$ and the vertical axis indicates $\\ell (\\ell +1)C_\\ell ^{AB}/(2\\pi )$ in units of $(\\mu  K)^2$, which is roughly equivalent to the power spectrum per unit of $\\ln\\ell$. The green curves indicate the TT, TE, and EE power spectra (from top to bottom) generated by the {\\it scalar} mode assuming the parameters from the best-fit model from WMAP three-year data. The BB scalar component (indicated by the heavy red curve) results from the gravitational lensing of the EE polarized CMB anisotropy by structures situated mainly around redshift $z\\approx 2$. The top four blue curves (from top to bottom on the left) indicate the TT, TE, BB, and EE spectra (BB is the heavy solid curve) resulting from the {\\it tensor} mode, assuming a scale-invariant ($n_T=0$) primordial spectrum and a tensor-to-scalar ratio $(T/S)$ of $0.1$. This value corresponds roughly to the upper limit established by WMAP. The bottom two blue curves indicate the tensor BB spectrum for $(T/S)$ equal to $0.01$ and $0.001$, respectively. For the TE cross-correlations we have plotted the log of the absolute value.} } \\label{Fig:UnderlyingAnisotropies} % \\end{figure} \\begin{figure}[th] \\includegraphics[width=9.cm,height=10.5cm,angle=90]{planck_bpol_fore_100_endOK_paperrevversion_thick.ps} \\caption{\\small \\baselineskip=8pt {\\bf B Mode Signal and Foregrounds.} CMB E and B polarization modes compatible with WMAP 3-yr data are compared to Galactic and extragalactic polarized (B-mode) foregrounds. The anticipated foreground subtraction residuals are compared to the B-mode induced by lensing (blue dots) and to the B-Pol sensitivity. The plots include cosmic and sampling variance plus instrumental noise (green dots labeled with cv+sv+n; black thick dots, noise only) assuming a multipole binning of 30\\%. The B-Pol frequency channel at 100~GHz is considered here. The corresponding Planck-HFI instrumental noise sensitivity is also displayed for comparison (four surveys, upper black thick dots). The E mode is denoted by the black long dashes. The B mode (black solid lines) is shown for $T/S = 0.1, 0.03, 0.01, 0.003, 0.001$ from top to bottom, at increasing thickness. Note that the cosmic and sampling (74\\% sky coverage) variance implies a dependence of the overall sensitivity at low multipoles on $T/S$ (again the green dots refer to $T/S =$ 0.1, 0.03, 0.01, 0.003, 0.001 from top to bottom), which is relevant for parameter estimation; instrumental noise only determines B-Pol's capability to detect the B mode. Galactic synchrotron (purple dashes) and dust (purple dot-dashes) polarized emissions produce the overall Galactic foreground (purple three dot-dashes) that is dominated by dust at 100 GHz. WMAP 3-yr power-law fits for uncorrelated dust and synchrotron have been used. For comparison, WMAP 3-yr results derived directly from the foreground maps are shown on a suitable multipole range: power-law fits provide (generous) upper limits for the power at low multipoles. Residual contaminations by Galactic foregrounds (purple three dot-dashes) are shown for 10\\% to 1\\% of the map level, at increasing thickness, as labeled on the right. The residual contribution by unsubtracted extragalactic sources, $C_\\ell^{\\rm resPS}$, and the corresponding uncertainty, $\\delta C_\\ell^{\\rm resPS}$, computed assuming a relative uncertainty $\\delta \\Pi / \\Pi = \\delta S_{\\rm lim}/S_{\\rm lim} = 10$\\% in the knowledge of their degree of polarization and in the determination of the source detection threshold, are also plotted as green dashes, thin and thick, respectively. } \\label{fig:science} \\end{figure}  A confirmation of inflation and determination of the inflationary potential would have profound implications for fundamental physics by providing new experimental data on the physics near the Planck scale. The constraints established would be indispensable for model building in string and M theory. The energy scales probed by polarization measurements lie many orders of magnitude beyond any conceivable accelerator experiment. Consequently, the quest for primordial gravitational waves from inflation constitutes a unique window for constraining the new physics near the Planck scale, which will help understand how quantum gravity unifies with the other fundamental interactions.  The implementation of this mission requires significant advances in three main areas: (1) a sensitivity to tensor modes of a factor of about 100 with respect to Planck, (2) control of systematic effects at the level of a few nK, and (3) a precise ($\\sim 1$\\%) knowledge of the galactic foreground polarization.  Aspects (1) and (2), related to the B-Pol design, will be extensively discussed in the following sections. Concerning (3), Fig.~\\ref{fig:science} compares the CMB B-mode (for various $T/S$ values) to the B-mode expected from the most relevant polarized foregrounds and their potential residuals assuming different levels of accuracy for their subtraction assuming B-Pol sensitivities. As is evident, a removal of the foreground, and in particular of the Galactic emission, at the level of about 1\\% is necessary to detect the primordial CMB B mode for $T/S \\sim 10^{-3}$. Foreground subtraction to this level can already be achieved exploiting already (or soon) available all-sky (or large area sky coverage) surveys. Important new data sets will become available in the next years regarding polarized foregrounds in the microwave (e.g., further WMAP surveys, QUAD, BICEPT, EBEX, BOOMERanG, QUIET, C$\\ell$OVER, SPIDER, Planck, etc.), radio (e.g. S-PASS, PGMS, C-BASS, GEM) and far-IR (e.g. PILOT) bands. In parallel, it will be crucial to generalize to polarization the component separation methods successfully applied to temperature data (e.g. Wiener filtering, maximum-entropy, Spectral-Matching ICA, CCA, phase methods) as well as to refine the already existing component separation methods for polarized data (e.g. template fitting methods, ICA, FastICA, PolEMICA), and to develop new techniques.  Obviously, the broad frequency coverage proposed for B-Pol is crucial to allow the application of these methods at the required level of sensitivity.  Finally we remark that because of its high sensitivity and accuracy in polarization, a mission devoted to B-modes would make substantial contributions in several other areas of astrophysics, such as the physical modeling of Galactic magnetic fields, interstellar dust and gas properties including turbulence effects \\cite{cho_lazarian}, and of cosmology, such as gravitational lensing of the CMB, cosmological reionization, and magnetic fields in the early universe (see e.g. respectively \\cite{lewis_challinor,lesgourgues_etal}, \\cite{burigana_etal}, \\cite{scannapieco_ferreira}, and references therein). ",
        "conclusions": "\\label{conclusions} The search and discovery of primordial B modes provide a unique window for exploring new physics near the Planck scale. There are no competing experimental probes able to access this domain. We have presented an implementation of a CMB polarization mission called B-Pol that fits within the tight requirements for medium-size missions within the ESA Cosmic-Vision 2015-2025 program. The instrument would survey the sky from the Lagrangian point L2 of the Sun-Earth system for two years, and produce maps of the {\\it polarized} microwave background anisotropy with sensitivity to $(T/S)$ two orders of magnitude better than Planck. B-Pol would explore the entire parameter space spanned by the ``large-field'' inflationary models, and consequently is capable to measure the energy scale of inflation. If, as suggested by current CMB and large scale galaxy surveys, the power spectrum of primordial perturbations is not exactly scale invariant, then in a wide class of inflationary models the level of gravitational waves will lie well within the range probed by B-Pol."
    },
    "0808/0808.3556_arXiv.txt": {
        "abstract": "{} % {We obtained VLT/FLAMES+UVES high-resolution, fibre-fed spectroscopy of five young massive clusters (YMCs) in M83 (NGC~5236). This  forms the  basis of a pilot study testing the  feasibility of  using fibre-fed spectroscopy  to measure the velocity dispersions of several clusters simultaneously, in order to determine their dynamical masses. In principle, this reduces the telescope time required to obtain a statistically significant sample of dynamical cluster masses. These can be used to assess the long-term survivability of YMCs by comparing their dynamical and photometric masses, which are necessary to ascertain the potential evolution of YMCs into second-generation globular clusters. } {We adopted two methods for determining the velocity dispersion of the star clusters:  cross-correlating the cluster spectrum with the template spectra and  minimising a $\\chi^2$ value between the cluster spectrum and the broadened template spectra. We also considered both red giant and red supergiant template stars. Cluster~805 in M83  (following the  notation of Larsen) was chosen as a control to test the reliability of the results obtained by this observational method, through a comparison with the results obtained from a standard echelle VLT/UVES spectrum obtained by Larsen \\& Richtler. } {We find no dependence of the velocity dispersions measured for a cluster on the choice of red giant versus red supergiant templates, nor on the method adopted. However, we do find that the standard deviation of the results obtained with only one method may underestimate  the true uncertainty. We measure a velocity dispersion of $\\sigma_\\mathrm{los}\\,=\\,10.2\\,\\pm\\,1.1\\,\\mathrm{km\\,s}^{-1}$ for cluster~805 from our fibre-fed spectroscopy. This is in excellent agreement with the velocity dispersion of $\\sigma_\\mathrm{los}\\,=\\,10.6\\,\\pm\\,1.4\\,\\mathrm{km\\,s}^{-1}$ determined from the standard echelle UVES spectrum of cluster~805. Our FLAMES+UVES velocity dispersion measurement gives $M_\\mathrm{vir}\\,=\\,(6.6\\,\\pm\\,1.7)\\,\\times\\,10^5\\,M_\\odot$, % consistent with previous results. This value of the virial mass is a factor of $\\sim 3$ greater  than the cluster's photometric mass, indicating a lack    of virial equilibrium. However, based on its effective star formation efficiency, the cluster is likely to virialise, and may survive for a Hubble time, in the absence of external disruptive forces. Unfortunately, our observations of the other M83 star clusters have insufficient signal-to-noise  ratios  to determine robust cluster velocity dispersions.} {We find that reliable velocity dispersions can be determined from high-resolution, fibre-fed spectroscopy. The advantages of observing several clusters simultaneously outweighs the  difficulty of accurate galaxy background subtraction,  providing that the targets are chosen to  provide sufficient signal-to-noise ratios, and are much brighter than the galaxy background.} ",
        "introduction": "The production of luminous, compact star clusters is characteristic of  starburst galaxies. However, these clusters have been observed in several different types of galaxy, including normal spirals (see e.g. \\citealt{larsen99}, and references therein). These bright clusters are termed Young Massive Clusters (YMCs), and are usually defined as those star clusters that are younger than $\\sim 1$~Gyr, more massive than $\\sim\\,10^5\\,\\mathrm{M}_\\odot$ and  compact, with half-light radii of a few to a few tens of parsecs. The sizes, luminosities and masses of these YMCs are consistent with the properties expected of young globular clusters. This led to the suggestion that YMCs represent globular clusters at an early phase of their evolution \\citep{whitmore93,schweizer93}. This would mean that studying YMCs would be tantamount to observing globular clusters at the epoch of their birth.  Since globular clusters are amongst the oldest building blocks in our Galaxy, and are benchmarks for stellar and galactic evolution, understanding their formation and evolution is fundamental in developing our understanding of the fields of large-scale star formation, as well as how galaxies build up and evolve. Determining whether or not YMCs are proto-globular clusters, therefore, is a vital  step in gaining insights into these questions.   The simplest criterion that YMCs need to fulfil in order to evolve into globular clusters is that they must be able to survive for a large fraction of a Hubble time. Since a cluster requires sufficient low-mass stars to survive beyond a few Gyr, the  present-day mass function (PDMF) of a cluster can be used to assess its survivability. It may be possible to get a handle on the PDMF of a cluster by comparing its dynamical mass to its photometric mass, which is the mass predicted by evolutionary synthesis modelling based on the observed luminosity of the cluster and assuming an initial mass function (IMF). If we assume that the dynamical mass of a cluster represents its true mass, any discrepancies between the two masses can be attributed to adopting incorrect assumptions in the photometric mass determination, such as the cluster IMF. Even for a very young cluster, which is not in virial equilibrium, its virial mass can still be used to assess the survivability of the cluster. This is done by by determining how far out of virial equilibrium the cluster is, as quantified by its effective star-formation efficiency (eSFE; \\citealt{goodwin06_bast}). Thus, dynamical mass determinations can potentially be used to test the scenario with respect to the long-term survivability of  YMCs, (e.g. \\citealt{ho96a,ho96b}; see \\citealt{degrijs07_parm} for an overview and references therein).   Since the velocity dispersion of a cluster needs to be measured in order to determine the dynamical mass of  a cluster (see Section~\\ref{sec:veldisp}), high-resolution spectroscopy of the cluster is needed. Large samples of  dynamical cluster masses need to be obtained to determine whether  YMCs might be proto-globular clusters, and this can be very  expensive in terms of telescope time. Observing several clusters at once with a high-resolution, fibre-fed spectrograph  limits this cost. However, it is unclear whether a sufficiently accurate galaxy background subtraction can be achieved from only a few fibres placed at a different position to the star cluster. In this paper, we explore the efficiency of this approach, using FLAMES+UVES to observe clusters in M83. Similar attempts at velocity dispersion measurements have been made with FLAMES+GIRAFFE (e.g. \\citealt{mieske08}), which uses an order of magnitude more fibres than FLAMES+UVES. Since large numbers of fibres can be dedicated to galaxy background observations, background  subtraction is more straightforward through the use of a master  background spectrum. This, however, is at the expensive of at least a factor of two in spectral resolution, and a significantly reduced spectral range. This reduced spectral resolution would, however, not permit the measurement of the small cluster velocity dispersions  anticipated for YMCs ($\\sim\\,7-15\\,\\mathrm{km\\,s}^{-1}$). Furthermore, \\citet{mieske08} observed ultra-compact dwarf galaxies (UCDs), rather than clusters. These are much  brighter relative to the background, and have a much uniform background  than  a spiral like M83.        M83, also known as NGC~5236, is a prime example of a nearby spiral galaxy hosting several YMCs \\citep{larsen04_cat}. At a distance of $4.5\\,\\pm\\,0.3$~Mpc \\citep{thim03}, with its face-on inclination and multitude of clusters, M83 is an ideal site to test whether  fibre-fed spectroscopy can be used to measure the  velocity dispersions of several clusters simultaneously. A further advantage of using M83 as the testing ground for this method is that the dynamical masses of two of its clusters -- clusters~502 and 805 (nomenclature from \\citealt{larsen04_cat}) -- have already been measured from standard high-resolution echelle spectroscopy \\citep{larsen04_rich}. Therefore, these clusters can provide a control sample to assess the reliability of the results found from fibre-fed spectroscopy.     M83 is a metal-rich galaxy, with  a central oxygen abundance of $12\\,+\\,\\log{\\mathrm{O/H}}\\,=\\,8.94\\,\\pm\\,0.09$ \\citep{bresolin05}, nearly twice the solar value of $12\\,+\\,\\log{\\mathrm{O/H}}\\,=\\,8.69$ \\citep{asplund}.  The grand-design spiral has a companion, NGC 5253 \\citep{rogstad74}, with both galaxies hosting  intense starburst activity \\citep{calzetti99,tremonti01}. M83 has a nuclear starburst  and ongoing star formation within its spiral arms \\citep{elmegreen98,harris01} and  has a star formation rate\\footnote{This was determined   using the equation of \\citet{kennicutt98} and adopting the average $\\mathrm{H}\\alpha$ flux of the galaxy as $7.1\\,\\times\\,10^{-11}\\,\\mathrm{erg\\,s}^{-1}$ (Kennicutt, priv. comm.) and a representative extinction of $E(B-V)\\,\\sim\\,0.5$~mag \\citep{hadfield05}.  } of $4\\,\\pm\\,1\\,M_\\odot\\,\\mathrm{yr}^{-1}$. The galaxy hosts  in excess of 1000 Wolf-Rayet stars \\citep{hadfield05}.    This paper is structured as follows. In Section \\ref{sec:obs}, we present our observations of M83 and subsequent data reduction, as well as photometry of the clusters and galaxy background regions. In Section \\ref{sec:veldisp}, we determine the velocity dispersion of cluster~805 and compare it with the results obtained from the standard UVES  spectrum of \\citet{larsen04_rich}. We  discuss  the possibility of using fibre-fed spectroscopy  to obtain velocity dispersions for  other clusters  in  Section  \\ref{sec:dis}. Finally, we  summarise our findings in Section \\ref{sec:sum}. ",
        "conclusions": "\\label{sec:dis}   Using  high-resolution, fibre-fed spectroscopy to determine velocity dispersions  would have a major impact on producing a large sample of dynamical cluster mass determinations, by making more efficient use of telescope time. By having a large sample of dynamical cluster masses, it should become possible to assess how likely YMCs are to survive for a large fraction of a Hubble time and evolve into globular cluster-type objects.   The primary concern over this method is the question of how successfully the galaxy background could be subtracted from fibre spectroscopy of faint clusters. Of course, sky subtraction  is not a problem with bright stars, despite a small amount of flux `bleeding' from the bright star fibres into the sky fibres. For the fainter objects, galaxy background subtraction is more challenging. Simply scaling and subtracting the galaxy background fibre with the most similar intensity background to the cluster being considered introduces more noise into the spectrum. Taking an average of two or three scaled galaxy background fibres minimises this, although still introduces noise. Subtracting an average of two galaxy background fibres decreases the signal-to-noise ratio of the cluster\\,+\\,background spectrum by a factor of  $\\sim 2.5$ for cluster~805, the brightest cluster, and a  factor of $\\sim 4$ for both cluster~347, the faintest cluster, and cluster~645, which has the lowest contrast with the galaxy background. It seems that scaling the galaxy background in this way is a reasonable approximation to the local galaxy background. All three galaxy background fibres are very similar, considering the level of noise, in spite of spatial variations in galaxy background, due, for example, to differing stellar populations and extinction. Nevertheless, some sky lines are imperfectly subtracted in the cluster~805 spectrum, although these are within the level of the noise. A reduced $\\chi^2$ between the low- and mid-background spectra and the scaled, high-background spectrum yields values of $\\chi^2\\,=\\,0.9$ and $\\chi^2\\,=\\,0.7$, respectively. While adopting the noise in the spectra as the uncertainty in the fit seems to overestimate the true uncertainty, the relative similarity in the galaxy background spectra, given the large degree of noise, is demonstrated. An alternative method would be to fit the continuum of the  scaled, extracted galaxy  background fibre and subtract this fit from the cluster spectrum, manually removing the sky lines  from the cluster spectrum. However, the extra complications involved in this process were not warranted for our data, since an appropriate galaxy background was achieved from averaged galaxy background fibres.       Also, the signal-to-noise ratio was much higher in the UVES spectrum of \\citet{larsen04_rich}, for which a shorter integration time of 7500\\,s was used. This would appear to indicate that long-slit echelle spectroscopy is a more efficient use of telescope time, especially since two clusters could potentially be observed in the same slit. However, finding  cases where two clusters can be observed simultaneously can be difficult, due to the relatively short slit length of many   echelle sepctrographs (e.g. 11\\,arcsec for UVES). The signal-to-noise ratio in the galaxy-subtracted UVES+FLAMES cluster spectra was substantially lower than anticipated from the ESO exposure time calculator. This predicted a signal-to-noise ratio in excess of 30 at $\\sim\\,8000\\,\\textrm{\\AA}$, neglecting the high galaxy background. Even considering this additional noise does not account for the much lower observed cluster signal-to-noise ratios. The most likely explanation for the greatly reduced signal-to-noise ratio is is that the fibres were positioned off-centre. However, this needs to be investigated further.     As discussed in Section \\ref{sec:comparison}, the signal-to-noise ratio of cluster~805 is a good representation of the lowest signal-to-noise ratio required to measure robust velocity dispersions. Therefore, making a conservative assumption that $50\\%$ of the flux of cluster~805 was missed in our spectrum indicates that results could be obtained for a cluster half as bright as cluster~805 (i.e., $m_I\\,\\approx\\,16.7$~mag). Approximately 18 of the 159 YMCs in M83, and 109 of all 1358 YMCs, contained in the catalogue of  \\cite{larsen04_cat} fulfil this criterion.    Another problem with this method, specific to FLAMES+UVES, is the limited choice of spectral region, as dictated by the need to use  a standard UVES setup. This meant that one of the strong lines of the Ca~{\\sc ii} triplet was lost. Furthermore, much of the data in the near-infrared cannot be used to determine velocity dispersions, due to telluric features and the lower CCD response in this spectral region. Clusters also need to be selected carefully, with $M_\\mathrm{vir}\\,\\gtrsim\\,10^5\\,M_\\odot$ in order to have high enough velocity dispersions to be measured. For a 10-Myr-old cluster, this corresponds to $M_V  = -11.5\\,\\mathrm{mag}$   ({\\sc starburst99}; \\citealt{leitherer99}),  or $m_V \\sim 16.8\\,\\mathrm{mag}$ for a distance of 4~Mpc % and $E(B-V)\\,\\sim\\,0.1$\\,mag.    While the drawbacks of this method must still be borne in mind, we have shown here that reliable velocity dispersions can be measured from fibre-fed spectroscopy. Therefore, this method can potentially be used to obtain a large sample of dynamical masses, which can be used to address the issue of the nature of YMCs as proto-globular clusters."
    },
    "0808/0808.1765_arXiv.txt": {
        "abstract": "Field star BD+20~307 is the dustiest known main sequence star, based on the fraction of its bolometric luminosity, $\\sim$4\\%, that is emitted at infrared wavelengths.  The particles that carry this large IR luminosity are unusually warm, comparable to the temperature of the zodiacal dust in the solar system, and their existence is likely to be a consequence of a fairly recent collision of large objects such as planets or planetary embryos.  Thus, the age of BD+20~307 is potentially of interest in constraining the era of terrestrial planet formation.  The present project was initiated with an attempt to derive this age using the {\\it Chandra X-ray Observatory} to measure the X-ray flux of BD+20~307 in conjunction with extensive photometric and spectroscopic monitoring observations from Fairborn Observatory.  However, the recent realization that BD+20~307 is a short period, double-line, spectroscopic binary whose components have very different lithium abundances, vitiates standard methods of age determination.  We find the system to be metal-poor; this, combined with its measured lithium abundances, indicates that BD+20~307 may be several to many Gyr old.  BD+20~307 affords astronomy a rare peek into a mature planetary system in orbit around a close binary star (because such systems are not amenable to study by the precision radial velocity technique). ",
        "introduction": "Hundreds of main sequence stars are known with orbiting dusty debris disks.  The dust in most of these systems is cold and appears to be analogous to dust in the Sun's Kuiper Belt.  Only a small percentage of debris disks contain warm dust analogous to that in the Sun's asteroid belt and zodiacal cloud (e.g., Su et al. 2006; Rhee et al 2008;  Trilling et al. 2008; Smith et al 2008). For stars that are at least 100 Myr old, the warm dust luminosity (L$_{IR}$/L$_{bol}$) is typically a few times 10$^{-4}$ or less.  Smith et al (2008) discuss and review the significance of this warm, mid-infrared, emission in the context of the Sun's zodiacal dust cloud.  They note that {\\it Spitzer Infrared Observatory} photometry (Hines et al 2006, Bryden et al 2006) indicates that warm dust is detected at only 2$\\pm$2\\% of the observed systems.  A more recent, more comprehensive, {\\it Spitzer} survey (Trilling et al. 2008) yielded a similar small percentage of warm excess emission stars in a sample of 213 F- and G-type stars. And a comparable small percentage also applies to warm excess A-type stars of age $>$400 Myr (Su et al. 2006) which, as will be seen below, is the age range of interest in the present paper. In all of these surveys and in other published {\\it Spitzer} surveys, the typical mid-infrared luminosity is sufficiently small that the warm dust could be produced by collisions of asteroids; for example, in a massive asteroid belt perturbed by the gravitational field of a nearby giant planet.  In contrast, Rhee et al (2008) discuss the dustiest main sequence stars currently known: Pleiad HD23514 and, dustiest of all, the G0 field dwarf BD+20~307 (HIP8920; Song et al. 2005, Weinberger 2008, Weinberger et al. 2008).  The dust luminosity at these two stars is orders of magnitude larger than at other mid-infrared excess main sequence stars and collisions of asteroids seem quite inadequate to explain the data.  A plausible model to account for so much dust so close to these stars involves a fairly recent collision of two planets or planetary embryos in the terrestrial planet zone (Rhee et al. 2008).  BD+20~307 was originally presumed to be a single star, and a measurement of its (presumably youthful) age would have constrained the duration of the era during which major collisional events occur in youthful planetary systems.  Using photospheric lithium content, Galactic space motion, and an upper limit to the X-ray flux measured with the ROSAT All-Sky Survey, Song et al. (2005) estimated the age of BD+20~307 to be $\\sim$300 Myr.  To refine this estimate, we obtained X-ray measurements with the {\\it Chandra X-ray Observatory} and hundreds of photometric optical observations to measure the star's rotation period and age, employing the technique of gyrochronology (Barnes 2007).  While we were carrying out these observations, Weinberger (2008) obtained the surprising result that BD+20~307 is a short period ($\\sim$3.5 day) spectroscopic binary of two nearly identical stars.  The reason this was a surprise is that the Song et al. (2005) echelle spectrum showed no evidence of two stars; by a low probability ($<$10\\%) quirk of fate, their 2004 epoch spectrum was obtained at a time when the two stars had nearly the same radial velocities.  While {\\it Spitzer} surveys for debris disks around stars have often explicitly omitted binaries, Trilling et al. (2007) found small quantities of cool dust around a substantial fraction of binary stars.  However, by analogy with the discussion of single stars in the first paragraph above, dust around such binary stars can readily be accounted for by collisions of Kuiper Belt objects.  Upon learning of the Weinberger (2008) result, we initiated a program to monitor the radial velocity of BD+20~307.  Our various observations are described in the next Section.  This is followed by a discussion of the age and some other aspects of this exceptionally dusty binary star, which points to the existence of rocky planets in the terrestrial planet region. ",
        "conclusions": "Until new techniques are available to image or otherwise detect planets with masses and semi-major axes similar to those of the terrestrial planets of our solar system, we must rely on more indirect observations and models. One promising indirect approach is detection of dusty debris associated with terrestrial planet formation (e.g., Gorlova et al. 2007; Currie et al. 2007; Rhee et al. 2008; Lisse et al. 2008).  To delineate the full extent of the era of terrestrial planet formation it is necessary to identify main sequence stars with copious quantities of warm dust grains and ages of tens to many hundreds of millions of years.  Thus, the goal of our research project was to establish the age of BD+20~307 which, along with HD 23514 (an F6 member of the Pleiades), are the only two stars known with age $\\geq$100 Myr and strong indications of the existence of orbiting terrestrial planets (Rhee et al. 2008).  Based on the findings of our study, it now appears likely that BD+20~307 is an old binary star and its copious quantity of warm orbiting dust has no direct relationship with the era of planet formation.  Rhee et al. (2008)  present a moderately detailed discussion of the origin and fate of the dust at HD 23514 and BD+20~307.  Given the large quantity of dust at BD+20~307, its short lifetime in orbit, and its concentration to a narrow range of warm temperatures (Weinberger et al. 2008), it is hard to escape the conclusion that something with the mass of a terrestrial planet was involved in a catastrophic collision, independent of whether that collision took place between two isolated objects or inside some sort of massive asteroid belt.  A remarkable aspect of BD+20~307 is that this collision apparently has taken place so late in its history.  Perhaps, as suggested by the discussion of possible systemic radial velocity variations in Section 3, there is an unseen stellar or brown dwarf object present whose gravitational influence recently destabilized the orbits of rocky objects of planetary mass.  In any event, whatever its age, if the massive amounts of dust at BD+20~307 do point toward the presence of terrestrial planets, then this represents the first known example of planets of any mass in orbit around close binary stars -- because the precision radial velocity technique is not sensitive to planets in such systems.  Because of its only modest sensitivity, IRAS was quite limited in the distance out to which it could detect warm dust at low luminosity, main sequence, stars of K- and M-type.  Should they exist, such very dusty, low luminosity, stars will be detectable by the Wide-Field Infrared Survey Explorer (WISE) mission.  A substantial sample of very dusty stars of known age could delineate the era of terrestrial planet formation while perhaps revealing the existence of additional remarkable stars like BD+20~307."
    },
    "0808/0808.3283_arXiv.txt": {
        "abstract": "Unexplained periodic fluctuations in the decay rates of $^{32}$Si and $^{226}$Ra have been reported by groups at Brookhaven National Laboratory ($^{32}$Si), and at the Physikalisch-Technische-Bundesandstalt in Germany ($^{226}$Ra).  We show from an analysis of the raw data in these experiments that the observed fluctuations are strongly correlated in time, not only with each other, but also with the distance between the Earth and the Sun.  Some implications of these results are also discussed, including the suggestion that discrepancies in published half-life determinations for these and other nuclides may be attributable in part to differences in solar activity during the course of the various experiments, or to seasonal variations in fundamental constants. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.2550_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec_intro}  The interstellar medium (ISM) of the Milky Way is a complex, multi-phase environment. Much of the ISM is an ionised gas at temperatures of $10^4-10^6$~K, requiring one or more vast sources of ongoing energy injection. In the modern picture of the ISM \\citep[see][for a review]{fer01}, ionised gas consists of two main phases: a hot ``coronal'' component at temperatures of $\\sim10^5-10^6$~K, and a warm ionised medium (WIM) of temperature $\\sim10^4$~K.  The WIM is most easily identified in H$\\alpha$ and other optical recombination lines, which show that much of this gas corresponds to \\HII\\ regions around massive stars. However, most \\HII\\ regions are seen only at low Galactic latitudes. Further from the plane, pulsar dispersion measures (DMs), free-free absorption of low-frequency Galactic synchrotron emission, interstellar scattering of compact radio sources, and faint H$\\alpha$ and H$\\beta$ emission all demonstrate the existence of a widespread, diffuse WIM, not associated with individual stars \\citep[see overviews by][]{kh87,rey90c}.  In the rest of this paper, we use WIM to refer only to this diffuse component, disregarding individual \\HII\\ regions.  The diffuse WIM has a typical volume-averaged free electron density $n \\approx 0.01-0.1$ cm$^{-3}$ and an electron temperature $T_e \\approx 8000$~K \\citep[e.g.,][]{rey90c,wsx+08}.  Scattering and dispersion of high-latitude pulsars demonstrate that the diffuse component of the WIM is distributed in a thick disk, extending more than a kpc above and below the Galactic plane with a roughly exponential fall-off in density \\citep{rd75,rey89,nct92b}. The power required to maintain the ionisation state of the thick-disk WIM is $\\sim5\\times10^{41}$~ergs~s$^{-1}$, a vast rate of energy injection that can seemingly only be supplied by the photo-ionising flux of all the hot stars in the Galaxy (Reynolds \\etal\\ 1984, 1990b\\nocite{rey84,rey90b}). Questions remain, however, over how these ultraviolet photons propagate to large distances above the Galactic plane \\citep{rey90,mc93}, and whether the spectrum of the photon field produced by massive stars is consistent with the observed line ratios in the WIM \\citep{rt95b,hklr96,mrh06}.  Many other edge-on spiral galaxies also have a faint, thick WIM disk, seen in deep H$\\alpha$ observations \\citep{ran98,det04}.  All these results make clear that the WIM is a key part of the feedback process between the stellar and gaseous components of a galaxy, and that it connects the energetics and dynamics of the disk and the halo. There is thus considerable motivation to understand the density, spatial distribution and extent of the WIM in the Milky Way and in other galaxies.  Our own Galaxy is unique in that the DMs of radio pulsars can provide the free-electron column of the WIM for thousands of lines of sight, allowing one to develop a smooth underlying model for the three-dimensional Galactic distribution of $n$ \\citep[e.g.,][]{tc93,gbc01}.  Integrating such a distribution allows one to also predict the emission measure (EM) along any sightline. However, even when interstellar extinction is taken into account, the predicted EMs are vastly smaller than those inferred from diffuse H$\\alpha$ emission, from low-frequency absorption in the spectra of non-thermal radio sources, or from free-free radio emission \\citep[e.g.,][]{pw02,srwe08}. Because DMs have a linear dependence on the free electron density while EMs depend on the square of the electron density, the inference is that even far from the plane, the WIM is not a smooth, homogeneous medium, but is clumped into discrete clouds or layers \\citep{rey77,rey91,pyn93}. In the following sections, we review some recent attempts to model this structure.  \\subsection{Modeling The WIM}  \\subsubsection{The NE2001 Model} \\label{sec_ne2001}  Recognising the limitations of a smooth, homogeneous WIM, Cordes \\& Lazio (2002, 2003\\nocite{cl02,cl03}) developed a detailed model, termed ``NE2001'', for the structure of ionised gas in the Galaxy. The NE2001 model consists of a thin disk of scale height 140~pc associated with low-latitude \\HII\\ regions; a thick-disk WIM layer as discussed above, with a scale height of 950~pc; large-scale structure associated with spiral arms; cavities and density enhancements corresponding to known features in the local ISM; and individual clumps and voids needed to model specific regions of enhanced/reduced scattering or dispersion.  The NE2001 model can predict the DM, EM, angular broadening, scintillation bandwidth and other parameters for any sightline integrated to any distance. It is widely used to estimate distances to radio pulsars for which there is no other distance indicator other than their DMs.  The NE2001 model attempts to provide a comprehensive description of the WIM, and can be continually improved as more DMs and other data are obtained. However, as \\cite{cl02} acknowledge, the NE2001 model fails to correctly predict the DMs of some pulsars in high-latitude globular clusters.  For example, the 22 pulsars in the globular cluster 47~Tuc have an average DM of 24.4~pc~cm$^{-3}$ and are at a distance of $\\approx5$~kpc \\citep{clf+00,gbc+03}. Pulsars in this cluster are included as an input to the NE2001 model, but the prediction of NE2001 is DM~$=42.1$~pc~cm$^{-3}$, in significant disagreement with the observations.  More broadly, \\cite{lfl+06} have shown that if NE2001 is applied to the new pulsars subsequently discovered in the various surveys that used the Parkes multibeam receiver, the inferred vertical distribution of these sources in the Galaxy has a scale height about a factor of two lower than expected.  The NE2001 model also has other difficulties at high latitudes.  At the north Galactic pole, NE2001 predicts\\footnote{We have used the software provided at {\\tt http://www.astro.cornell.edu/$\\sim$cordes/NE2001/}, v1.0.} EM~$\\approx 0.05$~pc~cm$^{-6}$. This is more than order of magnitude smaller than that observed \\citep{rey91,pw02,hrt+03}, suggesting that there is substantial clumping of the WIM at large distances from the Galactic plane, not included in NE2001 \\citep[see also discussion by][]{rey91,rey97}.  Meanwhile, \\cite{srwe08} have attempted to simultaneously reconcile extragalactic Faraday rotation measurements and all-sky maps of Galactic radio emission with the distribution of $n$ predicted by NE2001. They show that a joint fit to all these data requires an anomalously strong halo magnetic field, and an unrealistically low scale-height for Galactic cosmic rays. \\cite{srwe08} argue that if the scale-height of the thick-disk WIM in the NE2001 model were increased from 950~pc to $\\sim2$~kpc, a self-consistent model for the Galactic distribution of thermal gas, magnetic fields and cosmic rays could then be derived.  \\subsubsection{The BMM06 Model} \\label{sec_bmm06}  The arguments in \\S\\ref{sec_ne2001} suggest that the scale height and volume filling factor of the thick-disk WIM need to be re-evaluated. A detailed study of this issue has been carried by \\cite{bmm06}, hereafter BMM06, who used measurements of EM and DM for 157 pulsars to derive statistics on the density, scale height and filling factor of the WIM. BMM06 argued that both the internal cloud electron density, $N$ and volume filling factor, $f$, of the clumpy WIM evolve with a source's vertical distance, $z$, above or below the Galactic plane. They fit the data with a model in which $N$ decays exponentially, while $f$ grows exponentially, with $z$. The expressions for their best fits are: \\begin{equation} N(z) = N_0~e^{-\\frac{z}{H_N}} ,~ f(z) = f_0~e^{+\\frac{z}{H_f}} , \\label{eqn_bmm1} \\end{equation} with scale heights \\begin{equation} H_N = 710_{-120}^{+180}~{\\rm pc},~ H_f = 670_{-130}^{+200}~{\\rm pc}, \\label{eqn_bmm2} \\end{equation} and extrapolated mid-plane values \\begin{equation} N_0 = 0.41\\pm0.06~{\\rm cm}^{-3}, f_0 = 0.05\\pm0.01 . \\label{eqn_bmm3} \\end{equation}  The resulting volume-averaged density of the thick-disk WIM, $n \\equiv fN$, remains roughly constant at $n \\approx 0.02$~cm$^{-3}$ up to $z \\sim 1$~kpc, implying a much larger effective scale height for the WIM layer than had been assumed by all previous authors. At higher $z$, BMM06 argue that the filling factor takes on a constant value $f \\ga 0.3$, while $N$ continues to decay.  As discussed by \\cite{srwe08}, the NE2001 model cannot successfully predict the free-free emission seen at 23~GHz by the {\\em WMAP}\\ experiment, but incorporating the function $f(z)$ derived by BMM06 (see Eqn.~[\\ref{eqn_bmm1}] above) produces a good match to the data.  The BMM06 model clearly offers a number of improvements over previous work in its description of the WIM at high $z$. However, there are several important caveats and limitations associated with that study. First, for 95\\% of the pulsars used by BMM06, no independent estimates of distance (or hence of $z$) were available, so BMM06 used the distances predicted by the NE2001 model, and assumed errors of 20\\% in these distances.  However, as we noted in \\S\\ref{sec_ne2001}, the NE2001 model can be in error by much more than this at high Galactic latitudes. Conversely, BMM06 explicitly excluded from their sample all pulsars in globular clusters or in the Magellanic Clouds, eliminating a substantial number of sources with reliable, independent distances.\\footnote{A subsequent study by \\cite{bm08} has repeated the analysis of BMM06 but only using pulsars with independent distance estimates.  However, for reasons that they do not explain, \\cite{bm08} only consider 13 of the 25 globular clusters known to contain radio pulsars, and do not use information on pulsars in the Magellanic Clouds.}  A second limitation of BMM06 \\citep[and also of the subsequent study by][]{bm08} is that their data sets consist of paired estimates of DM and EM for each pulsar. While the DM is a direct integral of $n$ from the observer to the source, the EMs are derived from H$\\alpha$ data, which correspond to integrals to infinity, attenuated by dust extinction. BMM06 scaled down the EM values to account for the finite distance of each pulsar, but this calculation propagates the significant distance uncertainties in NE2001 into the EM values. Furthermore, scaling the EMs requires one to assume a WIM scale height, introducing a degeneracy into their subsequent fitting of the EMs to derive this same scale height. BMM06 corrected the EMs for extinction using standard reddening laws, but such corrections assume that dust and gas are well mixed along the line of sight. A inhomogeneous dust distribution, combined with significant distance uncertainties for most pulsars, can result in an erroneous extinction correction.  Finally, BMM06 considered pulsars with $60^\\circ < \\ell < 360^\\circ$ and $|b| > 5^\\circ$ (where $\\ell$ and $b$ are Galactic longitude and latitude, respectively). As we will show in \\S\\ref{sec_distrib}, the DMs of many pulsars at $|b| < 40^\\circ$ are contaminated by \\HII\\ regions and by other low-latitude structure.  When trying to extract the properties of the WIM in the thick disk, the DMs and EMs of such pulsars will increase the scatter in the data and will reduce the signal-to-noise ratio of the underlying signal.  Overall, we conclude that while the approach of BMM06 provides several new insights into the behaviour of the WIM as a function of $z$, there are large systematic uncertainties that need to be incorporated into the parameters they extract from the data.  \\subsubsection{Scope Of This Paper}  The discussion in \\S\\ref{sec_ne2001} and \\S\\ref{sec_bmm06} illustrates the strengths and weaknesses of the NE2001 and BMM06 models. Specifically, NE2001 is a detailed description of ionised gas along existing sightlines, but has limited predictive power at high $z$ where the number of existing sightlines is sparse. BMM06 provide a comprehensive joint analysis of EMs and DMs, and lay out a functional form for the $z$-dependence of the thick-disk WIM that is a significant improvement over that of NE2001. However, their analysis suffers from the significant uncertainties introduced by adopting distances from NE2001, and by trying to correct EMs for extinction and for the finite path length to the pulsars in their sample.  In this paper, we complement these previous studies, with a simple attempt to model the overall structure of the diffuse WIM. Our main thrust is to re-apply the analysis of \\cite{rey89}, who plotted DM$_\\perp\\equiv$ DM$.\\sin|b|$ vs. $z$ for $\\sim35$ pulsars with known distances, and compared this to the predictions of simple geometric models. While this approach has been repeated in various subsequent papers \\citep[e.g.,][]{nct92b,gbc01,bm08}, the much larger sample of pulsars now available affords us the luxury of being able to filter the data to avoid contamination by the clumps, spiral arms and \\HII\\ regions incorporated in NE2001, and to thus boost any signal produced by the truly diffuse WIM.  In \\S\\ref{sec_dist}, we review the sample of pulsars with known, reliable, distances and define our sample of 53 sightlines.  In \\S\\ref{sec_analysis} we use a cut on these DMs in Galactic latitude to derive a substantially revised scale-height of the thick-disk WIM, and then combine these DMs with separate data on EMs (i.e., not matched to pulsar sightlines, as carried out by BMM06 and Berkhuijsen \\& M\\\"uller 2008\\nocite{bm08}) to confirm the contention of BMM06 that the data are most simply explained if the volume filling factor of this layer evolves with~$z$.  In \\S\\ref{sec_disc} we discuss the implications of our calculations for the distances to pulsars, for the structure of the WIM, and for the transition from warm to million-degree gas in the Galactic halo. ",
        "conclusions": "We have used the dispersion measures of pulsars at known, reliable, distances and the emission measures traced by diffuse H$\\alpha$ emission to re-evaluate the structure and distribution of the thick disk of warm ionised gas in the Milky Way. Our main conclusions are as follows: \\begin{enumerate}  \\item The DMs of high-latitude pulsars are well fit by a plane-parallel distribution of free electrons, in which the volume-averaged electron density, $n$, decays exponentially with height, $z$, above the Galactic plane.  \\item Within a few kpc of the Sun, the exponential scale height for $n$ is $H_n\\approx1800$~pc, approximately double the value determined in most previous studies.  Estimates of distances to high-latitude pulsars using the ``NE2001'' electron model of Cordes \\& Lazio (2002, 2003\\nocite{cl02,cl03}) are thus likely to be substantial underestimates.  \\item Within a few kpc of the Sun, the mid-plane volume-averaged electron density of the diffuse WIM is $n_0 = 0.014\\pm0.001$~cm$^{-3}$, and the corresponding total vertical column density of the WIM is $\\approx26$~pc~cm$^{-3}$.  Assuming a total vertical emission measure EM$_0 = 1.9\\pm0.3$~pc~cm$^{-3}$, clouds in the thick-disk component of the WIM have a typical internal electron density of $N_0 = 0.34\\pm0.06$~cm$^{-3}$ at mid-plane, with a volume filling factor $f_0 = 0.04\\pm0.01$.  \\item Comparison of DMs and EMs shows that the internal electron density, $N$, of individual ionised clouds within the thick-disk WIM decays rapidly with $z$, with an exponential scale height $H_N \\approx 500$~pc.  The differing behaviour for $H_n$ and $H_N$ requires that the filling factor of the WIM grows exponentially with $z$ over the range $0 \\le z \\la 1.4$~kpc, with a scale height $H_f \\approx 700$~pc.  \\item The filling factor of the thick-disk WIM probably reaches a maximum of $f\\approx 0.3$ at $z \\approx 1.5$~kpc. At larger distances from the Galactic plane, the ISM becomes increasingly dominated by hotter, coronal, gas in the halo.  However, hot low-density halo gas makes only a negligible contribution to observed pulsar DMs.  \\end{enumerate}  A variety of future observations will allow us to test the above conclusions, and to better understand the structure of the diffuse WIM. Most immediately, from late 2008 WHAM will be relocated to the southern hemisphere, allowing studies similar to those carried out on the Perseus spiral arm to now be repeated for other regions.  Further improvements will come from an increase in the number of high-latitude radio pulsars with direct distance estimates.  For example, only $\\approx25\\%$ of the Galactic globular clusters at $|b| \\ge 40^\\circ$ contain known radio pulsars.  While some clusters at these latitudes have been searched for pulsations without success \\citep{and93,bl96,hrs+07}, at least some of these targets probably contain as yet undetected pulsars \\citep[e.g., NGC~5634;][]{shq02}, and should be searched more deeply. Furthermore, there are several high-latitude pulsars with 1.4~GHz radio fluxes at the level of a few mJy, which is sufficiently bright to derive interferometric parallaxes through future observations.  Most surveys for radio pulsars have concentrated on the Galactic plane, and those surveys that have covered high-latitude regions have been of comparatively modest sensitivity \\citep[e.g.,][]{mld+96}.  Forthcoming all-sky pulsar surveys with more sensitive instrumentation will greatly expand the sample of high-$|b|$ pulsars, whose DMs can then be used to better map out the structure of the WIM.  Finally, we note that the new model we have presented for the WIM can also be used to probe Galactic magnetic fields at high $z$.  In a subsequent paper, we will combine the density distribution derived here with Faraday rotation measures of a large number of polarised extragalactic sources at high $|b|$, with which we aim to derive the strength and geometry of magnetic fields in the Galactic halo."
    },
    "0808/0808.3774_arXiv.txt": {
        "abstract": "We present high-resolution Very Large Array imaging of the molecular gas in the host galaxy of the high redshift quasar \\bri\\ ($z=4.41$). Our \\bco\\ observations have a linear resolution of 0.15$''$ (1.0\\,kpc) and resolve the molecular gas emission both spatially and in velocity. The molecular gas in \\bri\\ is extended on scales of 5\\,kpc, and shows a complex structure. At least three distinct components encompassing about two thirds of the total molecular mass of 9.2$\\times$10$^{10}$\\,\\msol\\ are identified in velocity space, which are embedded in a structure that harbors about one third of the total molecular mass in the system. The brightest \\bco\\ line emission region has a peak brightness temperature of 61$\\pm$9\\,K within 1\\,kpc diameter, which is comparable to the kinetic gas temperature as predicted from the CO line excitation.  This is also comparable to the gas temperatures found in the central regions of nearby ultra-luminous infrared galaxies, which are however much more compact than 1\\,kpc. The spatial and velocity structure of the molecular reservoir in \\bri\\ is inconsistent with a simple gravitationally bound disk, but resembles a merging system.  Our observations are consistent with a major, gas-rich (`wet') merger that both feeds an accreting supermassive black hole (causing the bright quasar activity), and fuels a massive starburst that builds up the stellar bulge in this galaxy.  Our study of this $z$$>$4 quasar host galaxy may thus be the most direct observational evidence that `wet' mergers at high redshift are related to AGN activity. ",
        "introduction": "Great progress has been made in recent years both observationally and theoretically to further our understanding of galaxy formation and evolution from the early epochs of galaxy formation to the present-day universe.  One basic prediction of cosmological simulations is that during the early epoch of hierarchical galaxy formation, some of the most massive galaxies are already formed in major merger events (e.g., Springel \\etal\\ \\citeyear{spr05}). These mergers are believed to commonly trigger both AGN and starburst activity in such early systems, which is regulated via AGN feedback (e.g., Hopkins \\etal\\ \\citeyear{hop05}). Such feedback may be responsible for the present-day ``$M_{\\rm BH}$--$\\sigma_v$'' relation between black hole mass and bulge velocity dispersion (Ferrarese \\& Merritt \\citeyear{fer00}; Gebhardt \\etal\\ \\citeyear{geb00}).  Studies of molecular gas (the requisite material to fuel star formation) in young galaxies are an essential ingredient to understand their physical properties in more detail. Molecular gas (typically CO) has been detected in $\\sim$50 galaxies at $z$$>$1 to date, revealing large molecular reservoirs of $>$10$^{10}$\\,\\msol\\ in most cases (see Solomon \\& Vanden Bout \\citeyear{sv05} for a review). However, these studies rely almost exclusively on the integrated properties of the line emission, as the molecular reservoirs in these distant galaxies are difficult to resolve. To date, only the $z$=4.69 and $z$=6.42 quasars BR\\,1202--0725 and SDSS\\,J1148+5251, two of the most distant gas-rich galaxies could be resolved in molecular gas emission without the aid of gravitational magnification (Omont \\etal\\ \\citeyear{omo96}; Carilli \\etal\\ \\citeyear{car02}; Walter \\etal\\ \\citeyear{wal04}).  In this letter, we report on high angular resolution (0.15$''$; 1.0\\,kpc) Very Large Array (VLA)\\footnote{The Very Large Array is a facility of the National Radio Astronomy Observatory, operated by Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.}  observations of molecular gas in the host galaxy of BRI\\,1335--0417, a dust-rich, optically identified quasar at a redshift of 4.41, corresponding to only 1.4 Gyr after the Big Bang. Optical imaging with the Hubble Space Telescope reveals a single point source without any evidence for gravitational lensing (Storrie-Lombardi \\etal\\ \\citeyear{sl96}).  We use a concordance, flat $\\Lambda$CDM cosmology throughout, with $H_0$=71\\,\\kms\\,Mpc$^{-1}$, $\\Omega_{\\rm M}$=0.27, and $\\Omega_{\\Lambda}$=0.73 (Spergel \\etal\\ \\citeyear{spe07}). ",
        "conclusions": "\\begin{figure*} \\plotone{f4.eps} \\vspace{1.5mm}  \\caption{RGB composite color map representation and position-velocity ($p-v$) diagram of the \\bco\\ velocity channels shown in Fig.~\\ref{f3}, with three colors encoding the velocity range of the emission [red: redshifted (channels 1-2 in Figure \\ref{f3}), green: central (channels 3-5), blue: blueshifted (channels 6-7)]. The velocity range covered (relative to the line center) in \\kms\\ is indicated in the top left corner, and the linear scale is indicated in the bottom left corner of each panel. {\\em Left}: $p-v$ diagram along a position axis of 189$^\\circ$, connecting the two brightest peaks of the integrated line emission.  {\\em Middle}: All velocity channels (1-7). The dashed line indicates the orientation of the $p-v$ slice. {\\em Right}: All except the central channel (channel 4). \\label{f4}} \\end{figure*}  We present maps of the molecular gas distribution in a quasar host galaxy at $z$=4.4 at 1.0\\,kpc spatial resolution. The VLA data show that the molecular gas reservoir in this galaxy is not only distributed over a scale of $\\sim$5\\,kpc, but also structured in velocity space.  This is the first time that the molecular gas in a quasar host galaxy at such a high redshift (or, indeed, at any redshift greater than $\\sim$0) has ever been resolved both spatially and dynamically over multiple beams (except for the lensed system PSS\\,J2322+1944 at $z$=4.12; Riechers et al.\\ \\citeyear{rie08}). The emission in \\bri\\ is resolved into at least three distinct components harboring at least 1--2 $\\times$ 10$^{10}$\\,\\msol\\ of molecular gas each, and at least another 2 $\\times$ 10$^{10}$\\,\\msol\\ is found in the more diffuse molecular medium in between these concentrations. Each of the subcomponents hosts a few times the molecular gas mass of nearby ULIRGs such as Arp\\,220 (Downes \\& Solomon \\citeyear{ds98}). The brightest peak corresponds to a 1\\,kpc diameter region with a gas surface density of $\\Sigma_{\\rm H_2}^{\\rm peak}$=2.3 $\\times$ 10$^{10}$\\,\\msol\\,kpc$^{-2}$, which is by a factor of a few higher than in in the central regions of typical ULIRGs (e.g., Wilson et al.\\ \\citeyear{wil08}). The total molecular gas mass and distribution are reminiscent of those in the $z$=4.69 and $z$=6.42 quasars BR\\,1202--0725 and SDSS\\,J1148+5251 (Omont \\etal\\ \\citeyear{omo96}; Carilli \\etal\\ \\citeyear{car02}; Riechers \\etal\\ \\citeyear{rie06}; Walter \\etal\\ \\citeyear{wal04}), albeit concentrated on almost an order of magnitude smaller scales than in BR\\,1202--0725. The brightest peak of the \\bco\\ emission traces a region of 1\\,kpc diameter where the gas has an average temperature of at least 61\\,K (depending on optical depth and beam dilution), which is at the high end of but comparable to temperatures in the central regions of nearby ULIRGs (e.g., Downes \\& Solomon \\citeyear{ds98}). However, in ULIRGs, regions with molecular gas temperatures of 30--60\\,K are much more compact than 1\\,kpc. The CO excitation ladder in \\bri\\ is not constrained well at present. However, the excitation so far is consistent with being similar to the $z$=4.69 quasar host galaxy of BR\\,1202--0725, for which models of collisional excitation predict a kinetic gas temperature of $T_{\\rm kin}$=60\\,K and a median gas density of $n({\\rm H_2}) = 10^{4.1}\\,$cm$^{-3}$ (Riechers et al.\\ \\citeyear{rie06}). If indeed $T_{\\rm kin} \\simeq T_{\\rm b}^{\\rm peak}$, this would indicate that the brightest CO peak is close to being resolved. Also, it appears that the gas temperature and density are on the high end of but comparable to ULIRGs, while the total gas mass is by about an order of magnitude higher.  Based on the FIR luminosity and SED shape of the source ($L_{\\rm FIR}$=3.1$\\times$10$^{13}$\\,\\lsol; Benford et al.\\ \\citeyear{ben99}), a star formation rate (SFR) of 4650\\,\\msol\\,yr$^{-1}$ and a total dust mass of 2$\\times$10$^9$\\,\\msol\\ can be derived.\\footnote{Assuming SFR=1.5$\\times$10$^{-10}\\,L_{\\rm FIR}$(\\msol\\,yr$^{-1}$/\\lsol) (Kennicutt \\citeyear{ken98a}).} This does not account for possible heating of the dust by the AGN. However, if the dust in this source has an extension and non-uniform distribution similar to that of the molecular gas, local heating (rather than heating of a central source) is likely to be responsible for most of the dust emission. This is consistent with the finding that the (rest-frame) 4.0\\,cm radio continuum emission in this source dominantly originates from extended, kpc-scale regions and shows a peak $T_{\\rm b}$ of only few times 10$^5$\\,K (Momjian et al.\\ \\citeyear{mom07}). Such $T_{\\rm b}$ are typical for relativistic electrons from supernova remnants and the interstellar medium in nuclear starbursts, but by at least 2 orders of magnitude lower than found in AGN-powered environments. This finding is also reflected in the fact that \\bri\\ follows the radio-FIR correlation for star-forming galaxies (Carilli et al.\\ \\citeyear{car99}).  The FIR-derived SFR is by about an order of magnitude higher than in local ULIRGs; however, the ratio of SFR and total gas mass is consistent with the scaling relation for ULIRGs (Solomon et al.\\ \\citeyear{sol97}) and other high-$z$ galaxies (Solomon \\& Vanden Bout \\citeyear{sv05}; Riechers et al.\\ \\citeyear{rie06}). Moreover, in the centers of nearby ULIRGs, $\\Sigma_{\\rm H_2}^{\\rm peak}$ correlates with the SFR (Wilson et al.\\ \\citeyear{wil08}), which suggests that an increased availability of fuel for star formation at a certain density leads to an increased SFR. Given its high $\\Sigma_{\\rm H_2}^{\\rm peak}$ on kpc scales, this also motivates the high derived SFR for \\bri. Assuming that the gas is converted into stars at an efficiency of 5--10\\% as in giant molecular cloud cores (e.g., Myers et al.\\ \\citeyear{mye86}; Scoville et al.\\ \\citeyear{sco87}), the SFR corresponds to a gas depletion timescale of 2--4 $\\times$ 10$^8$\\,yr.  If the molecular gas in this system was gravitationally bound, the CO linewidth (see Sect.~2) and distribution would predict a dynamical mass of $M_{\\rm dyn}$=1.0 $\\times$ 10$^{11}\\,\\sin^{-2}\\,i$\\,\\msol, which could account for both the total molecular gas mass and the mass of the black hole of $M_{\\rm BH}$=6 $\\times$ 10$^9$\\,\\msol\\ (Shields \\etal\\ \\citeyear{shi06}), but not for a substantial fraction of a $\\sim$4$\\times$10$^{12}$\\,\\msol\\ stellar bulge as predicted if the local $M_{\\rm BH}$--$\\sigma_v$ relation were to hold (Ferrarese \\& Merritt \\citeyear{fer00}; Gebhardt \\etal\\ \\citeyear{geb00}).\\footnote{In principle, assuming an extreme inclination toward the line-of-sight would significantly increase $M_{\\rm dyn}$; however, the implied large CO linewidths suggest that such a dynamical structure would not be virialized (Riechers et al.\\ \\citeyear{rie08}).}  However, the overall structure of \\bri\\ looks rather disturbed, both spatially and in its velocity structure. While the general north-south extension of the source may be in agreement with rotating structure, this is not the case for the central part of the emission, in particular the compact peak seen in the central channel. This, and the structure seen in the high-resolution \\bco\\ map (Figure \\ref{f2}), are more reminiscent of the disturbed gas reservoirs in major mergers, such as seen in the nearby Antennae (NGC\\,4038/39, e.g., Wilson \\etal\\ \\citeyear{wil00}). If \\bri\\ was an interacting system, the northern component may be in the process of merging with the southern component, which likely hosts the luminous quasar. The connecting component then may be the part where the two galaxies overlap and merge. Note that this region alone would be massive enough to host more than 10 of the largest molecular complexes found in the overlap region of the Antennae (Wilson \\etal\\ \\citeyear{wil00}).  Cosmological simulations predict that the merger rates at $z$=4.4 are substantially higher than at $z$=0. Such scenarios imply that during the early epoch of hierarchical galaxy formation, some of the most massive galaxies form in major merger events (e.g., Springel \\etal\\ \\citeyear{spr05}). Such major, `wet' mergers are believed to commonly trigger both AGN and starburst activity, and lead to high excitation of the molecular gas during both the hierarchical buildup of the host galaxy and the quasar phase (e.g., Narayanan \\etal\\ \\citeyear{nar07}). In simulations, such objects often show multiple CO emisson peaks, arising from molecular gas concentrations that have not yet fully coalesced.  We conclude that the observed properties of \\bri\\ (AGN and extreme starburst activity, high CO excitation, disturbed morphology over 5\\,kpc scales) are connected to the ongoing buildup of the quasar host galaxy. Such a signpost of early galaxy assembly then could be considered direct observational proof of the scenarios proposed by cosmological simulations, and enable us to directly investigate the connection between quasar activity and high-mass merger events at early cosmic times."
    },
    "0808/0808.4154_arXiv.txt": {
        "abstract": "We  have studied the relationship between  the high- and low-ionization [O~IV]~$\\lambda$25.89~$\\micron$, [Ne~III]~$\\lambda$15.56~$\\micron$  and  [Ne~II]~$\\lambda$12.81~$\\micron$ emission lines with the aim of constraining the active galactic nuclei (AGN) and star formation contributions for a sample of 103 Seyfert galaxies. We used the [O~IV] and [Ne~II] emission as  tracers for the AGN power and star formation  to  investigate  the ionization state of the emission-line gas. We find  that Seyfert 2 galaxies have, on average, lower [O~IV]/[Ne~II]  ratios than those of Seyfert 1 galaxies. This result suggests two possible scenarios: 1) Seyfert 2 galaxies have intrinsically weaker AGN, or 2) Seyfert 2 galaxies have  relatively higher star formation rates than Seyfert 1  galaxies. We   estimate  the fraction of [Ne~II] directly associated with the AGN and  find that Seyfert 2 galaxies have  a larger contribution from star formation, by a factor of $\\sim 1.5$ on average, than what is found in Seyfert 1 galaxies. Using the stellar component  of [Ne~II] as a tracer of the current star formation we found similar  star formation rates in Seyfert 1 and Seyfert 2 galaxies. We examined the mid- and far-infrared continua  and  find that [Ne~II] is well correlated with the continuum luminosity at 60$\\micron$ and that both [Ne~III] and [O~IV] are better correlated with the 25$\\micron$ luminosities than with the continuum at longer wavelengths, suggesting that the mid-infrared continuum luminosity is dominated by the AGN, while  the far-infrared  luminosity is dominated by star formation. Overall, these results test the unified model of AGN, and suggest that the differences between Seyfert galaxies cannot be  solely due to  viewing angle dependence. ",
        "introduction": "Active Galactic Nuclei (AGN) are thought to harbor massive black holes surrounded by an accretion disk  responsible for the enormous energy rates observed in their unresolved nuclei \\citep{1984ARA&A..22..471R,2000ApJ...540L..13P,2004ApJ...613..682P}. Historically, Seyfert 1 and Seyfert 2 galaxies have been classified by the presence or absence of broad optical emission lines. In this regard, Seyfert 1 galaxies have  broad permitted  (FWHM$\\sim$ 1-5$\\times{\\rm 10^3~km~s^{-1}}$) and narrow (FWHM$\\sim$ 5$\\times{\\rm 10^2~km~s^{-1}}$) permitted and forbidden lines and  Seyfert 2 galaxies have only narrow permitted and forbidden line emission \\citep{1974ApJ...192..581K}. Using spectropolarimetry, \\cite{1985ApJ...297..621A} found broad Balmer lines and [Fe~II]  emission in the polarized spectrum of the Seyfert 2 \\objectname{NGC 1068} galaxy, characteristic of a Seyfert 1 spectrum. This represents the first observational evidence in favor of a unified model. In this model, Seyfert 1 and Seyfert 2 galaxies are intrinsically the same with their differences attributed to viewing angle. In Seyfert 2 galaxies, our line of sight to the broad line region (BLR) and the central engine  is obstructed by an optically thick  dusty torus-like structure, while in Seyfert 1 galaxies, our line of sight is not obstructed by the torus, allowing a direct view of the central regions of the active galaxy.   Although it has been observationally confirmed, the unified model for Seyfert galaxies does not address  the role of stellar activity. The fact that active galaxies can  also host massive star-forming regions \\citep[e.g,][]{1990MNRAS.242..271T,1997ApJ...485..552M,1998ApJ...493..650M,2004MNRAS.355..273C,2006MNRAS.366..480G,2007ApJ...671.1388D} suggests a connection between the AGN and star formation in the proximity of the super-massive black hole, typically on scales of a few hundred parsecs. These starbursts may have a significant impact on the fueling of the central black hole \\citep[e.g.,][]{1999MNRAS.303..173S,2007ApJ...671.1388D}. Many authors have suggested that violent star formation, in the circumnuclear region,  plays a fundamental role in the energetics of Seyfert 2 galaxies \\citep[e.g.,][]{1985MNRAS.213..841T,1998ApJ...493..650M,1995MNRAS.272..423C,2001ApJ...546..845G,2001ApJ...558...81C}. This extra source of energy, a young  stellar population in the vicinity of the nucleus,  could complement  the non-stellar component associated with  the AGN. \\cite{1997ApJ...485..552M} suggested that  asymmetric morphologies and bars, especially in Seyfert 2 galaxies, are an important factor in  star formation and Seyfert classification. These asymmetric morphologies can induce radial gas inflow and fueling of the nuclear region, thus  obscuring and  feeding  the active nucleus. \\cite{2001ApJ...546..845G} discuss the possibility of two kinds of Seyfert 2 galaxies based on their stellar population properties: those with young and intermediate age stars and those with the optical continuum dominated by old stars.   In principle, the richness of the infrared spectrum provides a unique opportunity to test the unified model of AGN, since  the mid- and far-infrared spectra appear to be different in Seyfert 1 and Seyfert 2 galaxies \\citep[e.g.,][]{1988ApJ...328L..35S,1992ApJ...401...99P,1993ARA&A..31..473A,2000A&A...357..839C,2007ApJ...656..148A}. However, there is  the technical difficulty of  isolating the AGN  from  contamination by the host galaxy emission and, more importantly, star formation features \\citep[e.g.,][]{2004A&A...418..465L,2005ApJ...633..706W}. In this work we will focus on deconvolving the different contributions (e.g., AGN+star formation) in the [Ne~II]~$\\lambda$12.81~$\\micron$ emission line and the mid- and far-infrared continua. In order to estimate the  component associated with the AGN  we will use the high-ionization potential ($\\sim {\\rm 54.9~eV}$) [O~IV]~$\\lambda$25.89~$\\micron$ emission line. We   found \\citep[][hereafter M08]{2008arXiv0804.1147M} a tight correlation in Seyfert 1 galaxies between the [O~IV] and the X-ray 14-195~keV continuum luminosities from  {\\it Swift}/BAT observations \\citep{2005ApJ...633L..77M}. A weaker correlation was  found in Seyfert 2 galaxies, which was due to   partial absorption in the 14-195~keV band. Overall, we proposed [O~IV] as a truly isotropic property of AGNs given its high ionization potential and that is basically unaffected by reddening, meaning that the [O~IV] strength directly measures the AGN power. In this work we will  isolate the    stellar component of the [Ne~II] emission to trace the instantaneous star formation rates in our sample of Seyfert galaxies. ",
        "conclusions": "We have investigated the ionization state of the emission-line gas  in Seyfert  galaxies  with the aim of constraining the active galactic nuclei (AGN) and star formation contributions in the mid- and far-infrared spectra for a sample of 103 Seyfert galaxies. We found the ratio  between the AGN power and star formation, as shown  by [O~IV]/[Ne~II], to be smaller in Seyfert 2 than in Seyfert 1 galaxies. In this regard, we also found a  correlation between [Ne~III] and [O~IV] versus  [Ne~II], with a clear separation between Seyfert groups. This separation suggests  that the [Ne~II] emission has two different contributions:  a component that could be associated with  the AGN ionizing continuum and a component produced by photoionization from  star formation regions. The evidence for the former is   the presence of  [Ne~II] in Seyfert galaxies with no detectable PAH features at 6.2~$\\mu$m and 11.5~$\\mu$m. We also obtained a strong correlation between [Ne~III] and [O~IV] confirming that [Ne~III] can also be used to  estimate  the intrinsic power of the AGN.   We used the [O~IV] and [Ne~II] correlation  from the sources with no detectable PAH features as a template  to deconvolve the star formation and AGN contribution in the [Ne~II] emission. We found that Seyfert 2 galaxies, on average, have a relatively higher star formation contribution in their [Ne~II], by a factor of $\\sim 1.5$, than Seyfert 1 galaxies, in agreement with other observations and theoretical work on circumnuclear regions of AGN \\citep[e.g.,][]{1985MNRAS.213..841T,1989ApJ...342..735H,1997ApJ...485..552M,2001ApJ...546..845G,2007ApJ...671..124D}. Using the stellar [Ne~II] luminosity as a SFR estimator we found that Seyfert 2 galaxies have similar  star formation rates in their spectra,  $8 \\pm 2 M_\\sun{\\rm ~yr^{-1}}$ to that found for  Seyfert 1 galaxies, $7\\pm 2 M_\\sun{\\rm ~yr^{-1}}$. This result must be interpreted carefully given the fact that [Ne~II] emission provides an instantaneous measure of the star formation rate, independent of the previous star formation history. Thus, caution most be taken when comparing with other SFR indicators on galaxies with recent starbursts, but are no longer forming massive stars, in which case our star formation predictions may account only for a small fraction of a give SFR. Overall we found that $77\\%$ of the Seyfert 2 galaxies in our sample show  some star formation contribution to their [Ne~II] observed luminosities while $56\\%$ of the Seyfert 1 galaxies have some stellar component.   We  used the correlations between the mid- and far-infrared continuum  with the [O~IV]  luminosities in the non-PAH sources  to estimate the star formation contribution in the mid- and far-infrared continua. The resulting star formation contribution represents a lower limit  given the fact that some fraction of the star formation from the host galaxy is mixed with the AGN contribution in the mid- and far-infrared continuum in our sample of pure AGN sources. Averaged over all the sample we found the  contribution from star formation to the 25$\\micron$, 60$\\micron$, 100$\\micron$ and FIR luminosities to be: $32\\pm 2\\%$, $45\\pm 5\\%$, $39 \\pm 4\\%$ and $42\\pm 4\\%$, respectively. We also found that Seyfert 1 galaxies exhibit  a narrower range of star formation contribution, $\\sim 26 \\pm 1\\%$, to their mid- and far-infrared continuum luminosities, than that found for Seyfert 2 galaxies, $\\sim 47 \\pm 9 \\%$.    We also found a good correlation between [Ne~II] and mid- and far-infrared  luminosities, with the strongest  correlation for the [Ne~II]-60$\\micron$. This result is in agreement with previous studies that link both the FIR and [Ne~II] emission with star formation. We  found a weaker correlation between [O~IV] and [Ne~III] versus the FIR luminosities. This result suggest that part of the FIR luminosity is reprocessed photons from the AGN, most likely  through reradiation  of  the AGN continuum by dust.   Assuming that the infrared continuum is dominated by emission from dust grains,  we found in our sample, that Seyfert 2 galaxies possess   cooler dust, with an average $\\alpha_{25-60}=-1.5\\pm0.1$,  than Seyfert 1 galaxies which have $\\alpha_{25-60}=-0.8\\pm0.1$.  In agreement, we also found that Seyfert 1 and Seyfert 2 galaxies are statistically different in terms of their  relative contribution from the AGN to the 60$\\micron$ and FIR but have a similar contribution in the $25\\micron$.    Summarizing, we found that Seyfert 1 and Seyfert 2 galaxies are statistically different in terms of the relative contribution from  the AGN and star formation ,as showed by the [O~IV]/[Ne~II] ratios. From this we found that Seyfert 2 galaxies have a higher star formation contribution, relative to the strength of the AGN, than that found in  Seyfert 1 galaxies but have  similar star formation rates. These results suggest that the differences between Seyfert galaxies cannot be  solely due to  viewing angle dependence."
    },
    "0808/0808.3890_arXiv.txt": {
        "abstract": "{Nearby young clusters are informative places to study star formation history. Over the last decade, the $\\sigma-$Orionis cluster has been a prime location for the study of young very low mass stars, substellar and isolated planetary mass objects and the determination of the initial mass function.} {To extend previous studies of this association to its core, we searched for ultracool members and new multiple systems within the 1\\farcm5$\\times$1\\farcm5 central region of the cluster. } {We obtained deep multi-conjugate adaptive optics (MCAO) images of the core of the $\\sigma-$Orionis cluster with the prototype MCAO facility MAD at the VLT using the H and K$_{\\rm s}$ filters. These images allow us to detect companions fainter by $\\Delta$H$\\approx$5~mag as close as 0\\farcs2 on a typical source with H=14.5~mag. These images were complemented by archival SofI Ks-band images and \\emph{Spitzer} IRAC and MIPS mid-infrared images} {We report the detection of 2 new visual multiple systems, one being a candidate binary proplyd and the other one a low mass companion to the massive star $\\sigma$~Ori~E. Of the 36 sources detected in the images, 25 have a H-band luminosity lower than the expected planetary mass limit for members, and H-K$_{\\rm s}$ color consistent with the latest theoretical isochrones. Nine objects have additional \\emph{Spitzer} photometry and spectral energy distribution consistent with them being cluster members. One of them has a spectral energy distribution from H to 3.6~$\\mu$m consistent with that of a 5.5~M$\\rm_{Jup}$ cluster member. Complementary NTT/SofI and \\emph{Spitzer} photometry allow us to confirm the nature and membership of two L-dwarf planetary mass candidates.} {} ",
        "introduction": "Over the last decade, the $\\sigma-$Orionis cluster has become one of the prime locations for the study of brown dwarfs (BDs) and planetary-mass objects (PMOs). It is young (2--3~Myr), free of extinction (A$_{\\rm V}<$1~mag) and it has a large low-mass population \\citep{1999ApJ...521..671B,2001ApJ...556..830B,2003A&A...404..171B,2004Ap&SS.292..339B,2004AJ....128.2316S,2005MNRAS.356...89K,2007AN....328..917C,2007A&A...470..903C,2008MNRAS.383..375C,2006A&A...460..799G, 2008arXiv0805.2914S}. The {\\it Hipparcos} parallax to the central OB pair $\\sigma$~Ori~AB is 320$^{+120}_{-90}$~pc, and most previous studies used 350~pc to $\\sigma$~Ori~AB as the cluster distance. Other authors estimated  distances at about 390~pc \\citep{2008MNRAS.383..750C,2008MNRAS.386..261M}. \\citet{2008AJ....135.1616S} recently refined the measurement using main-sequence fitting and derived an improved value of 440~pc for a solar metallicity \\citep{2006PhDT........Caballero}. The cluster mass function rises steadily from the very low mass stars through the BDs and into the PMO domain \\citep{2001ApJ...556..830B,2007A&A...470..903C}. No obvious discontinuities are seen either at the mass boundary between very low-mass stars and BDs or at the frontier between BDs and PMOs, which is not surprising as the initial mass function does not reflect the onset of nuclear reactions that may have occured during the subsequent evolution for stars and BDs. About two dozens PMOs candidates have been identified in the cluster \\citep{2002ApJ...578..536Z,2000Sci...290..103Z,2007A&A...470..903C}. About half of those have been confirmed spectroscopically \\citep{2001ApJ...558L.117M,2001A&A...377L...9B,2003ApJ...593L.113M,2000Sci...290..103Z}.  Evidence of disks has been found around BD and PMO members by the detection of large H$\\alpha$ emission, flux excesses at mid-infrared (mid-IR) wavelengths and large-amplitude photometric variability \\citep{2003AJ....126.2997B, 2003A&A...404..171B,2003ApJ...592..266M,2006A&A...445..143C, 2004A&A...419..249S, 2008ApJ...672L..49S,2007ApJ...662.1067H}. In particular the very strong H$\\alpha$ emission brown dwarf SOri~71, located just above the cluster deuterium burning limit, displays excess flux in IRAC band 4.5 at 8.0~$\\mu$m \\citep{2007A&A...470..903C}. If the cluster PMOs form in a similar way to the very low-mass stars and BDs, we expect that they would have dusty disks, probably of lower mass scaling with primary mass. \\citet{2007ApJ...662.1067H} suggested an increase of the disk frequency  towards low masses in the cluster, with a peak of 40\\% in the mass  interval 0.2--0.1~M$_{\\sun}$. According to \\citet{2007A&A...470..903C}, the disk  rate in the BD domain could be as high as 50\\%. \\citet{2007A&A...472L...9Z} and \\citet{2008ApJ...672L..49S} have reported mid-IR excesses in seven cluster PMOs, indicating that more than  30\\% of them have dusty disks. \\citet{2008A&A...481..423G} reported that the mass accretion rate of $\\sigma-$Orionis members harboring disks is significantly lower than that of e.g $\\rho-$Oph members. On the other hand, \\citet{2008arXiv0805.2914S} report a significantly larger fraction of accretors than in the neighboring $\\lambda-$Orionis association where star formation might have been triggered by a supernova explosion \\citep{2001AJ....121.2124D}.  The multiplicity of $\\sigma-$Orionis members has been the subject of a number of studies. Using a Fraunhofer micrometer, \\citet{1837AN.....14..249S} resolved $\\sigma$~Ori~AB for the first time as a close binary (0\\farcs26). A few decades later, \\citet{1893AN....131..329B} used the micrometer mounted on the 36 inch telescope at the Lick Observatory and confirmed the multiplicity of $\\sigma$~Ori~AB. \\citet{1904ApJ....19..151F}  identified a possible spectroscopic component in the $\\sigma$~Ori~AB binary system. More recently \\citet{2005AN....326.1007C} obtained adaptive optics (AO) images of the central region of the cluster, resolving a number of sources. \\citet{2005MNRAS.356...89K}  used high resolution spectroscopic measurements of a sample of candidate very low-mass stars and BDs of the association to confirm the youth and membership and to search for spectroscopic binaries. Their preliminary conclusions, based on small-number statistics, shows that the binary fraction among $\\sigma-$Orionis very low mass members is higher than that reported for their older field counterparts. Our understanding of the $\\sigma-$Orionis cluster was recently complicated by the discovery by \\citet{2006MNRAS.371L...6J} and \\citet{2007A&A...462L..23S} of two distinct kinematics populations with different ages. However, most of the objects belonging to the older kinematic group are located northward of the central region of the $\\sigma-$Orionis cluster, outside the area covered by the present study.  To extend the previous studies of $\\sigma-$Orionis to its core we conducted an AO assisted imaging survey of the central part of the $\\sigma-$Orionis cluster with the ESO multi-conjugated AO prototype instrument \\emph{Multi-Conjugate Adaptive Optics Demonstrator} \\citep[hereafter MAD][]{2006SPIE.6272E..21M}. These deep images have a resolution of $\\approx$0\\farcs1 on a field of view of 1\\farcm5$\\times$1\\farcm5. ",
        "conclusions": "Using multi-conjugate AO images of the core of the $\\sigma-$Orionis cluster, we have identified 6 new BD candidates and 25 planetary mass candidates. Five of these have additional mid-IR {\\it Spitzer} IRAC photometry consistent with that of sub-stellar members. With the current data, it is not possible to conclude on their membership of the association as they could also be foreground late-M dwarfs. The candidate proplyd $\\sigma$~Ori~IRS1 is resolved as a binary. The high spatial resolution MAD images resolve 5 pairs, including 2 previously unknown ones. The results presented in this paper illustrate the capacity of multi-conjugate AO to probe the immediate vicinity of young massive OB stars in great detail. Using complementary archival SofI and {\\it Spitzer} images, we were able to confirm the membership of two L-dwarf candidate members that exhibit near- and mid-IR excess associated with a disk. One planetary mass candidate newly detected in the MAD images (SigOri-MAD-31) displays a SED very similar to these latter two, suggesting that it is also a L-dwarf member of the association harboring a disk. The presence or absence of very low mass stars, BDs and planetary mass objects close to massive stars provides novel constraints on the models of formation. Follow-up observations of the candidates are very much needed to confirm the nature and membership of the new candidates, and to provide quantitative feedback on the models of formation."
    },
    "0808/0808.3224_arXiv.txt": {
        "abstract": "{We discuss the current knowledge of the Solar system, focusing on bodies in the outer regions, on the information they provide concerning Solar system formation, and on the possible relationships that may exist between our system and the debris disks of other stars.  Beyond the domains of the Terrestrial and giant planets, the comets in the Kuiper belt and the Oort cloud preserve some of our most pristine materials.  The Kuiper belt, in particular, is a collisional dust source and a scientific bridge to the dusty ``debris disks'' observed around many nearby main-sequence stars.  Study of the Solar system provides a level of detail that we cannot discern in the distant disks while observations of the disks may help to set the Solar system in proper context. } ",
        "introduction": "Planetary astronomy is unusual among the astronomical sciences in that the objects of its attention are inexorably transformed by intensive study into the targets of other sciences.   For example, the Moon and Mars were studied telescopically by astronomers for nearly four centuries but, in the last few decades, these worlds have been transformed into the playgrounds of geologists, geophysicists, meteorologists, biologists and others.  Telescopic studies continue to be of value, but we now learn most about the Moon and Mars from in-situ investigations.  This transformation from science at-a-distance to science up-close is a forward step and a tremendous luxury not afforded to the rest of astronomy.  Those who study other stars or the galaxies beyond our own will always be forced by distance to do so telescopically.  However, the impression that the Solar system is now \\textit{only} geology or meteorology or some other science beyond the realm of astronomy is completely incorrect when applied to the outer regions.  The Outer Solar system (OSS) remains firmly entrenched within the domain of astronomy, its contents accessible only to telescopes. Indeed, major components of the OSS, notably the Kuiper belt, were discovered (telescopically) less than two decades ago and will continue to be best studied via. astronomical techniques for the foreseeable future. It is reasonable to expect that the next generation of telescopes in space and on the ground will play a major role in improving our understanding of the Solar system, its origin and its similarity to related systems around other stars.  In this chapter, we present an up-to-date overview of the layout of the Solar system and direct attention to the outer regions where our understanding is the least secure but the potential for scientific advance is the greatest.   The architecture of the Kuiper belt is discussed as an example of a source-body system that probably lies behind many of the debris disks of other stars. Next, the debris disks are discussed based on the latest observations from the ground and from Spitzer, and on new models of dust transport.  Throughout, we use text within grey boxes to highlight areas where the Next Generation telescopes are expected to have major impact. ",
        "conclusions": ""
    },
    "0808/0808.1367_arXiv.txt": {
        "abstract": "We quantify an evolutionary channel for single sdB stars based on mergers of binaries containing a red giant star and a lower mass main sequence or brown dwarf companion in our Galaxy.  Population synthesis calculations that follow mergers during the common envelope phase of evolution of such systems reveal a population of rapidly rotating horizontal branch stars with a distribution of core masses between $0.32\\Msun - 0.7\\Msun$ that is strongly peaked between $0.47\\Msun - 0.54 \\Msun$. The high rotation rates in these stars are a natural consequence of the orbital angular momentum deposition during the merger and the subsequent stellar contraction of the merged object from the tip of the red giant branch.  We suggest that centrifugally enhanced mass loss facilitated by the rapid rotation of these stars may lead to the formation of single sdB stars for some of these objects. ",
        "introduction": "Subdwarf B (sdB) stars are core He-burning stars that have an extremely thin ($\\lesssim$ 0.02$\\Msun$) hydrogen-rich envelope (for recent reviews, see \\citealp{heb08,cat07,mon07}).  Observations indicate that 40 -- 70\\% of sdB stars in the field are in binaries \\citep{max01,ree04,nap04,lis05,mor06}.  \\citet{han02,han03} identified and performed extensive population synthesis calculations of two channels for the formation of sdB stars in binaries in the field: (1) one or two phases of unstable Roche lobe (RL) overflow leading to common envelope (CE) evolution and (2) one or two phases of stable RL overflow, and one channel for the formation of single sdB stars in the field:  the merger of two He white dwarfs (WDs).  Their model of the population of sdB stars in binaries has been very successful in explaining several features of the observed field population, such as the prevalence of systems with $P \\lesssim$ 10 days, and this model is generally well accepted.   Their model of the single sdB star population in the field predicts a relatively flat distribution of masses ranging from $\\sim 0.4\\Msun - 0.7\\Msun$ (see Fig.\\,12 in \\citealp{han03}).   This distribution differs from the mass distribution expected from single star horizontal branch (HB) models (e.g., \\citealp{dor93}), which predict a narrow distribution of sdB star masses strongly peaked at $\\sim0.47\\,\\Msun$.  Observationally, masses for seven single field sdB stars have been determined using asteroseismology (see \\citealp{ran07} and references therein).  Six are very close to 0.47$\\Msun$ and one is 0.39$\\Msun$.   However, it is not evident whether these masses are representative of the single sdB star population as a whole, since only $\\sim$ 5\\% of sdB stars pulsate.  Further, reliable mass determinations using asteroseismology depend on correctly identifying the pulsational modes of the star, which can be quite difficult.  The formation of single sdB stars also has been modeled using channels that do not require the merger of two WDs, but rather involve a single star with a phase of enhanced mass loss on the red giant branch (RGB).  Various suggestions for the cause of this enhanced mass loss include He mixing driven by internal rotation (Sweigart 1997), spin up of the primary via interactions with a planet in a close orbit \\citep{sok98}, and the existence of a sub-population of ZAMS stars with a high He-abundance (e.g., D'Antona et al. 2005).   In all of these proposed causes, the amount of mass loss on the RGB must be fine-tuned in order to yield the very small envelope masses observed in sdB stars.  However, it is not evident how such fine-tuning is incorporated naturally into the physical mechanisms that have been suggested to enhance the mass loss.  In this Letter, we quantify a channel for the formation of single sdB stars in the field of our Galaxy: the evolution of objects resulting from the merger of stellar or sub-stellar objects with RGB stars during a phase of CE evolution initiated by unstable RL overflow or by a tidal instability (hereafter, referred to as the \"CE merger channel\").  We carry out detailed population synthesis calculations of these mergers and predict the resulting population of HB stars at the present epoch, some of which may have very small envelope masses and could be observed as single sdB stars.  Although our treatment was developed independently, the idea that single sdB stars might result from the merger of an RGB star and a low mass companion during CE evolution was suggested by \\citet{sok98} and explored further by \\citet{sok00,sok07}. ",
        "conclusions": "In our model population of HB stars, a small fraction ($\\la 0.1$\\%) have envelopes with masses less than 0.1$\\Msun$ (see Fig.\\,1).  Approximately one-half of these systems have rotational velocities equal to $v_\\mathrm{crit}$, suggesting that these stars will continue to experience enhanced mass loss.  Follow-up evolutionary calculations of core-He burning stars with an initial envelope mass of 0.1$\\Msun$ (using the same {\\it STARS} stellar evolution code with an artificially large wind to simulate rapid mass loss) suggest that at some point ($\\sim0.07\\Msun$) the star will begin to contract upon further mass loss and move blueward along the HB.  It is conceivable, therefore, that these critically-rotating, very low envelope mass HB stars in our population may lose sufficient envelope material to become sdB stars.  These potential sdB systems constitute a very small fraction of our population.  Upon examination of Fig.\\,1, this fraction would increase substantially if the envelope masses in our population were systematically reduced by $\\sim0.2$$\\Msun$.  Such a reduction might be accomplished by assuming a higher rate of mass loss due to stellar winds on the RGB.   We have adopted a standard Reimers-like prescription for mass loss in our stellar models, but it has been suggested that a rapidly rotating envelope may lead to enhanced mass loss on the RGB (e.g., \\citealp{sw97}).  In addition, we find that our assumption of synchronization on the RGB results in small number of primaries being spun up to $v_\\mathrm{crit}$ prior to contacting their RL.  This number is increased significantly if we choose $v_\\mathrm{crit}<\\onethird\\,v_\\mathrm{br}$.  For example, if $v_\\mathrm{crit}=0.1\\,v_\\mathrm{br}$, two-thirds of the primaries are spun up to $v_\\mathrm{crit}$ prior to contacting their RL.  However, we are not able to properly model either possibility in our population synthesis calculations, since a self-consistent treatment of mass loss in such cases would require stellar evolutionary models in which mass loss via winds is coupled to the rotation and this is beyond the current capabilities of our code.  Reliable quantitative predictions concerning the absolute number of single sdB stars produced from our CE merger channel are not provided here, since our treatment of the merger process is too crude and is severely limited by the lack of models of stellar mergers relevant to CE evolution.  Rather, our intent in this Letter has been to investigate a possible evolutionary channel for the formation of single sdB stars and to motivate its further study.  In comparison with previously suggested single sdB star channels, this channel has certain advantages:   (1) the AM necessary to spin up the envelope prior to the HB is explicitly and naturally provided by the orbital AM of the companion (see also \\citealp{sok98}); (2) the CE phase may be initiated at any point along the RGB, not just near the tip, since mass loss from the envelope also occurs on the HB as a result of super-critical rotation caused by the star's contraction following He ignition in the core; and (3) the predicted spectrum of core masses is consistent with that expected in HB stars.  Further, our population synthesis calculations are, to the best of our knowledge, the first to explore in any substantive manner a possible observational counterpart to mergers resulting from CE evolution.  It is generally accepted that mergers are likely to occur during CE evolution (e.g., \\citealp{ibe93,taa00}) and with a frequency roughly comparable to that of the surviving post-CE binaries (e.g., \\citealp{pw07}).  Thus, there may exist a significant, but as yet unrecognized, population of objects in our Galaxy that are the result of a merger during CE evolution.  Assuming this is true, identification of possible observational counterparts of these systems is likely to stimulate the development of models necessary to study them."
    },
    "0808/0808.3362_arXiv.txt": {
        "abstract": "A significant population of stars with ages younger than the Pleiades exists in the solar neighborhood. They are grouped in loose young associations, sharing similar kinematical and physical properties, but, due to their vicinity to the Sun, they are dispersed in the sky, and hard to identify.  Their strong stellar coronal activity, causing enhanced X-ray emission, allows them to be identified as counterparts of X-ray sources.  The analysis presented here is based mainly on the SACY project, aimed to survey in a systematic way counterparts of ROSAT all-sky X-ray sources in the Southern Hemisphere for which proper motions are known.  We give the definition, main properties, and lists of high-probability members of nine confirmed loose young associations that do not belong directly to the well-known Oph-Sco-Cen complex.  The youth and vicinity of many members of these new associations make them ideal targets for follow-up studies, specifically geared towards the understanding of planetary system formation.  Searches for very low-mass and brown dwarf companions are ongoing, and it will be promising to search for planetary companions with next generation instruments. ",
        "introduction": "\\subsection{Overview} In a seminal paper, \\cite{Herbig1978} proposed the existence of post-T~Tauri stars (pTTS) and established a list of candidates. pTTS follow classical and weak line T~Tauri stars (cTTS, wTTS) in an evolutionary sequence. Although they were expected to outnumber the TTS, their discovery or identification remained difficult for a long time. Their relation with another category of young stars -- the isolated TTS (young low mass stars found spatially far away from any apparent dark or parental molecular cloud) -- remained an open problem. Interestingly, Herbig's list of pTTS and the one of isolated TTS \\citep{Quast1987} were actually quite similar. The most intriguing example was TW~Hya, a high Galactic latitude TTS \\citep{RucinskiKrautter1983} at a distance of at least 13$^\\circ$ from the nearest  dark clouds and located in a region characterized by the absence of any cloudlets from which it could have originated.  A systematic search for more isolated TTS was pursued with the  optical spectroscopic Pico dos Dias Survey (PDS) among  optical counterparts of the  IRAS Point Source Catalog \\citep{gregorio92, Torres1995, Torres1999}. One of the first results of the PDS was the discovery of four additional TTS around TW~Hya  \\citep{delareza89, gregorio92}. They concluded that this group was likely a very young association relatively close to the sun.   Only very few good candidates for isolated TTS were found within the PDS. But while the IR-excess selection criterion  effectively finds young stellar objects embedded in their placental material or with circumstellar disks, it fails to signal older objects whose disks have already been dissipated. Therefore, most stars  with ages between $\\ga10-70$~Myr (i.e. wTTS and pTTS) escaped discovery by this method.  Due to the enhanced X-ray activity of young stars \\citep{Walter86}, more efficient selection criteria for post and isolated TTS candidates were developed. The high sensitivity and full sky coverage of the ROSAT all-sky survey \\citep{truemper} revealed thousands of new X-ray sources projected in the direction of nearby star forming regions \\citep{Guillout98}. Ground-based spectroscopic follow-up studies showed that a large fraction of these X-ray sources were indeed wTTS  together with older pTTS and ZAMS stars \\citep[e.g.,][]{Alcala00}. Surprisingly, many of the newly found weak-line TTS were {\\it not} obviously connected to any molecular cloud region, raising again many questions about their origin \\citep{Sterzik95}.   Based on the similarity of the ROSAT X-ray fluxes, radial velocities, astrometry (Hipparcos) and spectroscopic characteristics of the four stars around TW~Hya, \\cite{Kastner97} confirmed that these stars formed a physical association with TW~Hya, about 20~Myr old and at a distance of 40 to 60 pc from Earth, which they called the TW~Hya Association.  Immediately after the existence of the TW~Hya Association was confirmed, several research groups became interested in this association and other members were found (see Section~4). Two main approaches were explored, a first one aiming to study the properties of its members, whereas other groups started to look for similar nearby young associations hidden among the ROSAT X-ray sources.   In 2000, as a result of this effort to find new associations, two new adjacent and similar associations were proposed, in Tucana \\citep{zuckerman00} and in Horologium \\citep{torres00}. To examine the physical relation between them and to search for other associations, we started the SACY (Search for Associations Containing Young stars) survey \\citep{torres06}. The SACY sample contains stars:\\\\ {\\it (i)} later than G0, in order to be able to use the Li $\\lambda$6707 line as a youth indicator \\citep{martin97};\\\\ {\\it (ii)} belonging to the TYCHO-2 or Hipparcos catalogs in order to have access to proper motions;\\\\ {\\it (iii)} that are candidate  optical counterparts of sources from the ROSAT All-Sky Bright Source Catalogue.  \\citet{torres06} presented a catalog of 1626 spectroscopically observed stars in the Southern Hemisphere.  In Figure{~\\ref{fig:thpiz}} we show the celestial distribution of the observed SACY sample, with additional stars observed from 2006 to 2008.  The SACY survey is now complete in the Southern Hemisphere (but for four stars) and the updated catalog has 2093 stars.  The survey enables us to define properties and membership probabilities for stars in these putative associations.  Preliminary results appear in \\citet{torres03a, torres03}.  In \\citet{torres06}, the prototypical methodology and analysis of the $\\beta$~Pic Association is presented and in forthcoming papers (in preparation) a more detailed analysis of other associations found will be given.  The Lithium abundances of the nine associations presented in this chapter are studied in \\citet{silva08} where they present the lists of the association members.  A similar survey is being pursued in the Northern Hemisphere and the first results are appearing now \\citep{guillout08}. Only a very low frequency of stars younger than the Pleiades is identified, more than one order of magnitude less than in the SACY survey. This strong hemispheric anisotropy could be explained, at least partially, by taking into account the different survey biases and completeness limits. Anyway, with only five young stars at this moment, the Northern survey does not help to find young associations.   \\subsection{Method of Analysis}  An association is a group of stars appearing {\\it concentrated} together in a small volume in space sharing some common properties such as age, chemical composition, distance and kinematics\\footnote{ In this sense we prefer to use the term association and not moving group.}.  However, if such a group is close enough to the Sun, its members will appear to cover a large extent in the sky (as an example, Orion at 50~pc would cover almost the whole sky).  Thus, to find a group, projected spatial concentrations (i.e., in terms of right ascension and declination only) and proper motions may not be enough. A better criterion is to look for objects sharing similar heliocentric space motions (UVW) all around the sky (U positive towards the Galactic center, V positive in the direction of Galactic rotation).  \\citet{torres06} describe the convergence method developed to search for members of an association and a corresponding membership probability model in detail. Both convergence method and probability model examine the stars in the hexa-dimensional space, UVWXYZ, as defined by the space motions relative to the Sun  and the physical space coordinates centered on the Sun (XYZ, in the same directions as UVW). We represent with m$_v$, M$_v$ and M$_{v,iso}$ the apparent visual magnitude, the resultant absolute magnitude with the distance obtained from the convergence method,  and the absolute magnitude given by  the adopted isochrone for the $(V-I)_C$ stellar color; and   $\\mu_\\alpha$, $\\mu_\\delta$ and V$_r$ are the proper motions and the radial velocity. Briefly explained, if there is no reliable trigonometric distance available\\footnote{We consider the trigonometric parallaxes as unreliable if they have errors larger than 2~mas and we do not use them.}, the convergence method finds the distance (d) for each star in the sample that minimizes the F value of Equation~\\ref{eq:f}. The first term is a photometric distance modulus and the second is a kinematical one. The method needs, as input, an assumed age and initial velocity values $(U_0, V_0, W_0)$ for the proposed association, and a cutoff value for F above which stars should be considered spurious. This cutoff value varies for each association but usually we begin with 3.5 (this approximately means 0.7 magnitudes for the distance modulus and 3~km~s$^{-1}$ for the velocity modulus). The method is iterative, and for each iteration a list of stars with new  $(U_0, V_0, W_0)$ is obtained. The process ends when the list of stars and the velocities $(U_0, V_0, W_0)$ do not change significantly. \\begin{eqnarray} \\label{eq:f} \\nonumber \\lefteqn{F(m_v,\\mu_\\alpha,\\mu_\\delta,V_r;d) =}\\\\ &[p\\times(M_v-M_{v,iso})^2+(U-U_0)^2+(V-V_0)^2+(W-W_0)^2]^{1/2} \\end{eqnarray} where $p$ is a constant weighting the importance of the evolutionary distance with respect to the kinematic distance. Actually we use $p>0$ only for the Oct Association, since it has no stars with trigonometric parallax. This means that in general our distances are only kinematical.  The  list of candidates serves as a training set for the probability model \\cite[k--NN model,][]{Sterzik95}. In this model we define around each star of the entire sample 6-dimensional spheres that contain a certain number $k$ of stars. A membership probability is then defined by the proportion of stars in these spheres that belong to the training set. The probabilities depend on the compactness of the association and the field density. Thus, for each association we can define a cutoff probability where we consider the stars as  probable members. Using this list of (high) probability members, we return to the convergence method until both lists match. Finally, a possible kinematical member becomes an actual good candidate if its Li content is compatible with the Li depletion for its age \\citep{neu97}.  As all the nearby loose associations proposed at this moment can be defined through their properties using SACY stars,  membership probabilities for  stars suggested elsewhere can be calculated whenever their basic kinematic data are available.  \\begin{figure}[!ht] \\centering \\includegraphics[draft=False,width=0.49\\textwidth]{EXP1.EPS} \\includegraphics[draft=False,width=0.49\\textwidth]{EXP.EPS} \\caption{{\\it Left:} The expansion in X direction for the young nearby associations, including the Sco-Cen and the R CrA (CrA) associations using the SACY data. We could not find expansion for the $\\epsilon$~Cha and the Carina associations, as they are very compact, and any expansion of the Oct Association should be confirmed -- see Section~6. The other abbreviations in the figure mean the associations: Columba, Tucana-Horologium, TW~Hya, $\\beta$~Pic, $\\epsilon$~Cha and Octans. {\\it Right:} The expansion of all young stars in the SACY sample with reliable parallaxes. Open circles are stars which are not members of any association.} \\label{fig:exp} \\end{figure}     When analyzing the entire SACY sample, we discovered an unexpected kinematical phenomenon that seems to be a general property of the young stars in the solar neighborhood and affects the application of the convergence method -- a positive correlation (r\\,=\\,0.99) between  U and X for stars younger than $\\sim$30~Myr. The correlation is seen in Figure~\\ref{fig:exp} where, in the left panel, we plot the mean values of U and X  derived from the SACY sample for these associations. This correlation can be interpreted as a physical expansion in the X direction.  This is not an artifact from the convergence method. The right panel of Figure~\\ref{fig:exp} shows the observed U and X values for all young SACY stars that have good quality Hipparcos parallaxes. While a strong correlation (r\\,=\\,0.98) of U and X persists for those stars that belong to young associations (filled circles), the correlation is only weak (r\\,=\\,0.46) for stars that do not belong to any association. Maybe one reason for the difference is the presence of not identified single lined spectroscopic binaries. The expansion is not only seen in the sample of nearby young associations as a whole, but also within each of these associations, with a similar spread of U and X values \\citep{torres06}. Some young associations have only a small extension in the X axis, and do not allow to determine their expansion ($\\epsilon$ Cha and Car associations). The Oct Association does not seem to follow the behavior shown in Figure~\\ref{fig:exp}, perhaps due to its higher Galactic latitude (see Table~\\ref{table:finalb}). For associations older than $\\sim$30~Myr this expansion is not present any more (AB Dor and Argus associations; see Table~\\ref{table:final}).  A similar phenomenon has been reported for a few individual associations, see for example \\citet{mamajek05} and \\citet{bobylev07}, but the global expansion must be considered first.  Unfortunately we have no explanation for the presence of this expansion.  It might reflect a more global motion like the Galactic arm epicycle movement which is still conserved from a recent local star formation event, and which has not yet been lost by higher dispersion acquired through Galactic dynamics on a longer timescale.  Regardless of the cause of this expansion, it must be taken into account when searching for kinematical young associations. It can be represented by the relation of  Equation~\\ref{eq:exp} that is incorporated in the convergence method in the appropriate cases:  \\begin{equation} \\label{eq:exp} U = 0.05 (X) - U_0 \\end{equation}   Reliable isochrones are essential as input to obtain the photometric distance modulus. However, none of the observed sequences can be represented by the published theoretical isochrones for these associations. Thus, to test the star membership, ad-hoc observational evolutionary sequences were used in the convergence method, represented by third degree polynomials:\\\\ \\\\ For 5~Myr: \\begin{equation} \\label{eq:iso5} M_v = 0.60 + 4.98(V-I)_C - 1.16(V-I)_C^2 +0.193(V-I)_C^3 \\end{equation} For 8~Myr: \\begin{equation} \\label{eq:iso8} M_v = 1.20 + 4.98(V-I)_C - 1.16(V-I)_C^2 +0.193(V-I)_C^3 \\end{equation} For 10~Myr: \\begin{equation} \\label{eq:iso10} M_v = 1.50 + 4.98(V-I)_C - 1.16(V-I)_C^2 + 0.193(V-I)_C^3 \\end{equation} For 30~Myr: \\begin{equation} \\label{eq:iso30} M_v = 1.18 + 6.28(V-I)_C - 1.68(V-I)_C^2 + 0.248(V-I)_C^3 \\end{equation} For 70~Myr: \\begin{equation} \\label{eq:iso70} M_v = 0.64 + 7.14(V-I)_C - 2.05(V-I)_C^2 + 0.314(V-I)_C^3 \\end{equation} valid in the color interval $-0.1 < (V-I)_C < 3.1$ \\\\ \\\\ These heuristic {\\sl absolute} ages obtained partially in comparison  with pre-main sequence models \\citep{siess00} must obviously be taken with caution. Nevertheless, the {\\sl relative} ages  are real: for example, the TW~Hya Association is older than the  $\\epsilon$~Cha Association and younger than the $\\beta$~Pic Association.   \\subsection{General Results}  Almost half of the young stars in the SACY sample belong to the large Oph-Sco-Cen Association (see the chapters by Wilking et al. and Preibisch \\& Mamajek). In this chapter we discuss  nine nearby young loose associations, which are kinematically well defined, but do not belong directly to the Oph-Sco-Cen Complex. Their main characteristics are summarized in Tables~\\ref{table:final} and \\ref{table:finalb}. Their distribution in the sky  can be seen in  polar projection in  Figures{~\\ref{fig:thpiz} to ~\\ref{fig:abpiz}}.  \\begin{table}[!ht] \\caption{Heliocentric space motions and expansion of the nearby associations} \\smallskip \\begin{center} { \\label{table:final} \\begin{tabular}{lrrrcr } \\tableline \\noalign{\\smallskip} Association & U~~~~~~~& V~~~~~~~& W~~~~~~~& Expansion& N\\\\ &[km s$^{-1}]$&[km s$^{-1}]$&[km s$^{-1}]$&&\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} $\\beta$~Pic   &-10.1$\\pm$2.1&-15.9$\\pm$0.8& -9.2$\\pm$1.0&yes&48\\\\ Tuc-Hor&       -9.9$\\pm$1.5&-20.9$\\pm$0.8& -1.4$\\pm$0.9&yes&44\\\\ Col    &       -13.2$\\pm$1.3&-21.8$\\pm$0.8& -5.9$\\pm$1.2&yes&41\\\\ Car    &       -10.2$\\pm$0.4&-23.0$\\pm$0.8& -4.4$\\pm$1.5&no?&23\\\\ TW~Hya           & -10.5$\\pm$0.9&-18.0$\\pm$1.5& -4.9$\\pm$0.9&yes&22\\\\ $\\epsilon$~Cha           & -11.0$\\pm$1.2&-19.9$\\pm$1.2&-10.4$\\pm$1.6&no?&24\\\\ Oct          &  -14.5$\\pm$0.9& -3.6$\\pm$1.6& -11.2$\\pm$1.4&no?&15\\\\ Argus          & -22.0$\\pm$0.3&-14.4$\\pm$1.3& -5.0$\\pm$1.3&no&64\\\\ AB~Dor&        -6.8$\\pm$1.3&-27.2$\\pm$1.2&-13.3$\\pm$1.6&no&89\\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} } \\end{center} \\end{table}  \\begin{table}[!h] \\caption{Space distribution, mean distances and ages of the nearby associations} \\smallskip \\begin{center} { \\label{table:finalb} \\begin{tabular}{l@{\\hskip8pt}r@{\\hskip8pt}r@{\\hskip8pt}r@{\\hskip8pt}r@{\\hskip8pt}r@{\\hskip8pt}r@{\\hskip8pt}r@{\\hskip8pt}r} \\tableline \\noalign{\\smallskip}   Assoc. & X~~~& % X Range& % Y~~~& Y Range& % Z~~~& Z Range& D~~~~~& % Age \\\\ &[pc]&[pc]~~~~&[pc]&[pc]~~~~&[pc]&[pc]~~~~&[pc]~~~~&[Myr] \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} $\\beta$~Pic&20&  -32/76  &-5   &-33/21   &-15 &-29/-1&   31$\\pm$21  &10\\\\ Tuc-Hor&  3 &    -61/43  &-24  &-47/-4   &-35 &-44/-30&  48$\\pm$7~~~&30\\\\ Col    &-42 &    -106/9  &-56  &-168/1   &-47 &-99/6& 82$\\pm$30  &30\\\\ Car    & 14 &     -2/33  &-94  &-154/-39 &-17 &-33/5&    85$\\pm$35  &30\\\\ TW~Hya & 15 &      2/34  &-44  &-61/-26  &21  &10/27&    48$\\pm$13  & 8\\\\ $\\epsilon$~Cha&50& 34/60 &-92  &-105/-78 &-28 &-44/-12& 108$\\pm$9~~~& 6\\\\ Oct&     22 &     -79/142&-106 &-138/-60 &-68 &-85/-38& 141$\\pm$34  &20?\\\\ Argus&    5 &     -55/64 &-115 &-154/-6  &-18 &-67/8&  106$\\pm$51  &40\\\\ AB~Dor&  -6 &     -94/73 &-14  &-131/58  &-20 &-66/23&   34$\\pm$26  &70\\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} } \\end{center} \\end{table}    \\begin{sidewaysfigure} \\label{fig:sacyxyz} \\centering \\includegraphics[width=0.28\\textwidth]{SACYUV.EPS} \\includegraphics[width=0.28\\textwidth]{SACYUW.EPS} \\includegraphics[width=0.28\\textwidth]{SACYVW.EPS}\\\\ \\includegraphics[width=0.28\\textwidth]{SACYXY.EPS} \\includegraphics[width=0.28\\textwidth]{SACYXZ.EPS} \\includegraphics[width=0.28\\textwidth]{SACYYZ.EPS}\\\\ \\caption{Combinations of the sub-spaces of the UVWXYZ--space for the SACY Associations showing the clusters in both kinematical and spatial coordinates. Associations symbols: filled stars - $\\beta$~Pic; filled circles - Tuc-Hor; open circles - Columba; crosses - Carina; filled triangles - TW~Hya; open triangles - $\\epsilon$~Cha; open stars - Octans; open squares - Argus; filled squares - AB~Dor.} \\end{sidewaysfigure}   There are interpretations connecting the local young associations with some more global stellar populations, like the Oph-Sco-Cen Complex.  A particularly interesting paper is the one by \\citet{fernandez08}, in which they propose that both the Sco-Cen Complex and the young local associations originated by the impact of the spiral shock wave against a giant molecular cloud.  To clarify the questions opened by these kinds of models, it is fundamental to have a precise definition of each association and a very good way to define their star memberships and, also, their age.  A  distinct and independent way to obtain the age and the origin of  nearby young stars is proposed by  \\citet{makarov07}. He traces back in time the Galactic orbits and evaluates near approaches in order to infer close conjunctions and clustering in the past of the stars. He concludes that the majority of nearby young stars were formed during close passages or encounters of their natal clouds with other cloud complexes today located at somewhat larger distances. The method requires excellent kinematical data and depends on the true membership of the proposed star for each association. The ages of the Tuc-Hor and TW~Hya associations agree reasonably well for both approaches.  These young associations have remarkably small velocity dispersions ($\\approx$1~km per sec, see Table~\\ref{table:final}). Their sizes are larger than implied from their dispersion velocities and their ages, but they are fully consistent with low-mass star forming regions and OB associations.  Most of these associations have a non-spherical distribution, and show distortions that seem age-dependent: the younger ones are almost spherical ($\\epsilon$~Cha and TW~Hya associations), the ones with intermediate age are extended in X direction ($\\beta$~Pic, Oct, Tuc-Hor associations) while the older ones in Y direction (Car, Col, Argus and AB~Dor associations).  If we approximate them with spheres, they will be within a radius of 25~pc for the $\\epsilon$~Cha and the TW~Hya associations, 40~pc for the Tuc-Hor Association, 60~pc for the $\\beta$~Pic, the Car and the Col associations, and 100~pc for the Oct, the Argus, and the AB~Dor associations.  The important review on nearby young stars and their properties by \\citet{zuckerman04} presents five associations known at that time ($\\beta$~Pic, Tuc-Hor, TW~Hya,\\linebreak $\\eta$~Cha and AB~Dor associations) and introduces a new one, \"Cha-Near\" \\footnote{It is very similar to the $\\epsilon$~Cha Association (that includes the $\\eta$~Cha cluster) defined here, see Section 5.}.  The focus of the present review is the re-definition of their nearby associations and the definition of new ones with the method described above and the significantly increased sample. We emphasize that the strength of our method is the availability of high quality radial velocities and proper motions together with spectral information, which is essential to find the associations presented here.  For that the SACY sample is of great importance to better define these associations.  Nevertheless it is of almost no help for the $\\epsilon$~Cha association, which is at the limit of the SACY sample possibilities, with most of the candidates coming from the literature.  Therefore, in the next sections, we will present the high probability members of these nine associations and their main characteristics, and, when pertinent, we compare our new definitions with those of \\citet{zuckerman04}.  \\begin{figure}[] \\begin{center} \\begin{tabular}{c} \\resizebox{1.0\\hsize}{!}{ \\includegraphics[draft=False]{SACY.EPS} \\includegraphics[draft=False]{GAYPIZ.EPS} } \\end{tabular} \\end{center} \\caption{{\\it Left:} Celestial polar projection of the actual SACY observed sample. This projection reaches out to +10$\\deg$. Young stars are in filled circles. Note the concentration at the Sco-Cen complex (from $\\alpha$ 12 to 18). Crosses are giant stars observed in the survey. {\\it Right:} Celestial polar projection of the associations in the GAYA complex (see Section~3): the Tuc-Hor (filled circles), the Col (open circles)  and the Car  (crosses) associations. The transverse curve represents the Galactic plane. } \\label{fig:thpiz} \\end{figure}  \\begin{figure}[] \\begin{center} \\begin{tabular}{r} \\resizebox{1.0\\hsize}{!}{ \\includegraphics[draft=False]{BPIPIZN.EPS} \\includegraphics[draft=False]{BPIPIZ.EPS} } \\end{tabular} \\end{center} \\caption{Celestial polar projections of the $\\beta$~Pic Association. {\\it Left:} Northern Hemisphere. {\\it Right:} Southern Hemisphere. The transverse curve represents the Galactic plane.} \\label{fig:bppiz} \\end{figure}  \\begin{figure}[!ht] \\begin{center} \\begin{tabular}{r} \\resizebox{1.0\\hsize}{!}{ \\includegraphics[draft=False]{TWAPIZ.EPS} \\includegraphics[draft=False]{ECHPIZ.EPS} } \\end{tabular} \\end{center} \\caption{{\\it Left:} Celestial polar projection of the TW~Hya association. {\\it Right:} Celestial polar projection of the  $\\epsilon$~Cha Association -- the crosses represent the field  $\\epsilon$~Cha Association members and the filled circles the $\\eta$ Cha cluster members. The transverse curve represents the Galactic plane. } \\label{fig:echpiz} \\end{figure}  \\begin{figure}[] \\begin{center} \\begin{tabular}{r} \\resizebox{1.0\\hsize}{!}{ \\includegraphics[draft=False]{OCTPIZ.EPS} \\includegraphics[draft=False]{ARGPIZ.EPS} } \\end{tabular} \\end{center} \\caption{{\\it Left:} Celestial polar projection of the Octans Association. {\\it Right:} Celestial polar projection of the Argus Association -- the crosses represent the field Argus Association members and the filled circles the IC~2391 members. The transverse curve represents the Galactic plane. } \\label{fig:argpiz} \\end{figure}  \\begin{figure}[] \\begin{center} \\begin{tabular}{r} \\resizebox{1.0\\hsize}{!}{ \\includegraphics[draft=False]{ABDPIZN.EPS} \\includegraphics[draft=False]{ABDPIZ.EPS} } \\end{tabular} \\end{center} \\caption{Celestial polar projections of  the AB~Dor Association. {\\it Left:} Northern Hemisphere. {\\it Right:} Southern Hemisphere. } \\label{fig:abpiz} \\end{figure}   As will be more clearly explained in the next sections, the definition of the associations, especially their age, is strongly dependent on their low mass star population.  The depth limit of the SACY sample can give low mass candidates only for the associations nearer than $\\sim$50~pc.  For the more distant associations for which we found no other way to obtain low-mass candidates, their definitions are less reliable.  From Figures~\\ref{fig:thpiz} to ~\\ref{fig:abpiz} we can see that the Columba, $\\beta$~Pic and AB~Dor associations have members in the Northern Hemisphere and surveys in this part of the sky can reveal new members for these associations (and perhaps for the Tuc-Hor Association too).  Also other authors are trying to obtain associations or new members, mainly by the convergence point method \\cite[see, for example,][]{debruijne99}.  Our experience shows that this method has low reliability.  \\citet{song02b} arrive at a similar conclusion in a fruitless search for new TW~Hya Association members in a list of candidates proposed by \\citet{makarov01} using the convergence point method.  Another example is the moving group in Carina-Vela proposed by \\citet{makarov00}, as it is discussed in Sections 3.3 and 7 that seems to be, at least partially, a mixture of Car and Argus associations.  In some cases, visual binary components (or cluster members) have distinct proper motions or radial velocities. This may result in distinct convergence values, membership probabilities and, sometimes, distances. In the worst cases, only one component is a  high probability member, and only this one is presented. Evidently this is an indication that the kinematical data of the system needs to be improved.  Close visual binaries or spectroscopic binaries may deteriorate the quality of the kinematical data. An instructive example is HD~202947 (BS~Ind),  proposed by \\citet{zuckerman04} as a member of the Tuc-Hor Association. It has been observed 11 times in SACY and  found to be a double line spectroscopic binary. The Hipparcos catalogue detected an eclipsing binary light curve with a period of 0.435~days. \\citet{guenther05} studied the radial velocities of BS~Ind and found a period of 3.3~years, and the true systemic velocity could finally be derived. (There are also broad peaks in the cross correlation function both from SACY and \\citet{guenther05} that can be assigned to the eclipsing binary.) Due to its strong Li line BS~Ind is considered a young object, but for its space motions (U=+0.1, V=--25.3, W=--8.6) it can not belong to the Tuc-Hor Association or to any of the nearby associations presented here.  The young loose associations, being near to the Sun and having an age distribution from 5\\,Myr to about 100\\,Myr, enable studies of stellar physics that depend on the initial phases of the stellar evolution, for example multiplicity, abundances, rotation, stellar activity, etc. Examples of papers on rotation and activity using these associations are \\citet{delareza04} and \\citet{scholz07}. A preliminary study \\citep{kastner03} of the X-ray emission properties of members of the TW~Hya, $\\beta$~Pic and Tuc-Hor associations shows trends in the hardness ratio distributions when compared to TTS, the Hyades or older stars.  At a glance, looking at the figures of the Lithium distribution in the next sections, we can see that the Lithium depletion appears to be more scattered in the older associations. The present collection of several associations with ages from 6 up to 70 Myr (Table~\\ref{table:finalb}) furnishes an excellent opportunity to investigate the Lithium depletion during this time interval in the PMS evolution \\citep{silva08}. The Lithium depletion in associations can be compared with that in open clusters (see \\citet{jeffries06} for a review on this subject). The effect of rotation on the Lithium depletion in the PMS phase is also important. In fact, stars that are fast rotators (vsin(i)~$\\ga$~20\\,km~s$^{-1}$) have higher Li line equivalent widths (upper points of the Lithium distribution in the figures of the next sections). This is in agreement with what is found for cluster stars, at least for K type stars, in the sense that fast rotators appear to have high Lithium abundances \\citep{jeffries06}.  Most of the stars of the lists presented in the next sections stem from the SACY sample, that is, X-ray selected low mass stars. Any study using our sample must be aware of this bias. The data presented in Tables~\\ref{table:final} and \\ref{table:finalb} define the main characteristics and properties of these associations and can be used as starting points to find additional members. For example, GAIA will be an excellent tool to construct unbiased samples.  Stars in the young nearby associations  are ideal targets for studies of planetary formation and very low mass sub-stellar objects. As a matter of fact, many members of the associations presented here have already been  studied in the context of the characterization of their protoplanetary disks and the search for sub-stellar companions. These works are reviewed in the last section. ",
        "conclusions": ""
    },
    "0808/0808.1151_arXiv.txt": {
        "abstract": "We present a distance measurement for the semiregular variable S~Crateris (S~Crt) based on its annual parallax. With the unique dual beam system of the VLBI Exploration for Radio Astrometry (VERA) telescopes, we measured the absolute proper motion of a water maser spot associated with S~Crt, referred to the quasar J1147$-$0724 located at an angular separation of 1.23$^{\\circ}$. In observations spanning nearly two years, we have detected the maser spot at the LSR velocity of 34.7\\,km\\,s$^{-1}$, for which we measured the annual parallax of 2.33$\\pm$0.13\\,mas corresponding to a distance of 430$^{+25}_{-23}$\\,pc. This measurement has an accuracy one order of magnitude better than the parallax measurements of HIPPARCOS. The angular distribution and three-dimensional velocity field of maser spots indicate a bipolar outflow with the flow axis along northeast-southwest direction. Using the distance and photospheric temperature, we estimate the stellar radius of S~Crt and compare it with those of Mira variables. ",
        "introduction": "Very Long Baseline Interferometry (VLBI) is a powerful technique for obtaining positions of celestial objects with milliarcsecond\\ (mas) level accuracy. The VLBI Exploration of Radio Astrometry (VERA) telescopes are a Japanese VLBI array dedicated to phase referencing VLBI~\\citep{kob03}. VERA consists of four 20\\,m diameter antennas  at Mizusawa, Ogasawara, Iriki and Ishigaki-jima (see figure~1 of \\cite{pet07}). To overcome phase fluctuations and the limited integration time in conventional fast-switching VLBI, VERA has a dual beam system that allows simultaneous observations of the target and reference sources separated by 0.3 to 2.2 degrees. This advanced capability of VERA can be used to measure annual parallaxes and proper motions of masers with a 10 micro arcsecond ($\\mu$as) level accuracy.  An annual parallax gives a simple geometrical measure of the distance, free from the complex assumptions required in other distance estimators. Studies of Asymptotic Giant Branch (AGB) stars by \\citet{whi08}, \\citet{kur05}, and \\citet{vle03} are based on the annual parallax derived from VLBI observations, and demonstrated the the ability of VLBI astrometry. More recently, successful measurements of distances and proper motions with VERA have been reported ~\\citep{hir07,hon07,ima07,hir08,sat07}. In the present paper, we report the observations of S~Crt, concentrating on the distance measurement. S~Crt (IRAS11501$-$0719, IRC$-$10259, AFGL4830S) is an AGB star, with a pulsation period of 155\\,days\\,\\citep{ben90}. S~Crt has a variability type of SRb in the General Catalog of Variable Stars~\\citep{ant05}. The apparent magnitude at $K$ band is 0.786$\\pm$0.314\\,mag \\citep{cut03}. The optical light curve in $V$ band is presented in figure~\\ref{lightcurves}(a) exhibiting $\\sim$0.8\\,mag variation.  The distance to S~Crt has been measured in various ways. \\citet{bow94} and \\citet{pat92} determined the distance to be 420\\,pc and 285\\,pc, respectively, using the period-luminosity\\,(PL) relation of Mira variables. \\citet{per97} used HIPPARCOS data to measure a parallax of 2.04$\\pm$1.31\\,mas, corresponding to a distance of 490\\,pc, but with an large uncertainty. Recently, the new HIPPARCOS catalog~\\citep{van07} gives the parallax of 1.27$\\pm$0.92\\,mas, corresponding to a distance of 787\\,pc. VERA offers the potential to determine the parallax with a better accuracy so that the inconsistencies in these distance measurements can be resolved.  The observations and data reduction are described in section~2. In section~3, we present absolute positions and internal motions of maser spots. The annual parallax of S~Crt is presented and converted to the absolute distance in this section. Finally, in section~4, we examine the properties of S~Crt in the context of the new distance measurement.  \\begin{figure*}[htpb] \\begin{center} \\includegraphics[width=130mm, angle=0]{fig1.eps} \\caption{Light curves of S~Crt. (a)\\,The optical light curve at $V$ band provided by the ASAS \\citep{poj97}. (b)\\,The water maser light curve at 22\\,GHz obtained with the single-dish monitoring program with VERA~\\citep{shi08}. Peak flux densities of major velocity components are presented with different symbols. The symbols of ``$+$'', ``{\\Large$\\bullet$}'', ``{\\Large$\\circ$}'', ``$\\times$'', and ``$\\square$'' indicate the peak fluxes of the individual components at the LSR velocities of 34.7, 36.9, 38.5, 39.5, and 40.0\\,km\\,s$^{-1}$, respectively. There were no available data before the year of 2003.8. The hatched area indicates the period of the present VLBI monitoring observations. } \\label{lightcurves} \\end{center} \\end{figure*} ",
        "conclusions": "\\subsection{Errors in the positions} Since it is difficult to identify contributions to positional errors in individual VLBI observations ~\\citep{kur05, hac06, hon07, sat07, hir07, ima07}, we applied errors in each observation by evaluating following three error sources: (1)\\,a zenith atmospheric delay offset at each station, (2)\\,station position errors, and (3)\\,image quality. The best estimates of the zenith atmospheric delay offsets fell within the range of $\\pm$3\\,cm. Errors in the \\textit{a priori} station positions can result in additional phase residuals. Therefore, we assumed the offset of 3\\,cm even after the atmospheric offset calibration was done.  For a given separation of zenith angles between two sources $\\Delta Z$, the difference of signal path lengths $\\Delta l$ between the two directions caused by the zenith atmospheric delay residual $l_0$ can be estimated as follows, \\begin{eqnarray} \\Delta l &= l_0\\,\\Delta Z\\,\\frac{d}{d Z} \\left(\\frac{1}{\\cos Z}\\right), \\label{deltaSecZ} \\end{eqnarray} where $Z$ is the mean zenith angle of the observed sources. We adopted 1.23$^{\\circ}$ as the approximation of $\\Delta Z$. During our observations, the elevation angles higher than 30$^{\\circ}$ were dominant. We assumed $Z=50^{\\circ}$ as the zenith angle for estimating errors of the worst case, and obtained $\\Delta l =$ 0.12\\,cm. Using a root sum squares of major and minor axes of the synthesized beam\\,($\\theta_b$ = 1.7\\,mas), we estimated a positional error of $\\sigma_Z =$ 156\\,$\\mu$as from the following expression, \\begin{eqnarray} \\sigma_Z = \\theta_b \\times \\frac{\\Delta l}{\\lambda_{H_2O}}, \\end{eqnarray} where $\\lambda_{H_2O}$ is the wavelength corresponding to the rest frequency of water masers. Since the position accuracy of antennas are determined to be $\\sim$3\\,mm based on geodetic observations at S and X bands, we estimated the positional error\\,($\\sigma_D$) attributed to the baseline errors to be 5\\,$\\mu$as using the following expression, \\begin{eqnarray} \\sigma_D = Sin\\theta_{SA} \\times \\frac{3\\,mm}{\\lambda_{H_2O}}, \\end{eqnarray} where $\\theta_{SA} = 1.23^{\\circ}$ is the separation angle of the source pair. In addition, we estimated the measurement errors \\,($\\sigma_I = \\theta_b/SNR$) which depends on the SNRs of the phase-referenced images. Finally, we obtained the position errors of each phase-referenced image by adding up quadratically the error factors, $\\sigma_Z, \\sigma_D$, and $\\sigma_I$. The error bars in figure~3 represent the errors in each observation, and the averaged value of the errors is 167\\,$\\mu$as.  \\subsection{Maser Morphology} Figure~\\ref{fig-map} shows the internal motions of the masers pots around S~Crt found in the present observations. Since the maser spots show a bimodal distribution about radial velocities and positions, we extract the kinematic essentials. The analytic tools based on the diagonalization of the velocity variance-covariance matrix (VVCM)~\\citep{blo00} and position variance-covariance matrix are used.  The elements of the VVCM matrix are presented as \\begin{eqnarray} \\sigma_{jk} = \\frac{1}{N-1} \\sum\\limits_{i=1}^{N} (V_{j,i}-\\bar{V_{j}})(V_{k,i}-\\bar{V_{k}}), \\label{VVCM} \\end{eqnarray} where the diagonal elements are the velocity dispersion. Here $j$ and $k$ denote three orthogonal space axes (e.g., R.A., Dec., and radial coordinate $z$), $i$ denotes the $i$th maser spot in a collection totalling N($=$26), and the bar indicates averaging over maser spots.  The vector $\\mathbf{V}_1$~(figure~\\ref{fig-map}), which is a two-dimensional projections of the eigenvector for the largest eigenvalue, is the major principal axis of the VVCM. This axis of maximum internal velocity dispersion is the outflow axis. The axis of $\\mathbf{V}_1$ lies at a P.A. of $\\sim255^{\\circ}$, and also make an inclination angle $\\sim43^{\\circ}$ with the plane of the sky, with the southwest lobe directed into the page (away from the observer). The VVCM was initially evaluated in a coordinate system with axes R.A., Dec., and radial coordinate z increasing toward the source, and then it was diagonalized.  We also considered the two-dimensional variance-covariance matrix to quantify the angular distribution of the maser spots. Here, the number of maser spots is 26, same as the case of VVCM. In the diagonalization of this matrix, one eigenvalue which is larger than the other by a factor of 3.1 was obtained. The corresponding eigenvector $\\mathbf{V}_2$ presented in figure~\\ref{fig-map} indicate the spatial elongation of the maser distribution. The position angle of this axis is found to be $\\sim208^{\\circ}$. Although maser velocity and position axes show a discrepancy of $\\sim47^{\\circ}$, the directions of these model-independent axes indicate the bipolar outflow with the major axis along northeast-southwest.  This dynamical property is not revealed by previous observation \\citep{bow94}.  We estimate the stellar position from the present result by using the kinematic information. Since the motion of each maser spot are powered by the forming star, its position can be estimated as an origin of these motions. Here, we take one of the simplest kinematic models, which assumes a linear motion with a constant velocities from a single origin. This technique is essentially same as that used in \\citet{ima00} and \\citet{hon07}. Based on this modeling, we obtained the origin of the maser motions which most effectively explain the observed positions and velocities. The stellar position from this kinematic analysis, the kinematic center was obtained to be ($X$, $Y$) $=$ ($-$7.9$\\pm$1.8\\,mas, $-$14.9$\\pm$13.3\\,mas), which is indicated with the cross with error bars in figure~\\ref{fig-map}. We used the first order moment with respect to the determined stellar position in order to find a typical distance of maser spots. The obtained distance was 14.2$\\pm$3.5\\,mas, and presented with dashed line in figure~\\ref{fig-map}. At the distance of 430\\,pc, the corresponding linear length is 6.11\\,AU. Since we estimated the inclination angle of the outflow axis, this projected length can be converted to the linear radius of 8.96\\,AU.  Nineteen years ago, the angular distribution of the masers was firstly observed with the NRAO Very Large Array in its A configuration \\citep{bow94}. The masers were uniformly distributed over a 40\\,mas$ \\times$40\\,mas region with the LSR velocities ranging evenly from 33.4 to 44.2\\,km\\,s$^{-1}$ (figure~13 of \\cite{bow94}). In their study, the blue-(red-)shifted spots were found in the northeast(southwest) of the maser region. In addition, masers with intermediate velocities were found in the central region. Although the extent of the distribution is consistent with our result, we did not detect the masers in central region. In figure~\\ref{fig-map}, the maser spots are mainly found at the peripheral region of the distribution. During the period of 2006.3 -- 2007.0, we could not find the emission with intermediate velocities(see figure~\\ref{fig-spectr}(b)), and it can explain the the absence of intermediate velocity maser spots . On the other hand, at the beginning and end of our VLBI observations, the masers with intermediate velocities are found in total power spectrum, and they distributed in the peripheral region. As \\citet{bow94} did not precisely mention the flux densities and sizes of each spot, it is difficult to specify the reason for our non-detection of the masers in central region. However, we can speculate as to the reason; (1)\\,Our image sensitivities were not high enough to detect them. The rms noise 38\\,mJy\\,beam$^{-1}$ of \\citet{bow94} was lower than that of the present VLBI observations, 90\\,mJy\\,beam$^{-1}$. (2)\\,The masers in central region have faded before the present VLBI observations. (3)\\,The masers were extended and diffuse, therefore they were resolved out with the synthesized beam of VERA. For example, we find the effect of resolve-out (e.g. $\\leq$30\\,\\% in 2006 March) from the comparison between cross-power spectrum (figure~\\ref{fig-spectr}) and total-power spectrum (figure~\\ref{lightcurves}(b)).  \\subsection{Stellar diameter and temperature} We extract the stellar properties of S~Crt using the distance. \\citet{ari99} determined the temperature of photosphere ($T_{\\mathrm{BB}}$) in S~Crt to be 3097$\\pm$100\\,K by fitting two blackbodies to its infrared spectrum. Using this temperature, we estimated an acceptable temperature range of 2600 -- 3500\\,K. Here we assume that the light variation of 0.8\\,mag (figure~\\ref{lightcurves}(a)) is same in the infrared and caused only by the temperature variation. \\citet{ari99} reported that Mira variables show $T_{\\mathrm{BB}} =$ 2418 -- 2902\\,K and SRVs shows $T_{\\mathrm{BB}} =$ 3005 -- 3500\\,K. The estimated temperature range of S~Crt is consistent with their finding.  It is worthwhile estimating the stellar radius of S~Crt, since there is no direct measurement. Using the apparent magnitudes in infrared\\,(J, H, and K), $T_{\\mathrm{BB}}$\\,(3097\\,K) and the distance\\,(430\\,pc), we estimated the radius of S~Crt photosphere ($R_{*}$) to be $1.81\\pm0.14\\times10^{13}$\\,cm\\,(260$\\pm$20\\,$R_{\\odot}$). This is presented with a shaded-circle in figure~\\ref{fig-map}, and the center of circle is placed on the kinematic center. Using this radius, we estimated an acceptable radius range of 213 -- 309$R_{\\odot}$, under the assumption that the light variation is caused, now, only by the radius variation. In the studies of \\citet{han95} and \\citet{van97}, $R_{*}$s of dozen of Mira variables are presented, and most $R_{*}$/$R_{\\odot}$ is larger than 300. The photospheric radius of S~Crt is close to the lower limits to those of Mira variables in the literature. Since the estimates of $R_{*}$ and $T_{\\mathrm{BB}}$ are now available, the luminosity ($L_{*}$) of S~Crt is estimated to be 2.29$\\times$10$^{30}$\\,[W], resulting in the ratio $L_*$/$L_{\\odot}$ of 5970.  A similar analysis of data for a larger number of Mira and SRVs would be very useful for a better understanding of the difference and similarities between Mira variables and SRVs."
    },
    "0808/0808.3425_arXiv.txt": {
        "abstract": "% Next generation probes of dark matter and dark energy require high precision reconstruction of faint galaxy shapes from hundreds of dithered exposures. Current practice is to stack the images. While valuable for many applications, this stack is a highly compressed version of the data.  Future weak lensing studies will require analysis of the full dataset using the stack and its associated catalog only as a starting point. We describe a ``Multi-Fit\" algorithm which simultaneously fits individual galaxy exposures to a common profile model convolved with each exposure's point spread function at that position in the image. This technique leads to an enhancement of the number of useable small galaxies at high redshift and, more significantly, a decrease in systematic shear error. ",
        "introduction": "Dark energy affects the cosmic history of the Hubble expansion H(z) as well as the cosmic history of mass clustering. If combined, different types of probes of the expansion history and structure history can lead to percent level precision in dark energy parameters. This is because each probe depends on the other cosmological parameters or errors in different ways. These probes range from cosmic shear, baryon acoustic oscillations, supernovae, and cluster counting -- all as a function of redshift z. Using the CMB as normalization, the combination of these probes will yield the needed precision to distinguish between models of dark energy (Zhan 2006).  Next generation surveys will measure positions, colors, and shapes of distant galaxies over such a large volume that the resulting stochastic (random) errors will be very small. It is necessary to control and reduce the systematic errors to even lower levels.  There are two primary systematic errors which can influence the data: Photometric Redshift errors, and Weak Lens Shear errors.  The work to date has employed highly idealized data models. Here we describe some of the image processing challenges associated with reconstruction of the galaxy images from many dithered exposures.  With its capability to go deep, wide, and fast, the LSST will yield continuous overlapping images of 20,000 - 25,000 square degrees of sky.  The baseline exposure time is 15 seconds, and each ``visit\" to a single patch of sky will consist of two such exposures separated by a 2 sec readout with the shutter closed. In order to meet the science goals, six bandpasses ($u, g, r, i, z,$ and $y$) covering the wavelength range 320-1050 nm will be used.  The system is designed and will be engineered to yield exquisite astrometric and photometric accuracy and superb image quality. The telescope and camera optics and the detector combine to deliver 80\\% energy within a 0.2 arcsecond pixel over the full 10 square degree field and full wavelength range. This LSST survey will take ten years to complete. In a ten-year survey, the LSST will make more than five million exposures. In current simulations, the sky is tiled with a fixed spherical covering of circular fields. This overlap leads to a significant fraction of area which is observed twice as frequently as the average. In practice, the position of each visit will be varied continuously across the sky to average out this extra exposure.  How is the precision of shear measurements of distant galaxies in weak lensing tomography affected by ground-based seeing? Galaxy shape measurement depends on three parameters: galaxy size, delivered PSF, and limiting surface brightness. New ground-based telescopes are routinely delivering 0.4-0.7 arcsec FWHM imaging without adaptive optics. Clearly there are unique advantages in space for UV or IR imaging.  Galaxies at 25 mag have mean half-light radius ~ 0.4 arcsec and FWHM $\\sim 0.8$ arcsec.  Angular sizes of galaxies change with redshift due to a number of effects including the cosmological angle-redshift relation, luminosity evolution, and surface brightness dimming. The net effect is a plateau over a range of z, out to z=3 (Cameron and Driver 2007). At the low surface brightness reached in hundreds of LSST exposures, typical galaxies at redshift $z<3$ can be resolved sufficiently to measure their ellipticity. This is shown in Figure 1.  One must convolve with the PSF and ask if the ellipticity can be measured. Galaxies have a large intrinsic ellipticity (rms $\\sim 0.3$), and it is most important to have many of them in order to average down the shot noise of this intrinsic ellipticity.  At 28-29 mag per sq. arcsec ground based seeing is sufficient to measure the large ellipticities of 40-50 galaxies per square arcminute to the required $z<3$ redshift limit for tomography. However, it is crucial that shape systematics are minimized.  \\begin{figure} \\center \\includegraphics[width=0.7\\textwidth]{O5-3_1} \\caption{Galaxy surface brightness vs radius (arcsec) in one redshift bin from z = 2 - 3 for $23 < i < 25$ AB mag. This plot is from HST ACS imaging and ground based spectroscopy.  At the 28 $i$ and 29 $r$ mag/sq.arcsec limit of the LSST survey most galaxies at $z<3$ are sufficiently resolved in 0.6 arcsec FWHM seeing to reconstruct their ellipticity. (courtesy H. Ferguson)} \\label{fig:O5.3_1} \\end{figure} ",
        "conclusions": ""
    },
    "0808/0808.3878_arXiv.txt": {
        "abstract": "Warps in the outer gaseous disks of galaxies are a ubiquitous phenomenon, but it is unclear what generates them. One theory is that warps are generated internally through spontaneous bending instabilities. Other theories suggest that they result from the interaction of the outer disk with accreting extragalactic material. In this case, we expect to find cases where the circular velocity of the warp gas is poorly correlated with the rotational velocity of the galaxy disk at the same radius. Optical spectroscopy presents itself as an interesting alternative to 21-cm observations for testing this prediction, because (i) separating the kinematics of the warp from those of the disk requires a spatial resolution that is higher than what is achieved at 21 cm at low HI column density; (ii) optical spectroscopy also provides important information on star formation rates, gas excitation, and chemical abundances, which provide clues to the origin of the gas in warps. We present here preliminary results of a study of the kinematics of gas in the outer-disk warps of seven edge-on galaxies, using multi-hour VLT/FORS2 spectroscopy. ",
        "introduction": "Warps in the outer disks of galaxies are a ubiquitous phenomenon. They are seen both in the distribution of stars \\cite[(Sanchez-Saavedra et al. 1990; Cox et al. 1996)]{sanchezsaavedra,cox} and neutral hydrogen \\cite[(e.g., Sancisi 1976; Bosma 1981)]{sancisi,bosma}, and surveys estimate that possibly $50\\%$ or more of all galaxies show evidence for warps beyond the isophotal radius R$_{25}$ \\cite[(Briggs 1990)]{briggs}. (which typically contains $\\sim 90\\%$ of the total light). Usually, these manifest themselves in the form of the outer disk bending away from the plane defined by the inner disk, defining either the shape of a bowl, or, more commonly, an integral sign as seen edge-on. This suggests that specific angular momentum of material in the warps is not aligned with that of the inner disk.  Numerous suggestions have been made over the years for the responsible mechanisms. Among the earliest such proposals were internal bending modes in the disk \\cite[(Lynden-Bell 1965)]{lyndenbell}, but such modes were soon recognized to be persistent only in a disk with an unrealistically sharp mass truncation \\cite[(Hunter \\& Toomre 1969)]{huntertoomre}. \\cite[Revaz \\& Pfenniger (2004)]{revaz} have revived this discussion by identifiying short-lived bending instabilities as a possible cause. Other proposed mechanisms focus on the interaction of the baryonic disk with its environment: A bending of the outer baryonic disk may be introduced by a misalignment between the angular momentum of the inner baryonic disk and the hypothetical non-spherical, dark matter halo that it is embedded in \\cite[(Toomre 1983; Dekel \\& Shlosman 1983; Kuijken 1991; Sparke \\& Casertano 1998)]{toomre,dekel,kuijken,sparke}, creating a gravitational torque on the disk. However, it has been argued that the inner halo would realign with the baronic disk over time, and the warp would dissipate \\cite[(Nelson \\& Tremaine 1999; Binney, Jiang \\& Dutta 1998; Dubinski \\& Kuijken 1995; New et al. 1998)]{nelson,bjd,dubinski,new}. \\cite[Ostriker \\& Binney (1989), Jiang \\& Binney (1999), Shen \\& Sellwood (2006)]{ostriker,hiang,shen} consider the impact of ongoing accretion onto an outer dark matter halo and argue that, since the angular momentum of infalling material will in general not be aligned with the present galaxy disk, this will create an ongoing torque on the outer galaxy disk. \\cite[Binney (1992)]{binney91} has furthermore suggested that, if infalling material loses angular momentum to the halo, it might penetrate as far as the outer edge of the baryonic disk itself.  It is this hypothesis that we wish to test with the present work. How galaxies acquire gas is one of the key questions in our understanding of how they evolve,  and determining whether warps are indeed signatures of such processes therefore is of great importance.  How can the direct accretion hypothesis be tested? The specific angular momentum vector of accreting material will, in general, neither be aligned exactly with that of the inner disk, nor have the same size. If such infalling material is indeed in direct contact and exchanging angular momentum with gas in the outermost baryonic disk, then it will introduce kinematic anomalies, i.e., a deviation from disk-like rotation, such as a lag or excess in the circular velocity. Measuring the circular velocity of gas in the warps therefore becomes an important observational discriminator.  Our project has measured line-of-sight velocities of gas in the outer disks of seven galaxies, the majority of which display clear signs of optical warps. Our measurements were obtained via optical spectroscopy of the H$\\alpha$ line. Although the tradititional way of observing gas in the outer disk is via the 21-cm line of neutral hydrogen, optical spectroscopy has proven a surprisingly successful alternative for studying the outer disk \\cite[(Bland-Hawthorn, Freeman \\& Quinn 1997;Christlein \\& Zaritsky 2008)]{bhfq,cz}, for a number of reasons: 1) Most importantly, the spatial resolution --- an order of magnitude better than even the highest-quality interferometric HI maps --- allows us to clearly separate gas in the warps from gas in the plane of the galaxy. This, in turn, allows us to access galaxies with smaller angular diameters at larger redshifts, greatly increasing the number of suitable targets. 2) Sporadic local star formation or illumination of gas in the warps by escaping UV flux from the inner star-forming disk guarantee low-level H$\\alpha$ flux from the outer disk far beyond what is usually perceived as the star formation threshold. It has been demonstrated \\cite[(Bland-Hawthorn, Freeman \\& Quinn 1997;Christlein \\& Zaritsky 2008)]{bhfq,cz} that multi-hour optical spectroscopy of such low-level emission can, in some cases, probe outer galaxy disks to similar extents as 21-cm. 3) Optical spectroscopy yields a plethora of ancillary data, particularly stellar continuum and metal lines. In determining the origin of gas in the outer disks and warps, metallicity indicators may provide valuable additional insight (although an analysis exceeds the scope of these proceedings). Our aim is to a) measure the rotational velocity of gas in the warps and b) determine whether there are kinematic anomalies in this rotational velocity, i.e., sudden breaks or upturns that are not consistent with an extrapolation of the rotation curve from the inner disk, which may be associated with the onset of the warp. ",
        "conclusions": "In an effort to measure the kinematics of gas in the warps in the outer disks of galaxies with high spatial resolution, we have targetted seven edge-on galaxies, most of them with known optical warps, using multi-hour long slit spectroscopy.  \\begin{itemize} \\item In at least three of seven cases, we have observed more extended H$\\alpha$ emission if the slit was aligned along the warp rather than along the major axis. This shows that we have been successful at detecting gas in the warps, and that we have thus attained our observational goal. Counting the single emission region detected in UGC 12423, this ratio is four out of seven. \\item In three of these four cases, there are no signs of kinematic anomalies in the rotation curve associated with the onset of the warp. \\item In one of these four cases, there is a sharp break in the rotational velocity roughly at the onset of the warp. In this case, the optical appearance of the ``warp'' is highly suggestive of a tidal feature. However, the break occurs in the major axis rotation curve. \\item In the remaining objects, the outer-disk features are sampled by the major-axis spectrum, but there is no evidence for kinematic anomalies in these regions. \\end{itemize}  Based on a small data set (four kinematic detections of extraplanar outer disk gas likely to be associated with a warp), we find evidence for kinematic anomalies only in the one case whose optical appearance strongly suggests a tidal feature. In other cases, the kinematics of the presumed warp gas are consistent with the extrapolation of the inner disk rotation curve. This sample is therefore not in support of the hypothesis that the outer, warped disks are in direct contact and exchanging angular momentum with material being newly accreted; however, the possibility cannot be ruled out yet that the real point of contact between the galaxy disk and accreted material lies at larger radii, and that the gas that we see at the onset of the optical warp has already settled into disk-like kinematics. We propose to explore three avenues for further investigation: (i) Gas at larger radii, substantially beyond the optical warps, should be targetted; ideally, targets should thus be selected on the basis of HI maps where available. (ii) The sample should be increased in size from the present three clear detections. (iii) Additional indicators that might constrain the origin of the warp gas, in particular, metallicity, should be included in the analysis.  Based on observations made with ESO Telescopes at the Paranal Observatories under programme ID $<$079.B-0426$>$."
    },
    "0808/0808.3068.txt": {
        "abstract": "Future missions such as Solar Orbiter (SO), InterHelioprobe, or Solar Probe aim at approaching the Sun closer than ever before, with on board some high resolution imagers (HRI) having  a subsecond cadence and a pixel area of about $(80km)^2$ at the Sun during perihelion.  In order to guarantee their scientific success, it is necessary to evaluate if  the photon counts available at these resolution and cadence will provide a sufficient signal-to-noise ratio (SNR).  For example, if the inhomogeneities in the Quiet Sun emission prevail at higher resolution, one may hope to locally have more photon counts than in the case of a uniform  source. It is relevant to quantify how inhomogeneous the quiet corona will be for a pixel pitch that is about 20 times smaller than in the case of SoHO/EIT, and 5 times smaller than TRACE.   We perform a first step in this direction by analyzing and characterizing the spatial intermittency of Quiet Sun images thanks to a multifractal analysis. We identify the parameters that specify the scale-invariance behavior. This identification allows next to select a family of multifractal processes, namely the Compound Poisson Cascades, that can synthesize artificial images having some of the scale-invariance properties observed on the recorded images.  The prevalence of self-similarity in Quiet Sun coronal images makes it relevant to study the ratio between the SNR present at SoHO/EIT images and in coarsened images. SoHO/EIT images thus play the role of \\lq high resolution' images, whereas the \\lq low-resolution' coarsened images are rebinned so as to simulate a smaller angular resolution and/or a larger distance to the Sun. For a fixed difference in angular resolution and in Spacecraft-Sun distance, we determine the proportion of pixels having a SNR preserved at high resolution  given a particular increase in effective area. If scale-invariance continues to prevail at smaller scales, the conclusion reached with SoHO/EIT images can be transposed to the situation where the resolution is increased from SoHO/EIT to SO/HRI resolution at perihelion. ",
        "introduction": "Many fine-scale structures in the corona do not seem to be well resolved by current imaging telescopes. Filamentary and threaded patterns are observed in coronal loops~(e.g.~\\citet{2007ApJ...661..532D}); in the Quiet Sun (QS), small dynamical events such as brightenings or blinkers point at unresolved sub-structures. Having sufficient resolution is also necessary to assess whether nanoflares occurring in QS may explain coronal heating ~\\citep{1988ApJ...330..474P,1998ApJ...501L.213K,2007ApJ...659.1673A,2001A&A...373..318M,2002ESASP.506..501B}. Several studies~\\citep{1998A&A...336.1039B,1998ApJ...501L.213K,2000ApJ...544..550A} present evidence in favor of a turbulent mechanism, with individual dissipative structure far below the instrumental resolution limit.  Scale-invariance has been observed down to the smallest reachable scales, namely down to the PSF dimensions of today\u0092s highest resolution instruments. This paper studies this scale-invariance (or self-similary) property in the case of a SOHO/EIT data set, first by computing the Fourier spectrum, and next by performing a multifractal analysis. But, should this property persist to even smaller scales? To answer this question, we can only rely on current observation extrapolations and on physical theory and modeling. We now briefly discuss some of those approaches.  A first element stems from the consistency between quantities observed by both EIT and TRACE, albeit at resolutions differing by a factor five in 1D ($\\times 25$ in 2D). For example, \\citet{2002ApJ...568..413B} show that EIT and TRACE studies agree about the distribution of energy released by flare-like events in the Quiet Sun. Power-law distributions fit both data, and they find similar values for their slopes. Along a similar line, a H\\\"older analysis in~\\citet{dch:spatial_noise} evidences a continuity between supra- and intra-pixel scales. Intra-pixel variability is therein estimated using a one-minute cadence EIT data set, for which solar rotation induces a displacement (from one image to the next) that is less than the pixel. To go to smaller scales, note first that regions of strong magnetic field, such as bright points, exhibit brighter EUV emission, while regions of weak magnetic field lead to a darker emission. Hence scale-invariance in the EUV emission might follow from similar properties in the magnetic field strength. The tectonics model of~\\citet{2002ApJ...576..533P} suggests a highly fragmented photospheric magnetic configuration in the quiet network, where the fundamental units of flux  are intense flux tubes having a 100km diameter. Finally, much literature has been devoted to the study of distribution of flare properties, see e.g.~\\citep{1998A&A...336.1039B,2000ApJ...535.1047A,2004ApJ...603L..61V}. Power-law distributions for e.g.~the energy release and volume occupied by flares might be explained by self-organization and the cascading nature of flare activities. Such statistical theory of flare activity are physically motivated by the turbulent nature of the solar corona~\\citep{2003matu.book.....B,1996ApJ...457L.113E} where the dissipative scale is estimated to be of a few meters~\\citep{1994SSRv...68...97E}.   One notices that the above reasoning relies on rather hypothetical guesses, and this is an additional motivation to actually plan high resolution observations of the QS in coronal lines. We are also aware that the observed intensity in a given EUV optically thin spectral line entangles the temperature and the emission measure. The above issue is worsened by the possible contamination from other lines in the passband, and by the projection effects. It remains nevertheless of interest to prepare the conceptual tools that will enable the analysis of better quality data when they come.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Beyond the Fourier spectrum,  a multifractal analysis allows to further characterize the spatial scale-invariance, or more generally the so-called \\lq intermittency' of a stochastic process~\\citep{cd95,frischbook95,amr97,chap99}. It has been extensively used for the statistical modeling of turbulent flows~\\citep{frischbook95}, Internet traffic~\\citep{fgw98}, natural~\\citep{tmp98,chainaistpami07} and meteorological images~\\citep{rad00,gtyh07}, as well as ionospheric indices~\\citep{PhysRevLett.76.4082}.  In solar physics, studies on the fractal dimension was achieved for Quiet Sun EUV network~\\citep{1998A&A...335..733G}, for the spatial extend of nanoflares events~\\citep{2002ApJ...572.1048A}, and for active regions~\\citep{2005ApJ...631..628M}. \\citet{law-cad-ruz:multiplic} present a  multifractal analysis of photoelectric images of line-of-sight magnetic fields in solar active regions and quiet photosphere. They consider a multiplicative cascade approach in order to infer a scale-invariant rule for the Ohmic dissipation measure. This rule is then used to re-estimate with greater accuracy  the multifractal spectrum. Multifractal analysis has been also applied to the analysis of solar magnetograms~\\citep{2005SoPh..228....5G,2005SoPh..228...29A}, and of the temporal variation of the emission observed in several RHESSI X-ray energy bands~\\citep{2007ApJ...662..691M}.   In this work, we aim at generating synthetic images at EIT resolution that capture several statistical properties of Quiet Sun images. To do this with minimum a priori, we propose to exploit the statistical scale invariance observed in data. We first perform a multifractal analysis. It yields a set of parameters, namely the multiscaling exponents, that quantitatively describe this scale invariance. Next we aim at finding a family of stochastic processes that obey the same statistical property with the same set of parameters, thus injecting minimum a priori in the model. The corresponding synthetic images are exempt of spurious artifacts such as  cosmic ray hits, projection effect, or error sources in the data. This is the first achievement necessary in the elaboration of a way to create synthetic images at arbitrary resolution. The analysis and parameter identification step is essential in order to know how to extrapolate the properties of EIT images at higher resolution. Moreover, our methodology for synthesizing Quiet Sun EUV images is new: having observed certain properties on EIT images, the associated parameters are then used within a multifractal stochastic process that can simulate EIT images.  The are two main applications of this procedure in solar physics. First, artificial EUV images can be used for the testing and calibration of automatic feature finder procedure, see e.g.~\\citet{2007A&A...464.1107G}. Second, and most importantly for the purpose of this paper, since QS images exhibit self-similarity behavior that can be reproduced with a multifractal stochastic process, it makes sense to study the relationships between intensities across different scales of observations. More precisely, we can look at the statistics of the ratio between intensity values at full resolution and at a rebinned version of the original images. The conclusion reached here can be transposed to higher resolutions as long as the scale-invariance continues to prevail. Our study therefore provides some guidance for radiometric studies related to future high resolution missions. The latter are much needed since the spatial resolution of current telescopes precludes definite conclusion about the fundamental processes that determine the existence and the underlying physics of the transition region and of the  quiet corona. In the following, we are considering in more details the case of High Resolution Imagers (HRI) on board Solar Orbiter (SO).   HRIs on board SO will produce at perihelion images having a pixel area at the Sun of $(80\\mbox{km})^2$. Evaluating the loss in SNR when going from low to high resolution is necessary in order to guarantee the scientific success of this high resolution mission. Indeed, the smaller the area covered by a pixel, the lesser the signal, and the fainter and more dynamical the target. Similarly, the higher the cadence, the shorter the exposure time, and the lesser the signal again.   Considering that the energy $E$ emitted from the Sun follows a Poisson process, and hence that $\\mbox{SNR} \\leq E/\\sqrt{E}$, $\\mbox{SNR}^2 \\leq E$, \\cite{2001ESASP.493..245H} describe the relationship between the quality of  observations (left hand side) and the experimental conditions (right hand side) as follows: \\begin{equation} \\label{E:SNRineq} \\frac{\\text{SNR}^2}{A_{\\mbox{sun}}  T_{\\mbox{ExpTime}}} \\leq \\frac{A_{eff} L}{d^2}~, \\end{equation} where $A_{sun}$ is the Sun area covered by one pixel, $T_{ExpTime}$ is the exposure time, $A_{eff}$ is the effective area of the telescope (in $m^2$),  $L$ is the radiance (in $W sr^{-1} m^{-2}$), and $d$ is the distance between the spacecraft and the Sun.   When comparing the situation of EIT with the one foreseen for an HRI on board SO, we observe that getting closer to the Sun provides a factor $d^2 = 25$, but since $A_{sun}$ is divided at the same time   from $(1800km)^2$ (EIT) to $(80km)^2$ (HRI), the effective area $A_{eff}$ needs to be enhanced by a factor of $(1800/80)^2 /25 \\simeq 20$ or more  in order to preserve a SNR for a uniform source in space. Moreover, if one takes into account the temporal variability of solar features, this factor gets even larger: ~\\citet{1998A&A...336.1039B} showed that the typical duration of an event was proportional to its size. If this argument still holds at a scale of $(80\\mbox{km})^2$, HRI should have an exposure time that is $(1800/80)^2$ times less than in the case of EIT. The effective area should then increase by a factor $\\left(\\frac{1800}{80}\\right)^4/25 \\simeq 10^4$ in order to have a SNR similar to EIT images.  Our main  contribution in this paper is to refine this argument by taking into account the inhomogeneities present in the solar corona. Indeed, these might provide an independent help in the quest for a sufficient SNR, with subsets of high-resolution pixels likely to carry a large part of the energy present at a lower resolution.   Towards quantifying the needed gain in $A_{eff}$, we first re-write~(\\ref{E:SNRineq}) in a more general form as \\begin{equation} \\label{E:SNRineq2} SNR^2 \\leq \\frac{A_{eff}}{d^2}  \\int_{T_{ExpTime}}\\int_{A_{sun}}  L(s,t)dtds~, \\end{equation} which now considers the time dynamics and possible spatial inhomogeneities in the emission process. In this paper, we consider a fixed exposure time and a fixed effective area, and we investigate  how the spatial distribution of the radiance will affect the SNR in the case of Quiet Sun data. As mentioned above, the temporal variability is also necessarily non-uniform;  the methods presented here will be adapted in the future to see how this variability affects positively the SNR in case of a high cadence instrument.  We  are interested in the evolution of the quantity \\begin{equation} \\label{E:eqQ} Q=  \\frac{\\int_{A_{sun}} L(s)ds}{d^2}~, \\end{equation} with a higher resolution. More precisely, we simulate the degradation in resolution starting from EIT images using an appropriate coarsening of the data. Taking  the ratio between the value of $Q$ obtained at low (coarsened) and high (EIT original)  resolution, we obtain an estimate of the gain in effective area needed in order to keep at high resolution the SNR (as defined in~\\refeq{E:SNRineq}) larger or equal to its value at low resolution.  If the scale-invariance continues to prevail down to a $(80km)$ scale, our conclusion about the needed gain in $A_{eff}$ may be transposed to the situation where the resolution is increased from EIT to HRI resolution at perihelion.   In future work, we aim at generating images directly at HRI resolution at perihelion. However, additional difficulties arise in this case because not only must the multifractal spectrum  be respected, but the histogram of rebinned images must fit the histogram of EIT images as well. Indeed, the multifractal stochastic model used in this paper to generate QS images does not constrain the simulated images to have the same histogram as the real ones. (Both histograms are however fairly close for images generated at the same scale of observation as the physical ones.)     This paper is structured as follows. Section~\\ref{S:obs} describes the observations considered here with a general overview of noise sources present in EIT images, followed by a presentation of the data set used in this paper. In Section~\\ref{S:waveletMF} we first recall the  basics behind wavelet-based multifractal analysis. Next we present our multifractal analysis of Quiet Sun images and discuss its limitation. With the set of parameters identified in Section~\\ref{S:waveletMF}, we are able to synthesize in Section~\\ref{S:synthe} Quiet Sun-like images. The real and artificial EIT data sets are used in Section~\\ref{S:fromEIT} to study  the effects of the transition from low to high resolution when taking into account the multifractal behavior of the observations. We then discuss the implications of our study in the construction of high-resolution instruments. %Finally, Section~\\ref{S:conclusion} details the difficulties behind generating stochastic images at any high resolution and outlines our work in progress in this respect.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{S:conclusion}   Amongst the many challenges faced by high resolution missions, the precise estimation of photon counts is of fundamental importance. We showed in this paper that an accurate radiometric model must take into account the spatial inhomogeneities present in the source.  We first showed how to characterize these spatial inhomogeneities through a multifracal analysis of Quiet Sun images: We computed the multifractal spectrum and derived a  set of parameters, namely the multiscaling exponents, that quantitatively describe the scale invariance properties. We then proposed a family of stochastic processes that obey the same statistical property with the same set of parameters, thus injecting minimum a priori in the model. This establishes firmly the spatial scale-invariance structure of EIT images.   Next, we compared a set of images (real and simulated via the above stochastic process) at EIT-resolution with an appropriate rebin of these images. The rebin is done so as to simulate a larger angular resolution, and a larger distance to the Sun. By comparing the intensity values at low and high resolutions, it is possible to estimate the needed gain in effective area such that the SNR at high and low resolutions remains the same. If the scale-invariance observed  at supra-pixel scales on EIT images continues to prevail at smaller scales, the distribution of this ratio will be the same when comparing current low resolution EIT images with future high resolution images. In this sense, we  provided a proxy for the needed gain in effective area for HRI telescopes.   With the new technologies, the future HRI instrument will have an effective area which is approximately 10 times larger than the one of EIT. In this case, our study shows that about $0.2\\%$ of the pixels in a HRI QS image will have a SNR similar to what is observed on EIT images. However, if the effective area of HRI becomes 15 times larger (instead of 10 times), about $1.6 \\%$ of the pixels (i.e.~eight times more pixels) will have an excellent SNR.  Hence a relatively small increase in instrumental performance may have a large impact on the quality of the data.  The results presented in this paper are based on images of the warm corona. X-ray images of the hot corona are likely to exhibit a more pronounced spatial intermittency;  a similar analysis on X-ray images should therefore conclude that  a smaller gain in effective area is needed. As a general conclusion we can say that 1/ all other things being equal, the quality of observations will improve when getting closer to the Sun\\footnote{However, as said before, a closer distance to the Sun might imply to build a telescope with a smaller pupil, and hence a reduced radiant intensity.}; 2/ regions where the source is uniform will be penalized by an increase in angular resolution, but irregular source regions will benefit from such an increase.    In order to help improve current radiometric model, our next goal is to generate synthetic images at arbitrary resolution (e.g.~EIT resolution or higher) which would capture as many statistical properties of Quiet Sun images as possible. Indeed, Figure~\\ref{figMFanal} shows that we have succeeded in building synthetic images from multifractal processes that obey some of the scale-invariance properties of Quiet Sun EUV images, more precisely that have the same multiscaling exponents. However, in Figure~\\ref{F:ratiolowNlarged} we see that the present model underestimates the proportion of bright (i.e. SNR-preserved) pixels.  This is related to  the discrepancy  between the histograms of real and artificial EIT images observed for high intensity values. Indeed, our proposed model captures the self-similar and multifractal nature of Quiet Sun images, but does not take into account other constraint such as the distribution of intensities in EIT images. %Despite its apparent straightforwardness, it is difficult from a mathematical point of view to accurately take into account such constraint.  In future work, we plan on exploring two alternatives to obtain a model that would both preserve the  histogram observed at a given resolution and satisfy the multifractal properties. The first one is to find a way to constrain the synthetic processes to comply with  some properties on their marginal distributions (like e.g. intensity histograms) as well as with the desired scale invariance properties. The difficulty here comes from the fact that the multifractal spectrum strongly constrains the global structure of stochastic process. Hence there are few degrees of freedom left to further adopt  other criteria. This issue remains thus an open question for mathematicians. The second possibility, more empirical, is to build  higher resolution images from low resolution images. The challenge here is  to extrapolate e.g.~EIT images at higher resolution using the scale invariance properties. In this case, the real image would serve as a boundary condition of the model, and by construction the histogram of the low resolution image would be preserved.    %%%%%%% APPENDIX -------------------------------------------- \\appendix  %%%%%%%%%%%%%%%%%%%"
    },
    "0808/0808.0389_arXiv.txt": {
        "abstract": "{Gravitational Wave Astronomy is becoming a reality as Earth-based interferometric gravitational-wave detectors reach the design sensitivities and move towards advanced configurations that may lead to gravitational-wave detections in the immediate future.  In this contribution, I briefly summarize the basic characteristics of this new area, the discovery prospects and the potential for fundamental physics.  Then, I present results of some investigations of two different sources of gravitational waves that are potential targets for present and future planned observatories.  First, I will discuss the generation of gravitational radiation by non-linear effects arising from the coupling between radial and non-radial oscillations of neutron stars, which may produce distinctive gravitational-wave signatures.  The gravitational radiation emitted by these sources is in the frequency band of Earth-based detectors. And second, I will discuss the gravitational-wave emission during the inspiral of extreme-mass-ratio compact binaries.  In this case, the gravitational waves have low frequencies, inside the frequency band of space observatories like LISA.}   \\FullConference{Supernovae: lights in the darkness (XXIII Trobades Cient\\'ifiques de la Mediterr\\`ania)\\\\ October 3-5 2007\\\\ Ma\\'o, Menorca, Spain}  \\begin{document} ",
        "introduction": "Detection and study of the gravitational radiation coming from astrophysical and cosmological sources is the main task of the emergent area of Gravitational Wave (GW) astronomy.  This task is coming closer to reality as Earth-based detectors like LIGO~\\cite{ligo} and VIRGO~\\cite{virgo} are improving their sensitivities. At the same time, there is a strong international effort to construct a GW observatory in space, the {\\em Laser Interferometer Space Antenna} (LISA)~\\cite{lisa}, an ESA-NASA mission that is expected to be launched in the next decade, with an ESA precursor mission, LISA PathFinder~\\cite{pathfinder}, that will be launched in 2010.  Gravitational radiation has physical properties that make it a very attractive tool for astronomy.  In particular, the weak coupling of gravity with matter implies that the GWs that our observatories will detect carry nearly uncorrupted information from the sources that produced them.  At the same time, this implies that GWs are quite difficult to detect since they will also interact weakly with the detector.  In many cases the signal will be buried in the instrumental noise, and in some cases, in GW backgrounds (for instance, LISA will be limited by the GW background produced by galactic stellar compact binaries).  Therefore, an important task of GW astronomy is to construct theoretical waveform templates in order to separate in an efficient way the signal from the noise using statistical techniques like match filtering. Depending on the GW source, these templates can be obtained by full numerical relativity simulations, relativistic perturbation theory, or post-Newtonian approximation schemes, or by a combination of them. This is a research area in GW astronomy where a significant part of the theoretical efforts have been directed to.  There are other important topics that theory is dealing with that are important for the development of GW astronomy.  In particular, the study of the formation mechanisms and rates of GW sources.  These studies are relevant for the design of the GW observatories.  They can also determine the regions of the configuration space of the potential GW sources that correspond to astrophysically realistic situations. This last point is crucial to establish the parameter space that GW template banks have to cover, and hence to determine the range of initial conditions for the simulations of GW sources. For instance, in the case of black hole (BH) binaries we are interested in the range of orbital parameters and spin orientations that we can expect from each of the possible astrophysical scenarios that yield BH binaries. Finally, another important arena for theoretical studies is the investigation of possible electromagnetic counterparts to GW events (GRBs, X-rays, etc.) and their impact on on GW astronomy, whereby synergies between electromagnetic and GW observations are explored.   The main sources of gravitational radiation for Earth-based interferometric detectors are: stellar-mass compact binaries [BH-BH, BH-Neutron Star (NS), and NS-NS binaries], supernovae core collapse, relativistic pulsars, oscillations of neutron stars (w-modes, r-modes, etc.), and stochastic sources (astrophysical and cosmological). On the other hand, the main sources for space-based detectors like LISA are: Galactic compact binaries [e.g. White Dwarf (WD) binaries], massive BH binary mergers, capture of stellar-mass compact objects by massive BHs sitting at galactic centers [known as extreme-mass-ratio inspirals (EMRIs)], and cosmological backgrounds. Remarkable advances in astrophysics, cosmology and gravitational physics are expected from observations of these systems. To illustrate this, let us mention a few outcomes of GW astronomy: Direct proof of the existence of Gravitational Radiation and its properties (we have an indirect proof of their existence through the observations of the binary pulsar PSR B$1913+16$~\\cite{binarypulsar} discovered by Hulse and Taylor~\\cite{hulsetaylor}); information about masses, spins, and orbital parameters for relativistic compact binaries; tests of the {\\em no-hair} theorem which, roughly speaking, says that astrophysical BHs are characterized only by their mass and intrinsic angular momentum; information about the equation of state of NSs, rotational properties, progenitor masses, etc.; tests of galaxy formation models and understanding of the growth history of massive BHs; study of the expansion history of the universe  by combining gravitational and electromagnetic observations of sets of GW \\emph{standard candles} like massive BH binary mergers or EMRI events; etc. As a consequence, it is expected that GW astronomy will open a new window to the exploration of the Universe in a similar way as it happened in the past with the opening of new electromagnetic windows.   Moreover, there is a considerable potential for observation of unexpected GW sources, and it is also possible that these sources may provide information relevant not only for astrophysics and cosmology but also for fundamental physics.  An outcome of GW astronomy that deserves a separate comment is the possibility of testing General Relativity and putting constraints on alternative theories of gravity~\\cite{cliffordwill}.  The first such constraints came from solar system  observations, but the gravitational fields involved in the solar system are weak and the velocities are small as compared with the speed of light. Therefore, these observations can only test the non-radiative weak-field limit of relativistic theories of gravity.  The situation changed after the discovery of the binary pulsar PSR B1913+16~\\cite{hulsetaylor} (which has been followed by other binary pulsars: PSR B1534+12, PSR J1141-6545, PSR J1829+2456, PSR J0737-3039), as this system involves regions with strong gravitational fields (inside and outside the components of the binary). The way in which certain alternative theories of gravity can be constrained by binary pulsar observations is by the timing of phenomenological \\emph{post-Keplerian} parameters (see~\\cite{damour} for a detailed and recent review of tests of gravity theories by using binary pulsars).  For a given theory, one can compute these post-Keplerian parameters (eight) in terms of the Keplerian ones (five) and the masses of the binary (two).  Therefore, the measurement of a post-Keplerian parameter determines a curve in the mass plane (the plane of the possible masses of the binary components). Then, the measurement of three post-Keplerian parameters constitutes a test of the gravity theory under consideration.   Such a test is passed if the three curves intersect at one point (if one adds the measurement error bars the {\\em thick} curves should just have a non-empty intersection).  In the cases in which $n$ ($n\\geq 3$) post-Keplerian parameters are measured, $n-2$ independent tests of the gravitational theory can be performed.   However, the tests just mentioned obviously do not involve GW measurements. To illustrate the potential of GW astronomy in this respect, let us mention the constraints that LISA observations can impose on the coupling parameter of scalar-tensor theories of the Brans-Dicke type, $\\omega^{}_{BD}$. In~\\cite{willyunes}, assuming a $10^3\\,M^{}_\\odot~BH\\;+\\;1\\,NS$ inspiral at $50\\;Mpc$  and $1$ year of integration prior the last stable orbit, it was shown that it is possible to get the following bound: $\\omega^{}_{BD} > 3 \\times 10^5$ (the Cassini mission~\\cite{cassini} yields $\\omega^{}_{BD} > 4\\times 10^4$). Moreover, it was also shown that LISA observations can also impose bounds on the Compton wavelength $\\lambda^{}_g$ of the graviton~\\cite{willyunes}: Assuming a $10^6\\,M^{}_\\odot$ BH binary inspiral at $3\\;Gpc$ we can get $\\lambda^{}_g > 5.4\\times 10^{16}\\;km$ (4 orders of magnitude larger than with solar system dynamics).  In~\\cite{lisaalternative}, these estimations were revised by including the effect of spin couplings.  It was found that the bound on the Brans-Dicke parameter is significantly reduced by spin-orbit and spin-spin couplings (factors of $10$ to $20$), while the bound on the graviton Compton wavelength is only marginally reduced (factors of $4$ to $5$).  As it has been mentioned in this introduction, NSs and BHs (of different sizes) are among the most important sources of GWs, mainly due to their strong gravity.  However, this same property also makes them very challenging for the task of making theoretical predictions of the shape and amplitude of the GWs that they emit. In what follows, I will describe some work on the topic of the description of sources of GWs in which I have participated with other researchers in the recent past.  More specifically, I will discuss some results on the description of two very different types of sources of GWs: (i) GWs from non-linear oscillations of relativistic stars, which are relevant for Earth-based GW observatories. (ii) Simulations of EMRI events and the GWs emitted by them, which is of interest for space-based observatories like LISA. ",
        "conclusions": ""
    },
    "0808/0808.0176_arXiv.txt": {
        "abstract": "The light curve of an exoplanetary transit can be used to estimate the planetary radius and other parameters of interest. Because accurate parameter estimation is a non-analytic and computationally intensive problem, it is often useful to have analytic approximations for the parameters as well as their uncertainties and covariances.  Here we give such formulas, for the case of an exoplanet transiting a star with a uniform brightness distribution. When limb darkening is significant, our parameter sets are still useful, although our analytic formulas underpredict the covariances and uncertainties. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.0206_arXiv.txt": {
        "abstract": "Stars in disks of spiral galaxies are usually assumed to remain roughly at their birth radii. This assumption is built into decades of modelling of the evolution of stellar populations in our own Galaxy and in external systems. We present results from self-consistent high-resolution $N$-body + Smooth Particle Hydrodynamics simulations of disk formation, in which stars migrate across significant galactocentric distances due to resonant scattering with transient spiral arms, while preserving their circular orbits. We investigate the implications of such migrations for observed stellar populations. Radial migration provides an explanation for the observed flatness and spread in the age-metallicity relation and the relative lack of metal poor stars in the solar neighborhood. The presence of radial migration also prompts rethinking of interpretations of extra-galactic stellar population data, especially for determinations of star formation histories. ",
        "introduction": "In a universe where mass assembles by accretion of progressively larger constituents, thin disks of galaxies form during quiescent evolution following the major merging epoch \\citep{brook:2004,Robertson:2006}. The inner parts, with lower angular momentum assemble first, followed by the higher angular momentum components, resulting in ``inside-out'' build up of disk material \\citep{Larson:1976,white:1991,Munoz-Mateos:2007}. Because gas can efficiently dissipate energy, its motion around the galaxy is mostly circular. Stars born in such gas disks are initially on nearly circular orbits, but their infant kinematic state is highly fragile; non-axisymetric perturbations such as bars or spiral arms readily drive the orbits away from circular. Despite the increase in eccentricity, mean orbital radii are assumed to remain relatively constant. As a result, interpretations of Galactic and extra-galactic stellar population observations invariably make the fundamental assumption that stars remain at roughly their birth radii throughout their lives. Consequently, this assumption is built into the vast majority of models of formation and evolution of galactic disks over the past few decades \\citep[e.g.][]{Tinsley:1975, Matteucci:1989, Carigi:1996,Chiappini:1997,Boissier:1999}.  Recent theoretical and observational evidence challenges this static picture. Spirals have long been known to be an important source of kinematic heating in galactic disks, gradually increasing the eccentricities of stellar orbits \\citep{Jenkins:1990}. However, if a star is caught in the corotation resonance of a transient spiral it may move inward or outward in radius while preserving the circularity of its orbit \\citep[][hereafter SB02]{Sellwood:2002}. Transient spirals with a wide range of pattern speeds are present throughout the evolution of the disk. Different pattern speeds result in spatially distinct locations of corotation resonances, allowing a star to undergo a series of resonant encounters in its lifetime, riding the spiral waves across large portions of the galaxy while retaining a mostly circular orbit. Stars born in-situ may therefore remain kinematically indistinguishable from those that have come across the galaxy.  In \\citet{Roskar:2008} we presented first results from our $N$-body Smooth Particle Hydrodynamics (SPH) simulations of disk galaxy formation in which stars migrated radially due to their resonant interactions with transient spirals. These migrations yielded radial stellar population gradient predictions, which have recently been observed in resolved-star \\citep{de-Jong:2007}, and integrated-light studies \\citep{Azzollini:2008, Bakos:2008}. In conjunction with our models, these observational data strongly imply that radial migration is an important process in observed galactic disks. In this Letter, we present further analysis of the simulation presented in \\citet{Roskar:2008}, focusing on the wide-ranging implications of radial migrations for the observable properties of stellar populations. ",
        "conclusions": "We have shown that radial migration caused by resonant interactions of stars with transient spirals has significant repercussions on the observed properties of a wide range of stellar population systems. The range of implications and possible solutions to outstanding problems given by radial migration is appealing. Further work is required to characterize fully the degrees and timescales of migration in spiral galaxies of different physical characteristics. On the other hand, future observational efforts such as RAVE, GAIA and LSST will provide intricate data sets that should further constrain the radial migration process in the Milky Way. Ongoing and future HST projects such as the ANGST and GHOSTS surveys, will require radial migration models in order to reconstruct the SFH from resolved-star observations of nearby spiral disks and thereby learn about the formation history of our galactic neighborhood."
    },
    "0808/0808.2148.txt": {
        "abstract": "We  present follow-up spectroscopy  and photometry  of 11  post common envelope  binary  (PCEB)  candidates  identified from  multiple  Sloan Digital  Sky Survey (SDSS)  spectroscopy in  an earlier  paper. Radial velocity   measurements   using   the   \\Lines{Na}{I}{8183.27,8194.81} absorption  doublet  were performed  for  nine  of  these systems  and provided  measurements of  six orbital  periods in  the  range $\\Porb= 2.7-17.4$\\,h.  Three  PCEB candidates did not  show significant radial velocity  variations  in  the  follow-up  data,  and  we  discuss  the implications for the use of SDSS spectroscopy alone to identify PCEBs. Differential  photometry confirmed  one of  our  spectroscopic orbital periods  and  provided  one  additional \\Porb\\,  measurement.   Binary parameters  are estimated  for the  seven  objects for  which we  have measured the orbital  period and the radial velocity  amplitude of the low-mass companion star, $K_\\mathrm{sec}$.   So far, we have published nine SDSS PCEBs orbital periods,  all of them $\\Porb<1$\\,d. We perform Monte-Carlo simulations  and show that $3\\sigma$  SDSS radial velocity variations  should still  be  detectable for  systems  in the  orbital period  range  of  $\\Porb\\sim1-10$\\,days.  Consequently,  our  results suggest  that   the  number   of  PCEBs  decreases   considerably  for $\\Porb>1$\\,day, and that during  the common envelope phase the orbital energy of the binary star is maybe less efficiently used to expell the envelope than frequently assumed.  % We discuss this  result in the context of  possible selection effects %and observational biases  and conclude that the paucity  of PCEBs with %orbital periods longer than a day is probably real. ",
        "introduction": "Most of the stellar sources are  formed as parts of binary or multiple systems. Therefore  the study of  binary star evolution  represents an important part  of studying stellar  evolution. While the  majority of the wide  main sequence binaries evolve  as if they  were single stars and  never interact,  a small  fraction are  believed to  undergo mass transfer  interactions.   Once the  more  massive  main sequence  star becomes  a   red  giant  it  eventually  overfills   its  Roche  lobe. Dynamically  unstable  mass  transfer  exceeding the  Eddington  limit ensues onto the companion  star, which consequently also overfills its own Roche  lobe.  The  two stars then  orbit inside a  common envelope (CE,  e.g.   \\citealt{paczynski76-1, livio+soker88-1,  iben+livio93-1, taam+ricker06-1,  webbink07-1}),  and  friction inside  this  envelope causes a rapid  decrease of the binary separation.  Orbital energy and angular momentum are  extracted from the binary orbit  and lead to the ejection  of the  envelope,  exposing a  post  common envelope  binary (PCEB).  After the envelope is expelled PCEBs keep on evolving towards shorter   orbital   periods   through   angular  momentum   loss   via gravitational  radiation  and, for  companion  stars  above the  fully convective mass limit,  by magnetic wind braking. In  binaries that do not undergo a CE phase both components evolve almost like single stars and keep  their wide binary  separations.  A predicted  consequence is hence  a  strongly  bi-modal  binary  separation  and  orbital  period distribution   among   post-main   sequence   binaries,   with   PCEBs concentrated at  short orbital periods, and non-PCEBs  at long orbital periods.  \\begin{table*} \\label{t-logobs} \\caption{Log of  the observations. Included are the  target names, the SDSS  $ugriz$ psf-magnitudes,  the period  of the  observations, the telescope/instrument  setup, the  exposure time,  and the  number of exposures.} \\setlength{\\tabcolsep}{1.0ex} \\begin{tabular}{lcccccccccrlr} % \\hline \\hline \\textbf{Spectroscopy}    &                   &       &        &             &           &     \\\\ SDSS\\,J  & $u$ & $g$ & $r$ & $i$ & $z$ & Dates & Telesc. & Spectr. & Grating/Grism & Exp.[s] & \\# spec. \\\\ \\hline 005245.11--005337.2 &  20.47 & 19.86 & 19.15 & 17.97 & 17.22 & 16/08/07-20/08/07 & NTT        & EMMI   & Grat\\#7     & 1200      & 21  \\\\ 024642.55+004137.2  &  19.99 & 19.23 & 18.42 & 17.29 & 16.60 & 05/10/07-09/10/07 & NTT        & EMMI   & Grat\\#7     & 650-800   & 19  \\\\ 030904.82--010100.8 &  20.77 & 20.24 & 19.50 & 18.43 & 17.77 & 07/10/07-12/10/07 & NTT        & EMMI   & Grat\\#7     & 1500      & 4   \\\\ 031404.98--011136.6 &  20.78 & 19.88 & 19.03 & 17.76 & 16.98 & 02/10/07-03/10/07 & M-Baade    & IMACS  & 600 line/mm & 600-900   & 12  \\\\ 113800.35--001144.4 &  19.14 & 18.86 & 18.87 & 18.15 & 17.53 & 16/08/06-23/03/07 & VLT ($sm$) & FORS2  & 1028z       & 900       & 2   \\\\ &        &       &       &       &       & 18/06/07-23/06/07 & WHT        & ISIS   & 158R        & 1500      & 1   \\\\ &        &       &       &       &       & 17/05/07-20/05/07 & M-Clay     & LDSS-3 & VPH-red     & 500-600   & 5   \\\\ 115156.94--000725.4 &  18.64 & 18.12 & 18.14 & 17.82 & 17.31 & 17/05/07-20/05/07 & M-Clay     & LDSS-3 & VPH-red     & 450-600   & 17  \\\\ 152933.25+002031.2  &  18.69 & 18.20 & 18.33 & 17.98 & 17.48 & 18/05/07-20/05/07 & M-Clay     & LDSS-3 & VPH-red     & 500-1000  & 18  \\\\ 172406.14+562003.0  &  15.83 & 16.03 & 16.42 & 16.42 & 16.52 & 25/09/00-29/03/01 & SDSS       &        &             & 900       & 23  \\\\ 224139.02+002710.9  &  19.62 & 18.82 & 18.40 & 17.33 & 16.58 & 02/10/07-03/10/07 & M-Baade    & IMACS  & 600 line/mm & 600       & 2   \\\\ &        &       &       &       &       & 07/10/07-12/10/07 & NTT        & EMMI   & Grat\\#7     & 750-800   & 3   \\\\ 233928.35--002040.0 &  20.40 & 19.68 & 19.15 & 18.07 & 17.36 & 07/10/07-12/10/07 & NTT        & EMMI   & Grat\\#7     & 1400-1700 & 15  \\\\ \\hline \\hline \\textbf{Photometry}      &                   &       &        &             &           &     \\\\ SDSS\\,J  & $u$ & $g$ & $r$ & $i$ & $z$ & Dates & Telesc. &  & Filter band   &  Exp.[s]   & \\# hrs  \\\\ \\hline 031404.98--011136.6 &  20.78 & 19.88 & 19.03 & 17.76 & 16.98 & 19/09/06-20/09/06 & CA\\,2.2\\,m  & & clear & 35-60  & 10.3 \\\\ 082022.02+431411.0  &  15.92 & 15.85 & 16.11 & 15.83 & 15.38 & 21/11/06-24/11/06 & Kryoneri\\,1.2\\,m & & $R$ & 60   & 8.2 \\\\ 172406.14+562003.0  &  15.83 & 16.03 & 16.42 & 16.42 & 16.52 & 04/08/06-10/08/06 & IAC80       & & $I$   & 80-180 & 12.2 \\\\ &        &       &       &       &       & 25/03/06-13/09/06 & AIP 70 cm   & & $R$   & 60-90  & 55.6 \\\\ \\hline \\end{tabular} \\begin{minipage}{\\textwidth} Notes: M-Baade  and M-Clay refer  to the two Magellan  telescopes at Las Campanas observatory. We use $sm$ to indicate that the data were taken in service  mode. CA\\,2.2 is the 2.2  metre telescope at Calar Alto observatory. \\end{minipage} \\end{table*}  The scenario outlined  above is thought to be  a fundamental formation channel  for a  wide range  of astronomical  objects such  as low-mass X-ray binaries, double degenerate white dwarfs, neutron star binaries, cataclysmic  variables and  super-soft X-ray  sources.  Some  of these objects  will eventually  end their  lives as  type Ia  supernovae and short gamma-ray bursts, which are of great importance for cosmological studies.  While the  basic scenario of  CE evolution has  been known for  a long time \\citep[e.g.][]{paczynski76-1}, the  physical details involved are complex and still poorly  understood. As a result, full hydrodynamical models for the  CE phase have been calculated only  for a small number of cases \\citep{sandquistetal00-1,ricker+taam08-1}. Similarly, orbital evolution  through  magnetic  braking  is not  well  understood,  with different prescriptions  in angular momentum loss  differing by orders of   magnitudes  \\citep[e.g.][]{verbunt+zwaan81-1,  andronovetal03-1}. Consequently, current binary population synthesis models still have to rely  on  simple  parametrisations   of  energy  or  angular  momentum equations \\citep{nelemansetal00-1, dewi+tauris00-1, nelemans+tout05-1, politano+weiler06-1, politano+weiler07-1}.  A fundamental  problem in advancing our understanding  of CE evolution and magnetic  braking in close  binaries is the shortage  of stringent observational constraints that  can be used to test  and calibrate the theory.  PCEBs  consisting of a white  dwarf and a  main sequence star are  a very  well-suited  class  of objects  to  provide the  required innovative   observational  input  since   they  are   numerous,  well understood in terms  of their stellar components, they  are bright and hence accessible  with 2--8 metre  telescopes, and their study  is not complicated by mass transfer. \\citet{schreiber+gaensicke03-1} analysed a sample  of 30 well studied  PCEBs and showed that  the population of PCEBs known so  far is not only small but  also heavily biased towards young systems with low-mass secondary stars, and can hence not be used for         comparison          with         population         models \\citep[e.g.][]{willems+kolb04-1}.  \\begin{figure*} \\includegraphics[width=0.83\\textwidth]{fig1.ps} \\caption{\\label{f-spec} SDSS  spectra of the 11 SDSS  white dwarf main sequence binaries studied in this work.} \\end{figure*}  The    Sloan    Digital    Sky    Survey    \\citep[SDSS]{yorketal00-1, stoughtonetal02-1, adelman-mccarthyetal08-1} is currently dramatically increasing the number  of known white dwarf plus  main sequence (WDMS) binaries \\citep{silvestrietal07-1,  schreiberetal07-1}, paving the way for  large-scale  observational  PCEB  population  studies.   We  have initiated an  observational programme to identify the  PCEBs among the SDSS  WDMS   binaries,  and  to  determine   their  binary  parameters \\citep{rebassa-mansergasetal07-1,  schreiberetal08-1}.   Those systems in  which  significant  radial  velocity variation  is  detected,  are classified  as close  binaries  or PCEBs.  Here  we present  follow-up observations     of    11     PCEB     candidates    identified     in \\citet{rebassa-mansergasetal07-1} and  provide accurate values  of the orbital periods as  well as estimates of their  stellar parameters and orbital inclinations for seven of these systems. ",
        "conclusions": "We have presented a study of 11 PCEB previously identified candidates, and measured  the orbital periods  of six and  one of them  from their \\Ion{Na}{I}  doublet radial velocity  variations and  its differential photometry,  respectively.   Combining  the $K_\\mathrm{sec}$  velocity amplitudes with  the results from  spectroscopic decomposition/fitting of their  SDSS spectra,  we constrained the  binary parameters  of the seven systems for which we could determine orbital periods.  No radial velocity   variations  were  detected   for  three   PCEB  candidates, suggesting  that they  may be  wide binaries.   We revisited  the PCEB candidate selection of \\citet{rebassa-mansergasetal07-1}, and conclude that  candidates with  low amplitude  velocity variations,  noisy SDSS spectra, and radial velocity shifts that only show up in the H$\\alpha$ emission  line definitely need  additional follow-up  spectroscopy for the confirmation/rejection of their PCEB nature.  Finally, we have had a first look at the  period distribution of PCEBs from SDSS identified from radial velocity  snapshots of SDSS WDMS binaries,  and noted that none  of the  ten  systems published  so  far have  an orbital  period $>1$\\,d.   Using a  Monte-Carlo simulation,  we demonstrated  that our method  of finding  PCEBs should  be efficient  also for  systems with periods $>1$\\,d,  and that, subject  to small number statistics, it appears  that the PCEB period distribution peaks at $\\Porb<1$\\,d.  This is in agreement with the period distribution of the  previously known PCEBs,  which represents  a heterogenous  mix of systems identified by various  different methods. In contrast, current binary population models predict a  large number of PCEBs with orbital periods $>1$\\,day  in clear  disagreement with the  currently observed sample. Additional effort is needed  to improve the size of the sample of known PCEBs,  as well as to fully model  its selection effects, but with more than  1000 SDSS WDMS binaries to draw  from, the outlook for more quantitative tests of CE evolution seem promising."
    },
    "0808/0808.3944_arXiv.txt": {
        "abstract": "The magnetospheric plasma distribution in the vicinity of a pulsar at various inclination angles is investigated using a new relativistic, parallel 3D Particle-In-Cell (PIC) code DYMPHNA3D. DYMPHNA3D uses a superposition of the electromagnetic fields associated with a rotating magnetised conducting sphere in a vacuum (\\textit{the pulsar fields}) and the electromagnetic fields due to the presence of the magnetospheric plasma surrounding the pulsar (\\textit{the plasma fields}), as the total fields. The plasma is moved self-consistently through the magnetosphere using the PIC methodology. Our initial simulation results are presented here. These show similar solutions to those obtained from previous numerical simulations, which show the fundamental plasma distribution in the vicinity of an aligned rotating neutron star to consist of two polar domes and an equatorial torus of trapped non-neutral plasma of opposite sign. The aligned case being the case in which the inclination angle, $\\alpha$, between the magnetic dipole moment and the rotation axis of the star is zero. Furthermore, our code allows for off-axis simulations and we have found that this plasma distribution collapses into a `Quad-Lobe' charge-separated non-neutral magnetospheric plasma distribution in the case of an orthogonal rotator, i.e., the case in which the magnetic dipole moment is oriented at right angles to the rotation axis of the neutron star, with the plasma remaining trapped close to the stellar surface by the force-free ($\\textbf{E}\\cdot\\textbf{B}=0$) surfaces. We find that if initialised with a Goldreich-Julian type distribution, the system is seen to collapse rapidly into these stable `Dome-Torus' structures. ",
        "introduction": "The problem as to exactly how and where the observed emission from pulsars is generated has defied explanation for nearly four decades. There are many popular models, each with its own advantages and disadvantages, each with its dedicated supporters. These models differ primarily in terms of the location within the magnetosphere at which the observed emission is generated, i.e., the location of the particle acceleration region. The primary high-energy emission models at present are the polar cap model \\citep{HM1998}, the outer gap model \\citep{CHR1986}, the slot gap model \\citep{HM2004,HARDINGETAL2008}, and the two-pole caustic model \\citep{DR2003}. The polar cap model proposes that the particle acceleration region is located close to the stellar surface in the open field region over the pulsar's polar caps. The outer gap model assumes that the vacuum gaps within which the particle acceleration occurs form in the outer magnetosphere along null charge surfaces. The slot gap model contains elements of both of these models. This model is an extension of the original polar cap model, required to overcome some problems that the polar cap model faced in terms of difficulties in producing emission beams wide enough to correlate with the wide pulse profiles observed. By allowing the gap to extend upwards near the polar cap rim, asymptotically approaching the boundary between the open and closed magnetospheric regions, delimited by the last closed field lines, the slot gap model is capable of producing these wider pulse profiles. Last, but not least, is the two-pole caustic model which assumes that the vacuum gap, in which the particle acceleration occurs, is a thin gap that extends from each of the polar caps to the light cylinder and is confined to the surface of the last open magnetic field lines. \\par \\cite{GJ1969} determined that a magnetised neutron star would be surrounded `everywhere' by a charge density given by, \\begin{equation} \\label{gjchargedensity} \\rho_{GJ}=\\frac{1}{4\\pi}\\mathbf{\\nabla}\\mathbf\\cdot\\mathbf{E}=-\\frac{1}{2\\pi c}\\frac{\\mathbf{\\Omega}\\mathbf\\cdot\\mathbf{B}}{1-|\\frac{(\\mathbf{\\Omega}\\mathbf\\times\\mathbf{r})}{c}|}, \\end{equation} where $\\rho_{GJ}$ is the Goldreich-Julian charge density and $\\mathbf{\\Omega}$ is the rotational frequency of the star. This was done for the aligned case which is a simplification of the more complex inclined rotator problem. They assumed that plasma on open field lines (i.e., the field lines which do not close within the light cylinder distance) can stream to infinity, while plasma on closed field lines corotates with the neutron star. They also assumed that a wind would be generated at the light cylinder, attributable to centrifugal forces due to the rigid corotation of this magnetospheric plasma with the star. \\par This model advocates a totally filled magnetosphere (TFM), in which the plasma is pulled from the pulsar surface to fill the magnetosphere with a charge density given by Eq.~(\\ref{gjchargedensity}). There are some critical problems with this model to include the fact that if the plasma has to follow magnetic field lines, then it is impossible for the magnetospheric region surrounding an aligned rotator to fill with plasma. Also, at large distance from the stellar surface, positive particles would have to travel along magnetic field lines which have their feet firmly rooted at the polar regions of the star, where only negative particles exist. The assertion that there needs to be acceleration of particles from the polar caps to replace particles lost at large distances from the pulsar is unfounded and these issues have been well addressed by \\cite{MICHEL1982,MICHEL2004,MICHEL2005}. \\par In stark contrast to the scenario in which the pulsar is surrounded `everywhere' by the Goldreich-Julian charge density are the results of explicit computer simulations, first performed by \\cite{KPM1985A,KPM1985B}. The results of such simulations demonstrate that it is possible to construct stable self-consistent solutions involving charge-separated non-neutral plasma distributions in the form of two polar domes of trapped charged particles separated by a vacuum gap from an equatorial torus of trapped charged particles of the opposite sign, for the case of the aligned rotator, $\\alpha=0^{\\circ}$. \\par The basic premise behind these early simulations was to let the quantised surface charge on the neutron star follow the electric forces quasi-statically, using an iterative approach until the force-free condition ($\\mathbf{E}\\cdot\\mathbf{B}=0$) is satisfied everywhere. This is an $O(N^{2})$ problem, as such it is computationally intensive and consequently, the number of particles had to be kept reasonably small by choosing a large charge quantisation level. In addition, in order to take advantage of the inherent symmetry involved in the system, the charge elements were taken to be rings. \\par These simulations were later reworked by \\cite{SMT2001}. In this case, the particles were taken to be `super-particles', with each `super-particle' representing a finite density of discrete particles and carrying a charge that is very large in comparison to the electronic charge. As before, the code moves the `super-particles' iteratively to their equilibrium positions, i.e., $\\mathbf{E}\\cdot\\mathbf{B}=0$ positions based on the electromagnetic fields in the magnetosphere which are due to the neutron star and all of the other particles. \\par As mentioned, these simulation methodologies are N-body in nature and as a result, are restrictive in terms of the number of particles that can be used due to computational overheads. The solution invoked to overcome this limitation is to use a Particle-In-Cell (PIC) approach, as was done by \\cite{SA2002}. To perform their simulations, they used a modified version of the publicly available 3D electromagnetic relativistic PIC code `TRISTAN' \\citep{BUNEMAN1993}. Their approach used a superposition of the vacuum rotator fields, i.e., the electromagnetic fields associated with a rotating conducting magnetised sphere in a vacuum \\citep{DEUTSCH1955} and the plasma fields, i.e., the electromagnetic fields due to the presence of the magnetospheric plasma, as the total electromagnetic fields. Deutsch wrote down by inspection, the vacuum rotator field equations (providing little insight into their derivation) which require a significant amount of processing to convert them into functions of $r,\\:\\theta\\:\\&\\:\\phi$. An explicit derivation of these field equations was performed by \\cite{MICHELLI1999}, Eq.~(\\ref{DEUTSCH_EQNS}). \\begin{eqnarray} \\label{DEUTSCH_EQNS} &&B_{r}=2B_{0}\\frac{a^{3}}{r^{3}}\\left(\\cos\\alpha\\cos\\theta+\\sin\\alpha\\sin\\theta\\cos\\phi_{s}\\right) \\nonumber \\\\ &&B_{\\theta}=B_{0}\\frac{a^{3}}{r^{3}}\\left(\\cos\\alpha\\sin\\theta-\\sin\\alpha\\cos\\theta\\cos\\phi_{s}\\right) \\nonumber \\\\ &&B_{\\phi}=B_{0}\\frac{a^{3}}{r^{3}}\\sin\\alpha\\sin\\phi_{s} \\nonumber \\\\ &&\\lefteqn{E_{r}=\\Omega{a}B_{0}\\Biggl(\\frac{2}{3}\\cos\\alpha\\frac{a^{2}}{r^{2}}+\\cos\\alpha\\frac{a^{4}}{r^{4}}\\left(1-3\\cos^{2}\\theta\\right)} \\nonumber \\\\ &&\\hspace{4em}-3\\sin\\alpha\\frac{a^{4}}{r^{4}}\\sin\\theta\\cos\\theta\\cos\\phi_{s}\\Biggr) \\nonumber \\\\ &&\\lefteqn{E_{\\theta}=\\Omega{a}B_{0}\\Biggl[-2\\cos\\alpha\\frac{a^{4}}{r^{4}}\\sin\\theta\\cos\\theta} \\nonumber \\\\ &&\\hspace{4em}+\\sin\\alpha\\left(\\frac{a^{4}}{r^{4}}\\cos2\\theta-\\frac{a^{2}}{r^{2}}\\right)\\cos\\phi_{s}\\Biggr] \\nonumber \\\\ &&E_{\\phi}=\\Omega{a}B_{0}\\sin\\alpha\\left(\\frac{a^{2}}{r^{2}}-\\frac{a^{4}}{r^{4}}\\right)\\cos\\theta\\sin\\phi_{s}, \\end{eqnarray} \\vspace{1em} \\par where $B_{0}$ is the magnetic field strength at the equator, r is the radial distance from the star, $\\theta$ is the stellar latitude, $\\phi$ is the stellar longitude ($\\phi_{s}=\\phi-\\Omega t$), and $\\alpha$ is the inclination angle between the magnetic dipole moment and the rotation axis of the neutron star. \\par Our simulations also use `super-particles', each of which represents a finite density of discrete particles and carries a charge that is very large in comparison to the electronic charge. The aligned rotator scenario mentioned above is the case of maximal symmetry. It is a simplification of the significantly more complicated inclined rotator scenario, in which the magnetic dipole moment is oriented at an angle, $\\alpha$, to the rotation axis, i.e., an orthogonal rotator being the case of minimal symmetry. It has been argued \\citep{MICHEL2005} that focusing efforts on the aligned case is a futile endeavour as this case does not produce an active pulsar. Also, the physics of this scenario is so drastically different to that of the inclined/orthogonal case, that the aligned model is incapable of generating any real insight into the operation of pulsars. Instead it is contested that the inclination between the magnetic and rotational axes of the neutron star is an essential attribute of any (\\textit{active}) pulsar model \\citep{MICHELSMITH2001,MICHEL2005}. \\par These simulations have independently verified the original KPM simulation results and showed the magnetospheric charge distribution to consist of these two polar domes of trapped negative charge (\\textit{electrons}), with an equatorial torus of trapped positive charge (\\textit{protons, positrons or ions}), separated by a vacuum gap. It is worth noting that the fundamentals of this `Dome-Torus' magnetospheric plasma distribution have also been investigated and verified independently by \\cite{RYLOV1976, RYLOV1989, SHIBATA1989, NEUKIRCH1993, ZACHARIADES1993, THIELHEIM1994}, and \\cite{PETRIETAL2002} as indicated by \\cite{MICHEL2004}. ",
        "conclusions": "The initial test application of this newly developed code has yielded some interesting results. We have found from our investigations into the plasma distribution in the vicinity of a rapidly rotating neutron star, that stable solutions involving charge-separated non-neutral plasma structures are not only viable, but are the natural equilibrium states for these systems. Our results agree well with the results of prior numerical simulations based on the aligned rotator scenario and we have extended these to look at the `off-axis' cases as well. \\par The general topology of the Crab nebula, as observed in images from the Chandra X-ray observatory, is in some fashion reminiscent of the `Dome-Torus' type plasma distribution, evident in the observed equatorial toroidal outflow and polar jets. If the observed equatorial toroidal outflow and polar jets are in some way a reflection of this kind of plasma structure in the vicinity of the pulsar, then given that this distribution is seen to collapse into a `Quad-Lobe' type structure in the case of the orthogonal rotator, this may suggest that a nebula involving an orthogonal rotator or near orthogonal rotator, would not posses this observable macroscopic structure. This would also indicate that the inclination of the Crab pulsar should be appreciably lower than than that of an orthogonal rotator, in order that it still possess this macroscopic quasi-`Dome-Torus' structure. Chandra observations of the young Crab-like pulsar J0205+6449 in the SNR 3C 58 also clearly show a macroscopic torus structure (seen edge-on) and some jet-like structure. The inclination angle is estimated to be of the order of $70^{\\circ}$, \\cite{SLANEETAL2002,SLANEETAL2004}, slightly in excess of typical Crab estimates of around $65^{\\circ}$, \\cite{ZHANGCHENG2002}. \\par Given the success of this fledgling code in its initial application and its inherent modularity, these pave the way for application to a plethora of interesting astrophysical kinetic plasma phenomena. Future plans for the code involve continuing the initial investigations presented in this paper in greater detail to include probing the possibility and effects of the diocotron instability, \\citep{SA2002,PETRIETAL2002,PETRIETAL2003}. Also planned as part of this work is the investigation of possible mechanisms required to perturb the observed stable distribution in order to allow the plasma to propagate beyond these trapping regions to form the observed equatorial outflow and polar jets. The nature of the developed code will allow us to numerically investigate the pulsar magnetospheric plasma distribution in unprecedented detail both spatially and temporally. Investigations of possible pulsar GRP mechanisms, necessarily over very short spatial scales are planned as well as simulations of a mechanism capable of generating broadband persistent electron cyclotron maser emission from Brown Dwarfs. The developed code is at its most basic, a `first-principles' approach to modelling plasmas and has been implemented in an extremely modular fashion facilitating the possibility of inclusion a wide variety kinetic phenomena into the computational model."
    },
    "0808/0808.3669_arXiv.txt": {
        "abstract": "\\input{Roche_abstract} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.3950_arXiv.txt": {
        "abstract": "XTE~J1701--407 is a new transient X-ray source discovered on June 8th, 2008. More than one month later it showed a rare type of thermonuclear explosion: a long type I X-ray burst. We report herein the results of our study of the spectral and flux evolution during this burst, as well as the analysis of the outburst in which it took place. We find an upper limit on the distance to the source of 6.1 kpc by considering the maximum luminosity reached by the burst. We measure a total fluence of 3.5$\\times 10^{-6}$ erg/cm$^2$ throughout the $\\sim$20 minutes burst duration and a fluence of 2.6$\\times 10^{-3}$ erg/cm$^2$ during the first two months of the outburst. We show that the flux decay is best fitted by a power law (index $\\sim$1.6) along the tail of the burst. Finally, we discuss the implications of the long burst properties, and the presence of a second and shorter burst detected by {\\it Swift} ten days later, for the composition of the accreted material and the heating of the burning layer. ",
        "introduction": "\\label{sec:intro}  The bulge of our Galaxy harbours a large fraction of the known population of accreting compact objects. During one of the regular monitoring observations of that region with the proportional counter array (PCA) onboard the Rossi X-ray timing explorer ({\\it RXTE}) on June 8th, 2008, a new transient X-ray source was discovered: XTE~J1701--407 \\citep{Markwardt08}. Three days later a follow-up observation with the X-ray telescope (XRT) onboard {\\it Swift} improved the source localization and showed a spectrum that could be fitted with an absorbed power law of index $\\sim$2 \\citep{Degenaar08}. The transient nature, flux level and spectrum of the source are characteristic of X-ray binaries, but the exact class of the binary (high-mass or low-mass) was unkown and the compact object (black hole or neutron star) remained unidentified for more than one month. On July 17th at 13:29:59 UT the burst alert telescope (BAT) onboard {\\it Swift} detected an X-ray flare at a position consistent with that of XTE~J1701--407 \\citep{Barthelmy08}. {\\it Swift} began slewing to the source 60 seconds after the trigger and the XRT began observing the field 97 seconds after the initial BAT trigger, detecting a fading X-ray source at the position of XTE~J1701--407 \\citep{Barthelmy08}. \\citet{Markwardt08b} suggested that the flare could be caused by a thermonuclear burst based on the early BAT data. By studying the evolution of the XRT spectrum, \\citet{Linares08} found a clear cooling curve and confirmed the speculation that this was a thermonuclear event, identifying the source as an accreting neutron star (NS) and the system, in all likelihood, as a low-mass X-ray binary (LMXB). As noted by \\citet{Linares08} this particular event, having a total duration of more than 15 minutes, belongs to the rare subclass of long duration bursts \\citep{Cumming06}.  After the long X-ray burst the source continued in outburst\\footnote{To clarify the terminology we stress that in the context of LMXBs an outburst is powered by accretion and lasts for weeks to years whereas a type I X-ray burst, or ``burst'', results from unstable thermonuclear burning on the surface of a NS and lasts for $\\sim$10 seconds to $\\sim$hours.} and {\\it RXTE} kept observing it. A refined {\\it Swift}-XRT position \\citep{Starling08}, a search for an IR counterpart \\citep{Kaplan08} and the discovery of kHz quasi-periodic oscillations \\citep[QPOs;][]{Strohmayer08} were all reported during the following two weeks. On July 27th BAT detected another X-ray flare from XTE~J1701--407 \\citep{Sakamoto08}. The soft spectrum and duration ($\\sim$10~s) of the flare indicated that this was a type I X-ray burst (http://gcn.gsfc.nasa.gov/notices\\_s/318166/BA/), and hereafter we refer to it as the short burst (even though from the BAT data the total duration is somewhat uncertain). At the moment of writing the source is still active, more than two months after the start of the outburst.  In this Letter we present our analysis of the bursting behaviour of this source. Long bursts are rare, and therefore probe quite unusual burning regimes. Their duration is between that of normal Type I X-ray bursts (durations $\\sim 10-100$ s, triggered by unstable burning of H/He) and superbursts (durations $\\sim$ hours, triggered by unstable burning of C). There are two scenarios in which such bursts could occur, both requiring the build-up and ignition of a thick layer of He. One possibility is that the system is ultracompact, so that it accretes nearly pure He \\citep{Cumming06}. The other possibility is that unstable H burning at low accretion rates builds up He \\citep{Peng07, Cooper07}. While several examples of the first type are now known (see for example \\citealt{intZand07, Falanga08} and references therein), the second class - which probe the boundaries of H burning stability - have so far proved much rarer \\citep{Chenevez07}. In this Letter we will demonstrate that XTE J1701-407 is a good candidate for membership of this second class. ",
        "conclusions": "\\label{sec:discussion}   Theoretical modelling of long bursts suggests that they involve the build-up and ignition of a thick layer of He. One possibility is that the binary is ultracompact, so that the accreted fuel is almost pure He.  At low accretion rates, and in the absence of heating from steady H burning, the temperature remains low and ignition can be delayed until a thick layer of He has built up \\citep{intZand05, Cumming06}. Although He burning is quick, the cooling time is long because of the thickness of the layer in which the heat is deposited. Long bursts can also be triggered in systems accreting a mix of H and He if the source accretes at low rates, close to the point where H burning stabilises \\citep{Fujimoto81, Peng07, Cooper07}.  In this case weak H flashes can build up He, which ignites sporadically, resulting in a long burst.  Layer thickness again results in a long cooling time: but the presence of H in the burning mix may also prolong nuclear energy generation.   \\begin{figure} \\begin{center} \\resizebox{0.6\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics[]{./decay.ps}}} \\caption{Fit to the long burst flux decay with a power law of index $\\sim$1.6, together with the values of the baseline-subtracted bolometric flux.} \\label{fig:decay} \\end{center} \\end{figure}   Using the measured fluxes $F$ from Table \\ref{table:bursts} and the upper limit on the distance $d$ that we obtained from the long burst, we infer an upper limit on the global accretion rate $\\dot{M} \\approx 4\\pi d^2 R F/GM$ of 3.7$\\times 10^{-10}$ $M_\\odot$ yr$^{-1}$ (for a neutron star mass $M = 1.4 M_\\odot$ and radius $R = 10$ km) around the times of the two bursts. Assuming isotropy, this corresponds to an upper limit on the local accretion rate $\\dot{m}$ of 1.9$\\times 10^{3}$ g/cm$^2$/s, (2.5\\% of the Eddington rate for an Eddington luminosity of 1.6$\\times 10^{38}$ erg/s). This accretion rate is close to the rate where H burning is expected to transition to stability \\citep{Fujimoto81, Peng07, Cooper07}, so we need to examine both scenarios for the generation of the long burst.   The BAT rise time of the long burst, $\\approx 50$ s, is much longer than most of the pure He bursts examined by \\citet{intZand07} and \\citet{Falanga08}. This points to it being a mixed H/He type burst, triggered by H ignition, although we note that burst lightcurves can change considerably depending on the energy range where they are observed \\citep[see][]{Chelovekov06,Molkov05}.  We next consider the ignition conditions. The long burst had a total energy output of at most $E_b = 1.6\\times 10^{40}$ ergs (6.1 kpc distance).  The corresponding ignition column depth $y = E_b (1+z)/4\\pi R^2 Q_\\mathrm{nuc}$, where $z$ is the redshift, $R$ the neutron star radius and $Q_\\mathrm{nuc} = 1.6 + 4 X$ MeV/nucleon the nuclear energy release, given an H fraction X at ignition \\citep{Galloway06}. Assuming $z=0.31$ (the redshift for a 1.4 $M_\\odot$, 10 km radius neutron star) we derive an ignition depth $y = 3.9\\times 10^8$ g/cm$^2$ for solar abundances (X=0.7) and $y = 1.1\\times 10^9 $ g/cm$^2$ for pure He (X=0). Ignition of pure He at this column depth would require a heat flux from the deep crust of $\\gtrsim$ 2 MeV/nucleon at 1\\% Eddington accretion rates (Figure 22, \\citealt{Cumming06}). This is higher than the crustal energy release of $\\sim$ 1.5 MeV/nucleon calculated by \\citet{haensel1990,haensel2003}, but close to the value of $\\sim$ 1.9 MeV/nucleon recently found by \\citet{gupta2007,haensel2008}. Recent results by \\citet{horowitz2008} demonstrate that even larger amounts of heat may be released in the crust of an accreting neutron star. However, the high heat requirements certainly put some strong constraints on the pure He accretion scenario. Using the maximum local accretion rate calculated above (1.9$\\times 10^{3}$ g/cm$^2$/s), these ignition conditions predict minimum recurrence times $\\Delta t_\\mathrm{rec} = y(1+z)/\\dot{m}$ of $\\sim$10 days (X = 0) and $\\sim$3 days (X = 0.7). The data provide only weak constraints on the time without bursts that elapsed before the long burst (we estimate that the source was in the field of view of BAT for about 20\\% of the two months analysed herein), but these numbers are at least compatible with the time since the start of the outburst.   \\begin{table} \\center \\footnotesize \\caption{Properties of both type I X-ray bursts detected from XTE~J1701--407} \\begin{minipage}{\\textwidth} \\begin{tabular}{ l c c} \\hline\\hline & Long & Short \\\\ \\hline BAT rise time (s) & 50 & 5 \\\\ Duration\\footnote{From start of rise to 1\\% of peak flux.} (min) & 21 & ? \\\\ Peak flux\\footnote{Bolometric} ($10^{-8}$ erg/s/cm$^2$) & 8.4 & $\\sim$4.9  \\\\ Total fluence\\footnote{Including fluence in BAT \\citep{Markwardt08b} and \\\\ during XRT data gap.} ($10^{-6}$ erg/cm$^2$) & 3.5 & ? \\\\ Persistent flux\\footnote{Unabsorbed 1-50~keV} ($10^{-10}$ erg/s/cm$^2$) & 8.6$\\pm$0.9 & 9.8$\\pm$1.0 \\\\ \\hline\\hline \\end{tabular} \\end{minipage} \\label{table:bursts} \\end{table}   The presence of a short, weaker burst at similar accretion rates is perhaps the strongest piece of evidence pointing to mixed H/He fuel. The short burst has a relatively slow rise time, $\\approx 5$ s, similar to that expected for mixed H/He bursts. The fact that it has a lower peak flux than the long burst also implies some H content: if the system were a pure He accretor we would expect both bursts to reach similar peak fluxes since He bursts, whether long or short, are expected to exhibit photospheric radius expansion \\citep{Cumming04b}. In the H-triggered burst scenario it is possible to have bursts with quite different properties at very similar accretion rates, as H burning transitions to stability \\citep{Cooper07}.  XTE~J1701--407 increases the number of known NS-LMXBs in the interesting low accretion rate bursting regime. The balance of evidence suggests that it may be a good probe of thermonuclear burning near the boundary of stable H ignition. However, further observations, in particular better constraints on recurrence times and burst energetics, are required to confirm this picture. Identification of the companion star and determination of the orbital parameters may also be able to confirm whether or not the system is ultracompact."
    },
    "0808/0808.3161_arXiv.txt": {
        "abstract": "% The W40 complex is a nearby site of recent massive star formation composed of a dense molecular cloud adjacent to an HII region that contains an embedded OB star cluster.  The HII region is beginning to blister out and break free from its envelope of molecular gas, but our line of sight to the central stars is largely obscured by intervening dust.  Several bright OB stars in W40 - visible at optical, infrared, or cm wavelengths - are providing the ionizing flux that heats the HII region. The known stellar component of W40 is dominated by a small number of partly or fully embedded OB stars which have been studied at various wavelengths, but the lower mass stellar population remains largely unexamined. Despite its modest optical appearance, at 600~pc W40 is  one of the nearest massive star forming regions, and with a UV flux of about 1/10th of the Orion Nebula Cluster, this neglected region deserves detailed investigation. ",
        "introduction": "The star forming region known as W40 consists of three interrelated components. First, there is the cold molecular cloud, which is designated using Galactic coordinates as G28.8+3.5, following \\citet{gos1970}.  This dark cloud has an angular extent on the order of one degree, and is centered around a dense molecular core with a diameter of approximately $20\\arcmin$, identified as TGU~279-P7 in the recent extinction atlas of \\citet{dob2005}.  Adjacent to this molecular cloud is the large blister HII region denoted as W40 \\citep{wes1958}, and also labeled as S64 by \\citet{sha1959}, as RCW~174 by Rodgers, Campbell, \\& Whiteoak (1960), or LBN~90 in the Lynds (\\citeyear{lyn1965}) catalog of bright nebulae. The W40 HII region is centered on J2000 coordinates $18^h31^m29^s$, $-2\\deg05\\farcm4$, and has a diameter of $\\sim6\\arcmin$.  The space between the hot, ionized HII region and the cold molecular cloud is marked by a thin CII region and an accompanying neutral interface \\citep{val1987}.  \\begin{figure}[tbp] \\plotfiddle{w40_fig1.eps}{6.3cm}{0.0}{85.0}{85.0}{-185.0}{0.0} \\caption{The W40 region seen on (left) the red Digitized Sky Survey and (right) at JHK with 2MASS. The field is approximately 17$\\times$17 arcmin, corresponding to a little less than 3 pc on the side at 600~pc. North is up and east is left. Courtesy K.~Getman. } \\label{fig:dss_2mass} \\end{figure}  Lastly, the W40 region hosts an embedded stellar cluster that is dominated by three bright IR sources \\citep{zei1978}.  These bright OB stars are the primary excitation sources for the W40 HII region \\citep{smi1985} and show evidence for substantial circumstellar envelopes \\citep{val1994}.  The W40 cluster is heavily obscured along our line of sight by the surrounding molecular cloud, which provides $A_V \\sim10$ magnitudes of visual extinction throughout, and as much as $A_V = 17$ mag at the dense center \\citep{rey2002}. Figure \\ref{fig:dss_2mass}  shows the W40 region in the optical from the Digitized Sky Survey and in the infrared from the 2MASS survey. ",
        "conclusions": ""
    },
    "0808/0808.2982.txt": {
        "abstract": "We present results from an adaptive optics survey for substellar and stellar companions to Sun-like stars.  The survey targeted 266 F5--K5 stars in the 3~Myr to 3~Gyr age range with distances of 10--190~pc.  Results from the survey include the discovery of two brown dwarf companions (HD~49197B and HD~203030B), 24 new stellar binaries, and a triple system.  We infer that the frequency of 0.012--0.072~$\\Msun$ brown dwarfs in 28--1590~AU orbits around young solar analogs is $3.2^{+3.1}_{-2.7}$\\% (2$\\sigma$ limits).  The result demonstrates that the deficiency of substellar companions at wide orbital separations from Sun-like stars is less pronounced than in the radial velocity ``brown dwarf desert.'' We infer that the mass distribution of companions in 28--1590~AU orbits around solar-mass stars follows a continuous $dN/dM_2 \\propto M_2^{-0.4}$ relation over the 0.01--1.0~\\Msun\\ secondary mass range.  While this functional form is similar to the that for $<$0.1~\\Msun\\ isolated objects, over the entire 0.01--1.0~\\Msun\\ range the mass functions of companions and of isolated objects differ significantly.  Based on this conclusion and on similar results from other direct imaging and radial velocity companion surveys in the literature, we argue that the companion mass function follows the same universal form over the entire range between 0--1590~AU in orbital semi-major axis and $\\approx$0.01--20~\\Msun\\ in companion mass.  In this context, the relative dearth of substellar versus stellar secondaries at {\\it all} orbital separations arises naturally from the inferred form of the companion mass function. ",
        "introduction": "The properties of brown dwarf companions to stars are important for understanding the substellar companion mass function (CMF), the formation of brown dwarfs, and the formation and evolution of low-mass ratio binary systems.  Widely-separated brown dwarf companions, in particular, are an important benchmark for studying the properties of substellar objects because of their accessibility to direct spectroscopic characterization and their relative ease of age-dating---from assumed co-evality with their host stars.  However, brown dwarf companions have been an elusive target for direct imaging. The main challenge has been the need to attain sufficient imaging contrast to detect secondaries that are $>$10$^3$ fainter than their host stars at angular separations spanning solar system-like scales ($<$40~AU~$=0\\farcs4$ at 100~pc).  The problem is alleviated at young ages when brown dwarfs are brighter.  In addition, nearby stars offer an additional advantage because the relevant angular scales are correspondingly wider and more accessible to direct imaging.  Young nearby stars are thus the preferred targets for substellar companion searches through direct imaging.  Nevertheless, early surveys for substellar companions, performed with seeing-limited or first-generation high-contrast imaging technology \\citep{oppenheimer_etal01, hinz_etal02, mccarthy_zuckerman04} had very low detection rates, $\\lesssim0.5$\\%.  This low brown dwarf companion detection rate was similar to that inferred from precision radial velocity surveys \\citep[$<0.5\\%$ over 0--3~AU;][]{marcy_butler00}, and prompted \\citet{mccarthy_zuckerman04} to conclude that the so-called ``brown dwarf desert'' extends far beyond the orbital separations probed by radial velocity surveys, out to at least $\\approx$1200~AU.  Over the past few years, advances in adaptive optics (AO) technology and high-contrast imaging methods have improved the chances for the direct imaging of substellar secondaries.  Modern AO systems, with $>$200 corrective elements spread across the beam of a 5--10~m telescope, are able to deliver high order rectification ($<$250~nm r.m.s.\\ residual error) of wavefronts perturbed by Earth's turbulent atmosphere at up to kHz rates.  In addition, our empirical appreciation of the local young stellar population has improved over the past decade, as demonstrated by the recent discoveries of a large number of young ($<500$~Myr) stellar associations within 200~pc from the Sun \\citep[e.g.,][and references therein]{kastner_etal97, mamajek_etal99, zuckerman_webb00, zuckerman_etal01a, montes_etal01a, zuckerman_song04}.  These have allowed us to select more suitable targets for direct imaging searches for substellar companions.  Several recent direct imaging surveys of nearby young stellar associations conducted on high-order AO-equipped telescopes (\\citealt{neuhauser_guenther04}, 25 A--M stars; \\citealt{chauvin_etal05a, chauvin_etal05b}, 50 A--M stars) or with the {\\it Hubble Space Telescope} ({\\it HST}; \\citealt{lowrance_etal05}, 45 A--M stars) have enjoyed higher detection rates (2--4~\\%) than the first generation of surveys.  In addition, at very wide ($>$1000~AU) separations, where the detection of brown dwarf companions to solar-neighborhood stars is not hindered by contrast, \\citet{gizis_etal01} have found that the frequency of substellar companions to F--K dwarfs is fully consistent with that of stellar companions to G dwarfs \\citep{duquennoy_mayor91}.  Thus, while the radial velocity ``brown dwarf desert'' remains nearly void within 3~AU even after the discovery of numerous extra-solar planets over the past decade, brown dwarf secondaries at $>$100--1000~AU separations seem to not be as rare.  The precise frequency of substellar companions in direct imaging surveys remains controversial. Several highly sensitive surveys performed with the {\\it HST} (\\citealt{schroeder_etal00}, 23 A--M stars; \\citealt{brandner_etal00}, 28 G--M stars; \\citealt{luhman_etal05e}, 150 B--M stars) and with high-order AO (\\citealt{masciadri_etal05}, 28 G--M stars; \\citealt{biller_etal07}, 54 A--M stars) have reported null detections of substellar companions.  Given the low (few per cent) detection rate of substellar companions in direct imaging surveys, it is now clear that, with $<50$ targets per sample, some of these surveys were too small to expect to detect even a single brown dwarf companion.  However, the non-detection of substellar secondaries in two largest surveys \\citep{luhman_etal05e, biller_etal07} is potentially significant.  Given current understanding of the importance of stellar mass for (1) stellar multiplicity rates \\citep[see review in][]{sterzik_durisen04} and (2) binary mass ratio distributions \\citep[see review in][]{burgasser_etal07}, it is imperative that any study of the substellar companion frequency is considered in the context of the mass distribution of primary stars in the sample.  Indeed, a large survey sample comprising primaries with identical masses is ideal.  The problem of the brown dwarf companion frequency is perhaps most comprehensively dealt with in the context of solar mass primaries.  For these a uniquely large body of stellar and substellar multiplicity data exist on all orbital scales.  On one hand, the exhaustive spectroscopic and imaging study of G dwarf multiples by \\citet{duquennoy_mayor91} provides an important anchor to the properties of 0.1--1.0~\\Msun\\ stellar companions to Sun-like stars.  On the other hand, the results from more than a decade of precision radial velocity surveys for planets around G and K stars allow a comparison with the planetary-mass end of the substellar companion mass range.  A large uniform sample of young Sun-like stars has been compiled by the Formation and Evolution of Planetary Systems (FEPS) {\\it Spitzer} Legacy team.  The purpose of the FEPS Legacy campaign with {\\sl Spitzer} was to study circumstellar disk evolution in the mid-IR.  However, the sample is also well-suited for a high-contrast imaging survey for substellar companions.  Seventy percent of the FEPS stars are younger than $\\sim$500~Myr, and all are within 200~pc.  As an auxiliary component to the FEPS program, we imaged most %(95\\%) %(228 stars) of the northern FEPS sample with the high-order AO systems on the Palomar 5~m and the Keck 10~m telescopes.  We further expanded our AO survey by observing several dozen additional nearby and mostly young solar analogs.  Preliminary results from the project were published in \\citet{metchev_hillenbrand04} and in \\citet{metchev_hillenbrand06}, including the discoveries of two brown dwarf companions: HD~49197B and HD~203030B.  The survey has now been completed, and no further brown dwarf companions have been found.  The results were analyzed in \\citet{metchev06}.  Here we present the AO survey in its entirety and focus on the statistical interpretation of the data.  The paper is organized as follows.  A full description of the survey sample is given in \\S~\\ref{sec_target_sample}.  The Palomar and Keck AO observing campaigns and the data reduction and calibration techniques are described in \\S~\\ref{sec_obsredcal}.  The candidate companion detection approach and the survey detection limits are discussed in \\S~\\ref{sec_detlim}.  The various methods used for bona fide companion confirmation are presented in \\S~\\ref{sec_candidates}.  Section~\\ref{sec_results} summarizes the results from our survey, including all of the newly-discovered and confirmed substellar and stellar secondaries.   Section~\\ref{sec_incompl_bias} contains a brief discussion of the various sources of incompleteness and a full discussion of the biases in the survey.  (A full-fledged incompleteness analysis is presented in the Appendix.) %(\\S~\\ref{app}). In \\S~\\ref{sec_anal} we estimate the frequency of wide substellar companions to young solar analogs, and present evidence for trends in the companion mass and companion frequency with semi-major axis and primary mass.  In \\S~\\ref{sec_discussion} we consider the results of the current investigation in the broader context of stellar multiplicity, and suggest the existence of a universal CMF.  Section~\\ref{sec_conclusion} summarizes the findings from our study. ",
        "conclusions": "\\label{sec_conclusion}  We have presented the complete results from a direct imaging survey for substellar and stellar companions to 266 Sun-like stars performed with the Palomar and Keck AO systems.  We discovered two brown dwarf companions in a sub-sample of 100 3--500~Myr-old stars imaged in deep coronagraphic observations.  Both were already published in \\citet{metchev_hillenbrand04, metchev_hillenbrand06}.  In addition, we discovered 24 % was 24 new stellar companions to the stars in the broader sample, five of which are in very low mass ratio $q\\sim0.1$ systems.  %Separately, we established the physical association of a bona fide companion at the stellar/substellar boundary to the Pleiad HII~1348.  Following a detailed consideration of the completeness of our survey, we found that the frequency of 0.012--0.072~\\Msun\\ brown dwarf companions in 28--1590~AU orbits around 3--500~Myr-old Sun-like stars is $3.2^{+3.1}_{-2.7}$\\% (2$\\sigma$ limits).  This frequency is %consistent with the frequency of extrasolar giant planets in 0--3~AU orbits, is marginally higher than the frequency of 0--3~AU radial velocity brown dwarfs, and is significantly lower than the frequency of stellar companions in 28--1590~AU orbits.  The frequency of wide substellar companions is consistent with the frequency of extrasolar giant planets in 0--3~AU orbits. %Therefore, we conclude that wide substellar companions are less common compared to wide stellar secondaries, although the deficiency may not be as stark as in the radial velocity brown dwarf desert. A comparison with other direct imaging surveys shows that substellar companions are most commonly detected at $\\gtrsim150$~AU projected separations from $\\gtrsim0.7~\\Msun$ stars.  However, because of bias against the direct imaging of faint close-in companions, brown dwarf secondaries are likely  also common at smaller projected separations.  Considering the two detected brown dwarf companions as an integral part of the broader spectrum of stellar and substellar companions found in our survey, we infer that the mass ratio distribution of 28--1590~AU binaries, and hence, the MF of 28--1590~AU secondary companions to solar-mass primaries, follows a $d N/d M_2\\propto M_2^{\\beta}$ power law, with $\\beta=-0.39\\pm0.36$ %$\\beta=-0.3\\pm0.3$ (1$\\sigma$ limits).  This distribution differs significantly from the MF of isolated objects in the field and in young stellar associations, and is inconsistent with random pairing of individual stars with masses drawn from the IMF.  In this context, the observed deficiency of substellar relative to stellar companions at wide separations arises as a natural consequence of the shape of the CMF, and does not require explanation through formation or evolutionary scenarios specific to the substellar or low-mass stellar regime.  Comparing our CMF analysis to results from other direct imaging and radial velocity surveys for stellar and substellar companions, we find tentative evidence for universal behavior of the CMF across the entire 0--1590~AU orbital semi-major axis and the entire 0.01--20~\\Msun\\ companion mass range.  Such a universal CMF is not inconsistent with the marked dearth of brown dwarfs in the radial velocity brown dwarf desert around Sun-like stars.  That is, the properties of brown dwarf companions at any orbital separation are conceivably an extension of the properties of stellar secondaries.  Hence, we predict that the peak in semi-major axes of brown dwarf companions to solar-mass stars occurs at $\\approx30$~AU.  Extrapolating the inferred CMF to masses below the deuterium burning limit, we find that if 0.003--0.012~\\Msun\\ ``planetary-mass'' secondaries can form through gravo-turbulent fragmentation, they should exist in $\\geq30$~AU orbits only around less than 1\\% of Sun-like stars."
    },
    "0808/0808.2810_arXiv.txt": {
        "abstract": "Due to the short settling times of metals in DA white dwarf atmospheres, any white dwarfs with photospheric metals must be actively accreting. It is therefore natural to expect that the metals may not be deposited uniformly on the surface of the star. We present calculations showing how the temperature variations associated with white dwarf pulsations lead to an observable diagnostic of the surface metal distribution, and we show what constraints current data sets are able to provide. We also investigate the effect that time-variable accretion has on the metal abundances of different species, and we show how this can lead to constraints on the gravitational settling times. ",
        "introduction": "There are two main classes of white dwarf stars: those with hydrogen-rich atmospheres (spectral type DA) and those with helium-rich atmospheres (non-DA spectral types). The reason for this is that the high surface gravities of white dwarfs lead to efficient gravitational settling, with the lightest elements rising to the surface. In addition, some DA white dwarfs have spectra showing metal lines of elements such as Ca and Mg \\citep{Zuckerman03}, and these stars are referred to as DAZs; about 20\\% of all DAs fall into this category. Recently, \\citet{Dufour07} have announced a new class of white dwarf with carbon-dominated atmospheres, the ``hot DQ'' stars, several examples of which have been found in the Sloan Digital Sky Survey \\citep{Liebert03}.     The presence of metals in the DAZs is intriguing since the settling time scale for the metals may be many orders of magnitude shorter than the evolutionary age of these objects.  Indeed, for DAZs with $\\teff \\sim$~12,000~K, the settling time scale can be on the order of days or weeks, meaning that these objects are experiencing ongoing accretion \\citep{Koester06}.  This ongoing accretion is consistent with the fact that nearly a dozen of these objects have detected dust disks \\citep[``debris disks,'' see][]{Tokunaga90,Becklin05,Kilic05,Reach05,Kilic06,vonHippel07a,Jura07a,Jura07b,Farihi08} and these disks are assumed to be the sources of the metal lines seen in these white dwarf atmospheres.  The best studied object of this DAZ class with an observed disk, G29-38, is also a multi-periodic variable white dwarf (DAV), pulsating in non-radial g-modes with periods of a few hundred to one thousand seconds.  The technique of asteroseismology uses the observed pulsation modes of a star to infer and constrain the interior structure of the star, thus obtaining information on the star's structure as a function of radius \\citep{Bradley94a,Kawaler94,Metcalfe02,Montgomery03}.  In contrast to this, our approach in this paper is to use the different \\emph{angular} dependence of the pulsations to constrain the accretion process. Since the accretion is most likely occurring through a disk, it is natural to suppose that the metals may not be uniformly distributed across the star's surface. We present calculations showing how we can place constraints on the non-uniformity of the accretion process. ",
        "conclusions": "We have shown how the temperature variations due to stellar pulsation can be used to constrain the metal distribution on the surface of a white dwarf, and we have shown that data currently in hand for the star G29-38 suggests that the metal distribution, and therefore the accretion, may be equatorial. If further studies support an equatorial distribution for Ca in G29-38, this argues against magnetic accretion onto spots near the poles.  Non-magnetic equatorial accretion would further imply that the inner edge of the disk is much thinner than the white dwarf's radius.  A physically thin disk is consistent with observations to date \\citep[e.g.,][]{Jura07a,vonHippel07a}.  We have also shown how the observed variations in atmospheric abundance of two chemical species can be used to place constraints on their settling rates. The element with the faster settling rate will show the larger fluctuations, and, at least within the context of this simplified model, the ratio of the magnitude of these excursions from the mean will be approximately proportional to the settling rate.  Thus, the time variability of these systems, both in terms of the pulsations of the white dwarf and in terms of the variable accretion rate, allows us to probe the physics of accretion and gravitational settling in these objects."
    },
    "0808/0808.0481_arXiv.txt": {
        "abstract": "We analyze the active-galaxy correlation reported in 2007 by the Pierre Auger Collaboration. The signal diminishes if the correlation-function approach (counting all ``source--event'' pairs and not only ``nearest neighbours'') is used, suggesting that the correlation may reveal individual sources and not their population. We analyze available data on physical conditions in these individual correlated sources and conclude that acceleration of protons to the observed energies is hardly possible in any of these galaxies, while heavier nuclei would be deflected by the Galactic magnetic field thus spoiling the correlation. Our results question the Auger interpretation of the reported anisotropy signal but do not contradict to its explanation with intermediate-mass nuclei accelerated in Cen~A. ",
        "introduction": "\\label{sec:intro} Astronomy of ultra-high-energy (UHE; energy $E \\gtrsim 10^{19}$~eV) cosmic rays (CRs) is a new branch of science whose birth we are currently witnessing. To identify UHECR sources is therefore an important and challenging task. Observational evidence in favour of particular sources is limited by low statistics, by systematic uncertainties in determination of the primary-particle type and energy, by deflection of charged particles in unknown cosmic magnetic fields and by other confusing effects in the analysis. Theoretically, several compelling mechanisms of UHECR acceleration have been developed and a number of potential astrophysical sources have been suggested (see e.g.\\ Refs.~\\cite{sources,comparative} for reviews and summary). Physical conditions in potential sources are often uncertain thus preventing one from firm identification of theoretically most favoured candidates.  Recently, considerable interest has been attracted by a claim of correlation of positions of nearby active galaxies with arrival directions of UHECR particles detected by the Pierre Auger Observatory (PAO)~\\cite{PAOagn,PAOagnLong}. The correlation was interpreted as an evidence that UHECR particles are protons either from nearby active galactic nuclei (AGN) or from other sources with a similar distribution in space. The AGN interpretation was considered likely based on the previous works favouring AGN as possible UHECR accelerators. Notably, Ref.~\\cite{PAOagnLong} cites Ref.~\\cite{BiermanStr} in this context; the latter reference however studies the most powerful AGN (BL Lac type objects and optically violent variable quasars), while the analysis of Ref.~\\cite{PAOagn} is based on the nearby objects listed in the V\\'eron-Cetti\\& V\\'eron catalog~\\cite{Veron} which are mostly low-power Seyfert galaxies. In Ref.~\\cite{paper1}, we analyzed general constraints from geometry of the source and from radiation losses which restrict potential astrophysical accelerators and compared them to the most recent astrophysical data. Our study suggested that only the most powerful AGN are able to accelerate UHE particles while Seyfert galaxies {\\em generically} are not. In this work, we apply the general constraints discussed in Ref.~\\cite{paper1} to {\\em particular} active galaxies correlated with Auger events.  The rest of the paper is organised as follows. In Sec.~\\ref{sec:PAO:corr}, we briefly review the Auger result emphasising particular details of the analysis of Refs.~\\cite{PAOagn,PAOagnLong} and demonstrate that the correlation signal is strong only if a single nearest-neighbour potential source is counted for each event and diminishes in the correlation-function approach when every pair ``event--source'' is counted both in the data and in the simulations. We then select these nearest-neighbour source candidates (mostly Seyfert galaxies) for the Auger events and study, in Sec.~\\ref{sec:PAO:individual}, their properties. We conclude that it is unlikely that these objects can accelerate protons up to the observed energies and discuss implications of this conclusion in Sec.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl} The correlations observed by PAO may be explained either by the AGN hypothesis or by chance coincidence. The AGN hypothesis implies that the primary cosmic-ray particles are protons and has three options:  (i) the correlated AGN are (some of) the sources;  (ii) the AGN as a whole are sources; particular correlated AGN may not be the actual sources if they are located in regions of enhanced AGN density;  (iii) the sources are not AGN but follow a similar distribution in space, that is, like AGN, they follow the local large-scale structure of the Universe.  The options (iii) and (ii) are disfavoured by the absence of the events from the direction of the Virgo supercluster~\\cite{Comment}; the options (i) and (ii) are disfavoured by the luminosity argument \\cite{Dedenko}.  In this paper we have noted that (ii) and (iii) are also disfavoured by the disappearence of correlations in terms of the correlation-function approach, contrary to the nearest-neighbour probability estimate. In order to test the option (i), we have studied physical conditions in particular correlated galaxies. Most of them are low-power Seyfert galaxies; we have demonstrated that {\\bf none of the correlated AGN can accelerate protons to the observed energies}, thus excluding the option (i) and giving further support to the explanation of correlations by means different than the AGN hypothesis.  Our results suggest that the correlated galaxies, notably Cen A, can accelerate nuclei to the observed energies. The deflections would however spoil the small-angle correlations, but the origin of a significant part of the detected particles in Cen~A jet and lobes remains one of the most probable explanations for the observed anisotropy."
    },
    "0808/0808.0354_arXiv.txt": {
        "abstract": "{New CCD photometry of NGC~6366 has lead to the discovery of some variable stars. Two possible Anomalous Cepheids (or Pop II Cepheids), three long period variables, one SX Phe and one eclipsing binary have been found. Also a list of 10 candidate variables is reported. The light curve of the RRab star, V1, has been decomposed into its Fourier harmonics, and the Fourier parameters were used to estimate the star's metallicity and distance; [Fe/H]~$= -0.87 \\pm 0.14$ and $d~=~3.2 \\pm 0.1$ kpc. It is argued that V1 may not be a member of the cluster but rather a more distant object. If this is so, an upper limit for the distance to the cluster of $2.8 \\pm 0.1$ kpc can be estimated. The $P-L$ relationship for SX Phe stars and the identified modes in the newly discovered SX Phe variable, V6, allow yet another  independent determination of the distance to the cluster of $d~=~2.7 \\pm 0.1$ kpc. The $M_V$- {\\rm [Fe/H]} relationship for RR Lyrae stars is addressed and the case of V1 is discussed.} ",
        "introduction": "\\label{sec:intro}  The importance of studying globular clusters is linked to the fact that they can provide insight into the structure and evolution of the Galaxy. Hence, the determination of their ages, chemical compositions, galactocentric distances and kinematics are of fundamental relevance.  Numerous works in recent literature  have been devoted to the estimation of the above mentioned quantities. Calculations of relative ages at a fixed metallicity provide an age scale with an uncertainty of $\\leq$ 1 Gyr. This has allowed Harris et al. (1997) to estimate that the very metal-poor ([Fe/H] $\\sim -2.14$) distant ($\\sim 90$ kpc) cluster NGC 2419 has similar age to the inner cluster M92, and conclude that the earliest stellar globular cluster formation began at about the same time everywhere in the Galaxy. By the same technique Stetson et al. (1999) found that some of the outer-halo clusters ($R_{GC} \\ge$ 50 kpc) (Palomar 3, Palomar 4 and Eridanus) are $\\sim$1.5-2 Gyr younger than the inner-halo clusters, ($R_{GC} \\leq$ 10 kpc) of similar metallicity (M3 and M5). These author recognise however that the age difference can be less than $\\sim$ 1 Gyr if their [Fe/H] or [$\\alpha$/H] abundances are overestimated. Similarly Sarajedini (1997) concluded that the outer-halo cluster Palomar 14 is 3 to 4 Gyr younger than inner clusters of similar metallicity. It seems that most of the outer-halo globular clusters are 1.5-4 Gyr younger than the inner-halo clusters of similar [Fe/H], and that NGC 2419 may be an exception (VandenBerg 2000).  Based on their metallicities and kinematics, two populations of globular clusters have  been distinguished; a halo population composed of metal poor clusters (i.e. [Fe/H] $\\leq -0.8$) which is slowly rotating ($V_{rot} = 20 \\pm 29 $ km/s with a large line-of-sight velocity dispersion of $127 \\pm 11$ km/s; Zinn 1996), and a disk population composed of metal rich globular clusters (i.e. [Fe/H] $\\ge -0.8$) which is rapidly rotating  ($V_{rot} = 157 \\pm 26 $ km/s with a small line-of-sight velocity dispersion of $66 \\pm 13$ km/s; Zinn 1996)  (see also Zinn 1985, Armandroff 1989; Zinn 1993; van den Bergh 1993). The metal-rich component may also include a bulge system ($R_{GC} \\leq$ 3 kpc) which according to Minniti (1995) formed after the halo. However, comparing the luminosity of the HB and the MS turn-off of the galactic bulge, the two metal-rich bulge globular clusters NGC~6528 and NGC~6553, and the inner Halo cluster NGC~104 (47 Tuc), allowed Ortolani et al. (1995) and Zoccali et al. (2003) to conclude that the bulge and the Halo are of the same age and that there are no traces in the bulge of an intermediate-age population. These authors recognise however that an age difference of $\\sim$ 2-3 Gyr in either direction is possible.   The clusters of intermediate metallicity ($\\sim~ -0.8$~dex) are very important since they can be used to trace the border between halo and disk. In this context, NGC 6366 (R.A.(2000)$=17^h27^m44^s$.3, DEC(2000)$=-05^{\\circ} 04'36''$ ; l=$18.41^{\\circ}$, b=$+16.04^{\\circ}$) is an interesting globular cluster since, on one hand, it is metal rich (with several [Fe/H] determinations that range between -0.65 and -0.99 Pike 1976, Johnson et al. 1982, Zinn \\& West 1984, DaCosta \\& Seitzer 1989, Da Costa \\& Armandroff 1995); and, on the other hand, it has a large heliocentric velocity of $-$123.2 $\\pm$ 1.0 km/s, which associates the cluster with the slowly rotating halo cluster system (Da Costa \\& Seitzer 1989), and which would make NGC 6366 the most metal-rich member of this cluster population. Its Galactic position near to the Galactic disk contributes to its high reddening, with estimates ranging between $E(B-V)= 0.65$ and $0.80$ (Harris 1976, Da Costa \\& Seitzer 1989, Harris 1993). Other Halo clusters with high metallicities  ([Fe/H] $\\geq -1.0$) in the list of Armandroff (1989) are NGC~6171, NGC~6569 and NGC~6712. An analysis of their kinematics led Cudworth (1988) and Da Costa \\& Seitzer (1989) to conclude that NGC~6712 is a halo cluster, while the other two cannot be classified unambiguously. Thus, it is likely that NGC~6712 is the second most metal-rich ([Fe/H] = $-1.01$) cluster of the halo population.  Photometric studies of globular clusters are of special interest because they can be used to generate a color magnitude diagram (CMD) that provides insight into the cluster age and evolutionary stage. The CMD of NGC 6366 displays a clear turn off point, a well developed red giant branch, and a very incipient horizontal branch (HB) (Pike 1976; Harris 1993, Alonso et al. 1997). The appearance of the CMD has triggered an interest in the age determination of the cluster and highlighted its importance in the study of galactic dynamics.  A few estimates of the age of NGC~6366 can be found in the literature, and their dispersion confirms the difficulty of dating globular clusters. Alonso et al. (1997) found an age of 18$^{+2}_{-3}$ Gyrs by comparing the CMD to the isochrones of Straniero \\& Chieffi (1991), and concluded that it is older by 4-6 Gyrs than other metal rich clusters. However, Rosenberg et al. (1999), using a homogeneous data base of 34 globular clusters and a differential approach relative to a group of coeval clusters of 13.2 Gyrs of age (Carretta et al. 2000), estimated an age of about 11.0 Gyrs and hence concluding that the cluster belongs to a group of clusters younger than the coeval group. Salaris \\& Weiss (2002), also applied a differential approach relative to a group of clusters whose ages are believed to be well determined (M15, M3, NGC 6171 and 47 Tuc) and made an age estimate of 9.4-9.6 $\\pm$ 1.4 Gyr for NGC~6366, depending on the metallicity adopted, i.e. $\\sim$ 2 Gyr younger than the reference clusters. This result further supports the idea that the cluster is young and it belongs to a group of young clusters linked to the galactic disk.  The lack of a developed HB in the CMD of NGC~6366 is the cause of the lack of RR Lyrae stars in the cluster, and the existence of a solitary known RR Lyrae variable in the cluster is therefore not surprising. This variable star was detected and studied by Sawyer Hogg (1973), and subsequently by Pike (1976) and Harris (1993). The latter aimed his $BV$ photometry with CCD observations to search for blue stragglers and variable stars. He reported 27 blue straggler candidates but found no variable stars. However, the photometry method used (PSF fitting directly to the images) does not allow precise measurements in crowded fields, limiting the possibility for detecting smaller amplitude variables (Harris 1993).  Difference image analysis (DIA) is a powerful technique allowing accurate PSF photometry of CCD images, even in very crowded fields (Alard \\& Lupton 1998; Alard 2000; Bramich 2008). In the present study, we apply a new algorithm for difference image analysis (Bramich 2008) to a set of $V$ and $R$ images of NGC 6366 in order to search for new variable stars down to $V=19.5$ mag. The paper is organized as follows: in section 2 we summarise the observations performed and describe the data reduction methodology. In section 3 we discuss our approach to searching for variable stars. In section 4 we report on the new variables found and discuss their nature. In section 5 we briefly discuss the Fourier decomposition of the light curve of the RR Lyrae star, V1, and its implications for the metallicity and distance of NGC 6366. In section 6 we present our conclusions. ",
        "conclusions": "\\label{sec:Concl}  The use of difference imaging has enabled us to perform precision photometry in the globular cluster NGC~6366 and thereby lead to the detection of confirmed and possible new variables. The difference imaging technique used employs a new algorithm for determining the convolution kernel as a pixel grid. Among the new variables that we have found, we have identified one SX Phe star with at least three active modes, two possible AC's (or P2C's), one eclipsing binary and three long period red variables. We have also detected possible variations in a group of long term variables and short period eclipsing binaries likely to be of the W~UMa type.  Despite the position of V1 very near the center of NGC~6366, the membership of V1 in the cluster is doubted mainly because the star is $\\sim 0.3$ mag fainter than the red HB in the CMD, and because an inneficient convection transport in the red giants in NGC~6366 cannot be invoked as a possible cause for a real underluminosity of V1. Therefore, the metallicity and distance estimated for V1 from the Fourier technique, cannot be considered as representative of the cluster. We note however that the metallicity found for  V1, [Fe/H]$_{\\rm ZW}=-0.87 \\pm 0.14$, is very similar to values ascribed to NGC~6366 by independent spectroscopic estimates for giant stars in the cluster, (e.g [Fe/H]$_{\\rm ZW}=-0.85~\\pm~0.10$, Da Costa \\& Seitzer 1989; [Fe/H]$_{\\rm ZW}=-0.58 \\pm 0.14$, Rutledge et al. 1997).  The distance to V1 was estimated as $d({\\rm V1}) = 3.2 \\pm 0.1$ kpc. If V1 is shifted 0.3 mag to the ZAHB, and considering that V1 might be evolved above the ZAHB, an upper limit for the distance to the cluster of $d({\\rm NGC~6366}) = 2.8 \\pm 0.1$ kpc can also be estimated. An independent determination of the distance to NGC~6366 from the $P-L$ relationship for SX Phe stars and the pulsation modes identified in the SX Phe star V6 found in the cluster, gives the distance $d({\\rm V6}) = 2.7 \\pm 0.1$ kpc, which is consistent with the upper limit determined from V1.  The position of V1 on the $M_V$-[Fe/H] plane suggests a linear extrapolation to the metal-rich domain. Inclusion of RR Lyrae stars in the metal rich clusters NGC~6388 and NGC~6441 seem to suppot the non-linear behavior of the $M_V$-[Fe/H] relationship, however it must be stressed that those RR Lyraes have unusual long periods for their amplitudes and then, the Fourier decomposition calibrations to determine their [Fe/H] and $M_V$ values may not be applicable.  The $M_V(RR)$-[Fe/H] relationship derived from the Fourier results compares well, within the uncertainties, with the clusters distribution from the analysis of Caputo et al. (2000) and with the $M_V(ZAHB)$-[Fe/H] relationship prediction from ZAHB models of VandenBerg et al. (2000) once evolution from the ZAHB is considered. The excellent agreement of the position of V1 with this theoretical prediction further supports the idea that V1 is an evolved object from the ZAHB and that its apparent underluminosity in the CMD is due to its non-membership in NGC~6366. This interpretation has to compete with the otherwise straight one that, based on the similarity of the metallicities of the V1 and of the cluster and sitting V1 so close to the center of the cluster, V1 is very likely a member of the cluster. In this later case the observed underluminosity of V1 relative to the red HB is yet to be understood.  \\bigskip  \\noindent We are grateful to the support astronomers of IAO, at Hanle and CREST (Hosakote), for their very efficient help while acquiring the data. AAF and VRL acknowledge support from DGAPA-UNAM grant through project IN108106. RF and PR are grateful to the CDCHT-ULA for finantial support through project C-1544-08-05-F. Very useful corrections, comments and suggestions made by an anonymous and very constructive referee are deeply appreciated, particularly those on the non-membership of V1 in NGC~6366. This work has made a large use of the SIMBAD and ADS services, for which we are thankful."
    },
    "0808/0808.1572_arXiv.txt": {
        "abstract": "While the topology of the Universe is at present not specified by any known fundamental theory, it may in principle be determined through observations. In particular, a non-trivial topology will generate pairs of matching circles of temperature fluctuations in maps of the cosmic microwave background, the so-called circles-in-the-sky. A general search for such pairs of circles would be extremely costly and would therefore need to be confined to restricted parameter ranges. To draw quantitative conclusions from the negative results of such partial searches for the existence of circles we need a concrete theoretical framework. Here we provide such a framework by obtaining constraints on the angular parameters of these circles as a function of cosmological density parameters and the observer's position. As an example of the application of our results, we consider the recent search restricted to pairs of nearly back-to-back circles with negative results. We show that assuming the Universe to be very nearly flat, with its total matter-energy density satisfying the bounds $0<|\\Omega_{0}-1|$ $\\lesssim10^{-5}$, compatible with the predictions of typical inflationary models, this search, if confirmed, could in principle be sufficient to exclude a detectable non-trivial cosmic topology for most observers. We further relate explicitly the fraction of observers for which this result holds to the cosmological density parameters. ",
        "introduction": "One of the central open questions regarding our understanding of the Universe concerns its shape (topology), and in particular whether it is finite or infinite (see, e.g., the reviews~\\cite{CosmTopReviews}). A major difficulty in this regard is that general relativity, being a metric theory, does not specify the topology of the Universe.% \\footnote{In this work, in line with the usage in the literature, by topology of the Universe we mean the topology of its three dimensional spatial sections $M$.} Despite our present-day inability to predict the topology of the Universe from a fundamental theory, it may in principle be determined through observations. This requires high resolution observations, which have increasingly become available in recent years. Most notably, the on-going accumulation of data by the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{wmap} and other Cosmic Microwave Background (CMB) surveys have, on the one hand, provided strong support for the inflationary scenario~\\cite{Linde2}, and  the very near flatness of the Universe, and on the other hand made it feasible to perform systematic searches for possible evidence of a non-trivial topology of the Universe~\\cite{Cornish-etal03,CitSRedux,% Roukema-etal2004,Aurich-etal2006} (see also the related Refs.~\\cite{RelatedCirc}).  In the context of general relativity, the observable Universe seems to be well described by a $4$-manifold $\\mathcal{M} =\\mathbb{R}\\times M$ with locally homogeneous and isotropic spatial sections $M$ and endowed with a Robertson-Walker metric \\begin{equation} ds^{2}=-dt^{2}+a^{2}(t)\\left[ d\\chi^{2}+S_k^{2}(\\chi)(d\\theta ^{2}+\\sin^{2}\\theta d\\phi^{2}) \\right]  \\;, \\label{RWmetric} \\end{equation} where $a(t)$ is the scale factor and $S_k(\\chi)$ takes the forms $\\chi\\,,$ $\\sin\\chi\\,,$ or $\\sinh\\chi$ depending upon whether the geometry of the spatial sections is euclidian, spherical or hyperbolic with curvature parameters $k=0,\\pm 1$. These geometries in turn are determined by finding out whether the total energy-matter density of the Universe, $\\Omega_0$, is equal to, greater than or smaller than 1. Often the homogeneous and isotropic spatial sections $M$ are assumed to be the simply connected $3$-manifolds: euclidian $\\mathbb{R}^{3}$, spherical $\\mathbb{S}^{3}$, or hyperbolic $\\mathbb{H}^{3}$. However, given the metrical (local) nature of general relativity, they can also be multiply connected $3$-manifolds (which we assume to be compact and orientable) $M=\\widetilde{M}/\\Gamma$, where the covering space $\\widetilde{M}$ is respectively $\\mathbb{R}^{3}$, $\\mathbb{S}^{3}$ or $\\mathbb{H}^{3}$ depending on $k$, and $\\Gamma$ is a discrete and fixed point-free group of isometries of $\\widetilde{M}$ called the holonomy group. The local geometry of the spatial sections $M$ thus constrains, but does not dictate, its topology.  The immediate observational consequence of such multiple connectedness is that an observer could potentially detect multiple images of radiating sources. In particular, in a universe with a detectable non-trivial topology the  last scattering surface (LSS) intersects its topological images in the so called circles-in-the-sky~\\cite{CSS1998}, i.e., pairs of matching circles of equal radii, centered at different points of the LSS with the same distribution of temperature fluctuations along both circles. In this way, to observationally probe a non-trivial topology on the largest available scales, one needs to scrutinize the CMB sky-maps in order to extract such correlated circles in order to determine the topology of the Universe. Thus, a detectable non-trivial cosmic topology is an observable attribute, which can be probed through the circles-in-the-sky for all locally homogeneous and isotropic universes with no assumptions on the cosmological density parameters.  The conditions for the detectability of cosmic topology were studied in Refs.~\\cite{TopDetec} for classes of hyperbolic and spherical manifolds, as a function of cosmological density parameters. These studies were extended to the case of \\emph{generic} manifolds and general qualitative results were obtained for the \\emph{inflationary limit} in Ref.~\\cite{Mota-etal} (see also Ref.~\\cite{Mota-etal-b}). Furthermore, the inverse question of whether the detection of a non-trivial cosmic topology can be used to set constraints on cosmological density parameters has been studied for \\emph{specific} topologies ~\\cite{Previous,Related}.  Our aims here are twofold. Firstly, to extend the previous general results by the authors~\\cite{Mota-etal}, in order to give, for \\emph{generic} detectable non-trivial topologies, a concrete relationship between the angular parameters associated with the circles-in-the-sky and the cosmological density parameters in the \\emph{inflationary limit}. Secondly, by using these relations, and taking into account the observer positions, we provide a set up to interpret the negative results of recent and future searches for circles-in-the-sky in the WMAP data. In this way, we can concretely specify, for example, the extent to which the recent searches for circles-in-the-sky may be used to exclude detectable non-trivial cosmic topologies for most observers.  The structure of the paper is as follows. In Section~\\ref{Sec2} we give a brief account of the  pre-requisites necessary for the following sections, including a brief discussion of the so called inflationary limit, which we use to obtain bounds on the detectability of cosmic topology. In Section~\\ref{Sec3} we combine these bounds with our previous results to obtain bounds for the parameters of detectable holonomies in the inflationary limit.  We quantify these bounds by giving a concrete measure of the observers for whom the detectability holds. In Section~\\ref{Sec4} we recast our bounds on detectable holonomies in terms of bounds on the angular parameter that measures the deviation from antipodicity in pairs of matching circles generated by detectable holonomies. These bounds are then compared in Section~\\ref{Sec5} to the parameters adopted in the recent searches for circles-in-the-sky to determine whether and under which conditions it is possible to rule out a detectable non-trivial cosmic topology given the inflationary limit. Finally, in Section~\\ref{Sec6}, we conclude with a brief discussion of the significance of our results to draw quantitative conclusions from such negative results of searches for circles. ",
        "conclusions": ""
    },
    "0808/0808.3607_arXiv.txt": {
        "abstract": "%  The DAMA/NaI and DAMA/LIBRA annual modulation data, which may be interpreted as a signal for the existence of weakly interacting dark matter (WIMPs) in our galactic halo, are examined in light of null results from other experiments: CDMS, XENON10, CRESST~I, CoGeNT, TEXONO, and Super-Kamiokande (SuperK). We use the energy spectrum of the combined DAMA modulation data given in 36 bins, and include the effect of channeling.  Several statistical tools are implemented in our study: likelihood ratio with a global fit and with raster scans in the WIMP mass and goodness-of-fit (g.o.f.).  These approaches allow us to differentiate between the preferred (global best fit) and allowed (g.o.f.) parameter regions. It is hard to find WIMP masses and couplings consistent with all existing data sets; the surviving regions of parameter space are found here.  For spin-independent (SI) interactions, the  best fit DAMA regions are ruled out to the 3$\\sigma$ C.L., even with channeling taken into account.  However, for WIMP masses of $\\sim$8~GeV % some parameters outside these regions still yield a moderately reasonable fit to the DAMA data and are compatible with all 90\\% C.L.\\ upper limits from negative searches, when channeling is included.  For spin-dependent (SD) interactions with proton-only couplings, a range of masses below 10~GeV is compatible with DAMA and other experiments, with and without channeling, when SuperK indirect detection constraints are included; without the SuperK constraints, masses as high as $\\sim$20~GeV are compatible. For SD neutron-only couplings we find no parameters compatible with all the experiments. Mixed SD couplings are examined: \\eg\\ $\\sim$8~GeV mass WIMPs with $\\anSD = \\pm \\apSD$ are found to be consistent with all experiments.  In short, there are surviving regions at low mass for both SI and SD interactions; if indirect detection limits are relaxed, some SD proton-only couplings at higher masses also survive. ",
        "introduction": "Introduction}  Among the best motivated candidates for dark matter are Weakly Interacting Massive Particles (WIMPs).  Direct searches for dark matter WIMPs aim at detecting the scattering of WIMPs off of nuclei in a low-background detector. These experiments measure the energy of the recoiling nucleus, and are sensitive to a signal above a detector-dependent energy threshold \\cite{Jungman:1995df}.  The discovery of an annual modulation by the DAMA/NaI experiment \\cite{Bernabei:2003za}, confirmed by the new experiment DAMA/LIBRA~\\cite{Bernabei:2008yi} of the same collaboration, is the only positive signal seen in any dark matter search.  The DAMA collaboration has found an annual modulation in its data compatible with the signal expected from dark matter particles bound to our Galactic Halo \\cite{Drukier:1986tm,Freese:1987wu}.  However, other direct detection experiments, \\eg\\ CDMS~\\cite{Akerib:2003px,Ahmed:2008eu}, CoGeNT~\\cite{Aalseth:2008rx,Collar:2008pc}, COUPP~\\cite{Behnke:2008zz,Collar:2008pc2}, CRESST~\\cite{Angloher:2002in,Stodolsky:2004dsu}, KIMS~\\cite{Lee.:2007qn}, TEXONO~\\cite{Lin:2007ka,Wong:2008pc}, and XENON10~\\cite{Angle:2007uj,Angle:2008we} have not found any signal from WIMPs. It has been difficult to reconcile a WIMP signal in DAMA with the other negative results~\\cite{previous}, for the case of canonical WIMP masses $\\sim 100$GeV with standard weak interactions motivated by Supersymmetry (SUSY).  Yet, for light WIMPs, such compatibility was possible.  Papers written prior to the latest DAMA/LIBRA results found regions in cross-section and WIMP mass that reconciled all null results with DAMA's positive signal, even assuming a standard halo model: WIMPs with spin independent interactions~\\cite{Gelmini:2004gm,Gondolo:2005hh} in the mass range 5--9~GeV, and WIMPs with spin dependent interactions~\\cite{Savage:2004fn} in the mass range 5--13~GeV were found to be compatible with all existing data. In addition, Ref.~\\cite{Gelmini:2004gm,Gondolo:2005hh} studied not only the case of the standard halo model, but also models with dark matter streams; they found that nonstandard velocity distributions could reconcile all the existing data sets including DAMA.  Other dark matter explanations invoke inelastic scattering of WIMPs \\cite{TuckerSmith:2001hy}.  Since these studies of a few years ago, DAMA/LIBRA has reported new experimental results.  First, of particular interest is the new presentation of the modulation data in 36 separate energy bins, which allows much more detailed investigations of the data.  Second, in the past year another possible complication has come to light. The DAMA collaboration has pointed out that ``channeling'' may affect the interpretation of their data.  In general only a fraction (known as the quenching factor) of the recoil energy deposited by a WIMP is transferred to electrons and is converted into useful signal in the DAMA detector (\\eg\\ ionization or scintillation); the remainder is converted into phonons and heat and goes undetected.  Hence measured energies must be corrected for this behavior to obtain the proper recoil energy, by dividing by this quenching factor.  The DAMA detector is composed of NaI, with quenching factors $Q_{Na} = 0.3$ and $Q_I = 0.09$. Yet, recently, the DAMA collaboration pointed out that some fraction of the nuclear recoils travel in straight paths along the crystal in such a way as to lose very little energy to other atoms and to heat, giving nearly all their energy to electrons; \\ie, for these cases the measured energy corresponds very nearly to the energy deposited by the WIMP and thus have $Q \\approx 1$.  As a consequence, the detector is sensitive to lower mass WIMPs than previously thought.  The DAMA threshold is 2~keVee (electron equivalent). Using the above quenching factors this was thought to correspond to 7~keV recoil energy for sodium and 22~keV recoil energy for iodine; yet, due to channeling the actual threshold for some of the events can be as low as 2~keV.  The new channeling studies revived the possiblity of DAMA's compatibility with other data sets, particularly at low masses.  Light neutralinos as WIMPs with masses as low as 2~GeV \\cite{Gabutti:1996qd} or 6~GeV \\cite{Bottino:2002ry, Bottino:2008mf} have been considered, even with the cross sections we find necessary for WIMPs to be compatible with the DAMA signal and all other negative searches (for spin independent interactions)~\\cite{Bottino:2008mf}.  In any event, in this paper we proceed in a purely phenomenological way in choosing the WIMP mass and cross section, with either spin-independent or spin-dependent cross sections.  We do not attempt to provide an elementary particle model to support the values of masses and cross sections.  As justification of our approach, let us recall that there is no proven particle theory of dark matter.  The candidates we are considering are stable neutral particles which have very small cross sections with nucleons. Regarding their production in accelerators, they would escape from the detectors without interacting.  Unless there is a concrete specific model relating our neutral candidate to other charged particles (which can, if fact, be observed) there is no way such particles could be found in accelerators. The usual signature searched for in accelerators, for example at LEP, Tevatron or LHC, is the emission of a charged particle related to the neutral particle in question.  For example, searching for ``neutralinos'' one puts bounds on one of its cousins, a ``chargino'', or another relative, a ``slepton''. Without a detailed model there are no accelerator bounds on neutral dark matter candidates.  An alternative way to search for WIMP dark matter of relevance to this paper is via indirect detection of WIMP annihilation in the Sun.  When WIMPs pass through a large celestial body, such as the Sun or the Earth \\cite{indirectdet:solar,indirectdet:earth}, interactions can lead to gravitational capture if enough energy is lost in the collision to fall below the escape velocity.  Captured WIMPs soon fall to the body's core and eventually annihilate with other WIMPs.  These annihilations lead to high-energy neutrinos that can be observed by Earth-based detectors such as Super-Kamiokande \\cite{Desai:2004pq} (SuperK), which currently provides the tightest indirect detection bounds at light WIMP masses, AMANDA \\cite{Ahrens:2002eb,Ackermann:2005fr}, IceCube \\cite{Ahrens:2002dv} and ANTARES \\cite{Blanc:2003na}. The annihilation rate depends on the capture rate of WIMPs, which is in turn determined by the WIMP scattering cross section off nuclei in the celestial body.  While the Earth is predominantly composed of spinless nuclei, the Sun is mostly made of hydrogen, which has spin. Thus the spin-dependent cross section of WIMPs off nucleons can be probed by measuring the annihilation signals from WIMP annihilation in the Sun.  Other indirect detection methods search for WIMPs that annihilate in the Galactic Halo or near the Galactic Center where they produce neutrinos, positrons, or antiprotons that may be seen in detectors on the Earth \\cite{indirectdet:galactichalo,indirectdet:galacticcenter}. Of course, these rely upon additional physics (WIMP annhilation) beyond what direct detection experiments rely on (WIMP scattering).  Thermal WIMPs can only achieve the correct relic density today if their annihilation cross sections are fixed to be the value we use when computing the SuperK limits; yet one can imagine alternatives in which case the SuperK bounds should not be used.  In this paper, we reconsider the issue of compatibility of the new DAMA/NaI and DAMA/LIBRA (hereafter DAMA) results with all other experimental results.  We use the 36 bins of modulation data and take into account the channeling effect.  We restrict our studies to the standard isothermal model for the dark matter halo.  We carefully investigate both spin-independent as well as spin-dependent couplings.  Our procedure is the following. Assuming the standard dark halo model, we find the viable region by making sure that we produce the correct amplitudes for the DAMA modulation.  We then compare this region with the most stringent current experimental constraints from negative searches.  We use three different statistical methods to find the DAMA region: the likelihood method with (i) a global fit and (ii) raster scans in the WIMP mass, and (iii) goodness-of-fit.  These methods and the different information derived from them are described in detail. ",
        "conclusions": "Conclusions}  The latest DAMA/LIBRA annual modulation data, which may be interpreted as a signal for the existence of weakly interacting dark matter (WIMPs) in our Galactic Halo, are examined in light of null results from other experiments.  We investigate consequences for both spin-independent (SI) and spin-dependent (SD) WIMP couplings as well as mixed SI \\& SD couplings.  The positive signal for annual modulation is compared to negative results from CDMS, CoGeNT, CRESST~I, TEXONO, XENON10, and Super-Kamiokande (SuperK). In addition to the new experimental results, two new ingredients are included compared to our previous studies: 1) the energy spectrum of DAMA/LIBRA data is now given in 36 bins, and 2) the effect of channeling is taken into account, which lowers the nuclear recoil energy threshold and increases sensitivity to low WIMP masses. Several statistical tools are implemented in our study: best fit regions are found; goodness-of-fit (g.o.f.); and raster scans over the WIMP mass. These approaches allow us to differentiate between the most preferred (best fit) and allowed (g.o.f.) parameters. It is hard to find WIMP masses and couplings consistent with all existing data sets; the surviving regions of parameter space are found here.  For SI interactions, the best fit DAMA regions of any WIMP mass are ruled out to 3$\\sigma$, even with channeling taken into account.  However, some parameters outside the best fit regions still yield a moderately reasonable fit to the DAMA data: we find parameters in the SI case at masses of $\\sim$8--9~GeV that are compatible with each of the experiments at better than the 3$\\sigma$ level when channeling is included. For SD interactions, we find no neutron-only couplings that are compatible with the DAMA data to within the 3$\\sigma$ level and satisfy other experimental constraints, with or without channeling included. For SD proton-only couplings, a range of masses below 10~GeV is compatible with DAMA and other experiments, with and without channeling, when SuperK indirect detection constraints are included; without the SuperK constraints, masses as high as $\\sim$20~GeV are compatible. Mixed couplings with $\\anSD = \\pm \\apSD$ are found to be consistent with all experiments at $\\sim$8~GeV. The general case of mixed SD proton and neutron couplings are examined as well. In short, there are surviving regions at low mass for both SI and SD interactions (although only moderately so for the SI case) and at high masses for the case of SD proton-only couplings if indirect detection limits are relaxed.  While in the process of preparing this paper, several papers were presented that also examine the compatibility of the DAMA signal with other negative experimental results, always assuming that the dark matter consists of WIMPs \\cite{Bottino:2008mf, Petriello:2008jj,Chang:2008xa,Chang:2008gd,Fairbairn:2008gz, Hooper:2008cf} (although other non-WIMP candidates were studied as well, such as mirror~\\cite{Foot:2008nw}, composite~\\cite{Khlopov}, and WIMPless~\\cite{Feng:2008dz} dark matter). Our results are compatible with the results of these other studies of WIMPs, although our conclusions sometimes differ. In all of them channeling is taken into account.  Ref.~\\cite{Bottino:2008mf} concludes that the recent CDMS results are compatible with the DAMA regions as derived by the DAMA collaboration, for WIMPs with spin independent interactions with mass in the 7 to 10 GeV range, a results similar to ours. Ref.~ \\cite{Petriello:2008jj} used only the raster scan and the DAMA data binned into two bins to study WIMPs with spin independent interactions (while here we use 36 data bins with this method) and find light WIMPs compatible with all experimental results, when channelling is included.  Refs.~\\cite{Chang:2008xa} and~\\cite{Fairbairn:2008gz} both studied WIMPs with spin independent interactions using global fits with two parameters and the data binned into many bins, thus including the spectral modulation amplitude information.  Their results are similar to those of our global fits, but they conclude that light WIMPs with spin independent interactions ``cannot account for the data''~\\cite{Chang:2008xa} or ``are strongly disfavoured''~\\cite{Fairbairn:2008gz} within the standard halo model (they considered also the limits imposed by the DAMA total spectrum, not included here, which, however, affect heavier WIMPs). Ref.~\\cite{Hooper:2008cf} studies indirect detection bounds which may be very important for light WIMPs, but depend on the WIMP-annihilation channels."
    },
    "0808/0808.0118_arXiv.txt": {
        "abstract": "{} {The aim of this paper is to study the basic equations of the chemical evolution of galaxies with gas flows.  In particular, we focus on models in which the outflow is differential, namely in which the heavy elements (or some of the heavy elements) can leave the parent galaxy more easily than other chemical species such as H and He.} {We study the chemical evolution of galaxies in the framework of simple models, namely we make simplifying assumptions about the lifetimes of stars and the mixing of freshly produced metals.  This allows us to solve analytically the equations for the evolution of gas masses and metallicities.  In particular, we find new analytical solutions for various cases in which the effects of winds and infall are taken into account.} {Differential galactic winds, namely winds carrying out preferentially metals, have the effect of reducing the global metallicity of a galaxy, with the amount of reduction increasing with the ejection efficiency of the metals.  Abundance ratios are predicted to remain constant throughout the whole evolution of the galaxy, even in the presence of differential winds.  One way to change them is by assuming differential winds with different ejection efficiencies for different elements.  However, simple models apply only to elements produced on short timescales, namely all by Type II SNe, and therefore large differences in the ejection efficiencies of different metals are unlikely.} {Variations in abundance ratios such as [O/Fe] in galaxies, without including the Fe production by Type Ia supernovae, can in principle be obtained by assuming an unlikely different efficiency in the loss of O relative to Fe from Type II supernovae.  Therefore, we conclude that it is not realistic to ignore Type Ia supernovae and that the delayed production of some chemical elements relative to others (time-delay model) remains the most plausible explanation for the evolution of $\\alpha$-elements relative to Fe.} ",
        "introduction": "\\label{intro}  When a stellar system is formed out of gas, the dying stars enrich the remaining gas so that new stars are metal-enriched compared to their ancestors.  Simple models of chemical evolution have been used since decades to interpret the evolution of the chemical composition of stars and gas.  Pioneering works have determined the evolution of the metallicity in the Galaxy (Schmidt \\cite{sch63}; Searle \\& Sargent \\cite{ss72}; Tinsley \\cite{tin74}; Pagel \\& Patchett \\cite{pp75}) by making very simplifying hypotheses about the initial mass function (IMF), the lifetime of stars and the mixing of chemical elements with the surrounding interstellar medium (ISM).   These models remain useful guides for understanding the main chemical properties of galaxies, but it was soon realized that they cannot give a complete picture of galaxy evolution.  In particular, simple closed box models cannot reproduce the distribution of metallicity of long-living stars in the galaxy ({\\it G-dwarf problem}, Lynden-Bell \\cite{lyn75}; Rana \\cite{rana91} and references therein).  The most obvious solution to the G-dwarf problem is to relax the hypothesis that the solar neighborhood evolved as a closed box and to allow for infall of gas (Tinsley \\cite{tin80}).  Increasingly complex models of chemical evolution including gas flows (infall, outflow, radial flows), but always assuming the instantaneous recycling approximation (IRA)\\footnote{IRA assumes that stars with masses larger than 1 M$_{\\odot}$ die instantaneously whereas those below 1 M$_{\\odot}$ live forever (Matteucci \\cite{mat01})}, have been proposed, covering a large variety of galactic morphologies (Larson \\cite{lar76}; Clayton \\cite{clay88}; Edmunds \\cite{edm90}; K\\\"oppen \\cite{koe94}; Martinelli \\& Matteucci \\cite{mm00}).  These models consider the global metallicity Z and neglect metallicity-enriched outflows, so they cannot be used to predict variations of abundance ratios.  Pagel \\& Tautvai\\v{s}ien\\.{e} (\\cite{pt98}) have calculated abundance ratios using analytical models, but they had to assume that chemical elements such as iron are ejected with a fixed time delay $\\Delta$ after the time of star formation.  This is a rough approximation since it is known that the progenitors of SNeIa span a very wide lifetime (see e.g. Matteucci \\& Recchi \\cite{mr01}).  Only detailed numerical models relaxing IRA and including the chemical enrichment from Type Ia SNe (e.g. Matteucci \\& Greggio \\cite{mg86}) allow one to follow in detail the evolution of single elements.  Matteucci \\& Greggio (\\cite{mg86}) showed in detail the effect of the time-delay model, already suggested by Tinsley (\\cite{tin80}) and Greggio \\& Renzini (\\cite{gr83}), in particular the effect of a delayed Fe production by Type Ia SNe on abundance ratios involving $\\alpha$-elements (O, Mg, Si).  Outflows and galactic winds are produced when the thermal energy of a galaxy (produced by energetic events such as SN explosions and stellar winds) is enough to unbind a large fraction of the ISM (Mathews \\& Baker \\cite{mb71}).  The existence of {\\it differential winds}, namely galactic winds in which the metals (or some of them) are ejected out of the parent galaxy more efficiently than the other elements, has been first hypothesized and applied to the chemical evolution of galaxies by Pilyugin (\\cite{pil93}) and Marconi et al. (\\cite{mmt94}). In particular, Marconi et al. (\\cite{mmt94}) suggested that the chemical properties of gas-rich dwarf irregular galaxies, most remarkably the spread in H/He vs. O/H and N/O vs. O/H, can be more easily explained by assuming differential winds, in which products of SNeII (mostly $\\alpha$-elements and one third of iron) can escape the potential well of a galaxy more easily than other elements formed in low and intermediate mass stars.  Hereafter we will call {\\it selective winds} the differential winds in which different metals have different ejection efficiencies.  The existence of differential and selective winds has been theoretically confirmed by several detailed chemodynamical simulations (MacLow \\& Ferrara \\cite{mf99}; D'Ercole \\& Brighenti \\cite{db99}; Recchi et al. \\cite{rmd01}; Fujita et al. \\cite{fuj04}) and it has been observationally confirmed by determining the metallicity of outflows in starburst galaxies (Martin et al. \\cite{mkh02}; Ott et al. \\cite{owb05}).  Recently, Salvadori et al. (\\cite{sfs08}) have suggested differential winds coupled with relaxation of IRA in massive stars to explain the steep decline of the [O/Fe] ratio observed in the stars of Sculptor, without invoking any Fe production from Type Ia SNe.  In this paper we apply the hypothesis of differential and selective winds in the framework of simple models of chemical evolution with IRA.  The chemical evolution of galaxies is a complex phenomenon that can be properly treated only by means of detailed numerical codes, possibly taking into account also the hydrodynamical evolution of the system, as we have explained above.  However, in spite of the oversimplifications introduced by simple chemical evolution models, the results they produce are very linear and simple to interpret. They can therefore illustrate in a direct way the effect of differential winds on the chemical evolution of galaxies, as well as the effect of relaxing IRA.  Moreover, for clarity purposes, we show also the results of detailed numerical models (Lanfranchi \\& Matteucci \\cite{lm03}; \\cite{lm04}) as a comparison.  The plan of the paper is as follows: in Sect. 2 we summarize the assumptions of the simple models of chemical evolution and the results of models including gas flows (infall and outflow).  In Sect. 3 we introduce the assumptions necessary to solve the equations of chemical evolution of galaxies with differential winds and produce some illustrative result.  In Sect. 4 we study the abundance ratios of some chemical elements predicted by models with differential and selective winds.  Finally, we discuss some implications of the results and draw the main conclusions in Sect. 5. ",
        "conclusions": "\\label{discussion}  In this paper we have studied general properties of simple chemical evolution models of galaxies with infall and differential winds, namely with a selective loss of metals from the parent galaxy.  We have seen that, assuming gas flows proportional to the SFR, the equations of the evolution of gas mass, total mass and metallicity Z of gas can be solved analytically, leading to a simple expression of the variation of Z as a function of the gas mass fraction $\\mu$ and of the ejection efficiency of metals $\\alpha$ (which is defined as the ratio between the metallicity of the galactic wind and the metallicity of the ISM).  As expected, the resulting metallicity of the galaxy turns out to be systematically lower than the metallicity attained by models with non-differential winds.  In particular, the metallicity of models with differential winds tends asymptotically (for low values of $\\mu$, namely during the late evolution of the galaxy) to a constant value which decreases almost linearly with $\\alpha$.  We have checked that, for exponential infall laws and star formation rates proportional to the gas mass (linear Schmidt law) a constancy between the star formation rate and the infall rate arises quite naturally during the late evolution of the galaxies, provided that the star formation timescale does not differ significantly from the infall timescale.  In the first few Gyr the ratio between infall rate and star formation rate can vary quite significantly but, in spite of that, the resulting evolution of the metallicity as a function of time does not change appreciably compared to the case with linear flows.  We have also analyzed models in which the metallicity of the infalling gas varies with time.  A realistic case of infalling gas with increasing metallicity can be envisaged in the interaction galactic winds/ICM.  Indeed, this interaction must certainly lead to some pollution of the gas external to the galaxy, therefore to a constant increase of the ICM metallicity, in particular when differential winds are considered.  In order to take it into account, we have made two extreme assumptions:  \\begin{itemize}  \\item the metals carried out of the galaxy instantaneously and uniformly mix with the surrounding ICM ({\\it instantaneous ICM mixing}).  \\item The metallicity of the infalling gas is set to be equal to the metallicity of the outflow, representing a physical situation in which the very same gas carried out of the galaxy by outflows (or a fraction of it) falls back to the galaxy due to the pull of its gravitational potential ({\\it galactic fountain}).  \\end{itemize} \\noindent Real physical situations should lie between these two extreme assumptions since the distribution of metals in the ICM is expected to be patchy (Domainko et al. \\cite{dom06}; Sanders \\& Fabian \\cite{sf06}; Tornatore et al. \\cite{tor07}).  The metallicities predicted by instantaneous ICM mixing models are of course larger than models in which the infalling gas has a pristine chemical composition. The magnitude of this difference depends on the relative amount of ISM and ICM masses.  Under normal circumstances, in which the ICM mass is larger than the ISM one, the difference in the final metallicity is only $\\sim$ 0.1 - 0.2 dex.  A larger ISM metallicity is obtained by galactic fountains models.  This result due to the fact that galactic fountains maximize the chemical enrichment of the infalling gas.  This metallicity tends however to the one attained by the model with $Z_A = 0$ if very large metal ejection efficiencies ($\\alpha \\gg 1$) are assumed.  We have then studied the metallicity ratios produced by differential wind models.  We have seen that the models with variable infall metallicities (both instantaneous ICM mixing and galactic fountain) simply predict that abundance ratios between chemical elements are given by the ratio of their respective stellar yields (if these do not change with metallicity).  Assuming a fixed infall metallicity $Z_A \\neq 0$ this ratio also depends on the value of $Z_A$ and on the infall parameter $\\Lambda$, but nevertheless it does not change during the evolution of the galaxy.  However, we stress here that simple models can be properly applied only to SNeII products, whose release is very prompt.  Iron is for instance mostly produced by SNeIa, whose explosion timescales range from a few tens of Myr to several Gyr, depending on the star formation history of the considered galaxy (Matteucci \\& Recchi \\cite{mr01}), therefore only the fraction of iron produced by SNeII ($\\sim$ 1/3 of the total) can be treated by the simple models.  A constant O/Fe ratio, as predicted by simple models is therefore expected because, in this case, both chemical elements are produced by massive stars, on similar timescales and they should therefore have a similar evolution.  However, O/Fe ratios are not constant in galaxies and larger ratios are observed in older, more metal-poor stars.  The most natural explanation of this trend is that the delayed production of iron by SNeIa reduces the initially high O/Fe ratio of the oldest stars ({\\it time-delay model}; Matteucci \\& Greggio \\cite{mg86}).  Simple models can produce variable abundance ratios either if we assume different ejection efficiencies for different chemical species (selective winds) or an IMF variable in time.  Both assumptions should be quite ``ad hoc''.  We have tested the case of selective winds for the O/Fe ratio showing that, if the O ejection efficiency is significantly larger than the Fe one, O/Fe decreases with time (and with Fe/H).  We have confirmed this result also by means of detailed numerical simulations in which IRA is relaxed showing that, neglecting SNeIa, the only way to get a knee in the [O/Fe] vs. [Fe/H] relation is by assuming $\\alpha$-enhanced galactic winds.  In a model suitable for dwarf spheroidal galaxies, in which the outflow rate exceeds the infall rate, the O/Fe is almost constant up to [Fe/H] $\\sim$ -1.5, then it bends down.  This behavior is consistent with the observation of many dwarf spheroidal galaxies of the Local Group (Venn et al. \\cite{venn04} and references therein).  If instead iron is ejected more efficiently than oxygen, the opposite happens, with O/Fe ratios increasing with time.  On the other hand, an iron ejection efficiency larger than the oxygen seems plausible, as it has been shown by chemodynamical simulations of dwarf galaxies experiencing violent bursts of star formation (Recchi et al.  \\cite{rmd01}; \\cite{rec02}). However, also these models predict O/Fe ratios decreasing with time (and with [Fe/H]) because SNeIa were not neglected and their contribution is more important than selective winds in determining the evolution of the O/Fe ratio.  Moreover, Recchi et al.'s model found larger iron ejection efficiencies precisely because of the delayed explosion of SNeIa.  These SNe explode in a medium already heated and diluted by the previous activity of SNeII, therefore SNeIa products can be very easily channeled along the galactic funnel.  If we neglect SNeIa, O and Fe are produced by the same stars and they should therefore share the same fate.  Our main conclusions can be summarized as follows:  \\begin{itemize}  \\item  The differential winds have the effect of reducing the global metallicity of a galaxy.  The magnitude of this reduction increases with the ejection efficiency of metals.  \\item  This remains true also if we allow a variable metallicity of the infalling gas.  Larger galactic metallicities are obtained if we assume that the infall has the metallicity of the outflowing gas (galactic fountain model).  \\item  Simple models with constant (i.e. non metallicity-dependent) yields predict constant abundance ratios, even in the presence of differential winds or variable infall metallicities.  The same occurs for detailed numerical models.  \\item   One way to get variable abundance ratios in simple models is by assuming selective winds, namely differential winds with different ejection efficiencies for each chemical element.  However, simple models apply only to chemical elements promptly released on short timescales, therefore large element-to-element variations of the ejection efficiencies are not very likely.  Abundance ratios predicted by simple models of chemical evolution are then likely to remain constant during the evolution of the galaxy.  \\end{itemize}"
    },
    "0808/0808.2437_arXiv.txt": {
        "abstract": "{Propagation of charged cosmic-rays in the Galaxy depends on the transport parameters, whose number can be large depending on the propagation model under scrutiny. A standard approach for determining these parameters is a manual scan, leading to an inefficient and incomplete coverage of the parameter space.} {In analyzing the data from forthcoming experiments, a more sophisticated strategy is required. An automated statistical tool is used, which enables a full coverage of the parameter space and provides a sound determination of the transport and source parameters. The uncertainties in these parameters are also derived. } {We implement a Markov Chain Monte Carlo (MCMC), which is well suited to multi-parameter determination. Its specificities (burn-in length, acceptance, and correlation length) are discussed in the context of cosmic-ray physics. Its capabilities and performances are explored in the phenomenologically well-understood Leaky-Box Model. } {From a technical point of view, a trial function based on binary-space partitioning is found to be extremely efficient, allowing a simultaneous determination of up to nine parameters, including transport and source parameters, such as slope and abundances. Our best-fit model includes both a low energy cut-off and reacceleration, whose values are consistent with those found in diffusion models. A Kolmogorov spectrum for the diffusion slope ($\\delta=1/3$) is excluded. The marginalised probability-density function for $\\delta$ and $\\alpha$ (the slope of the source spectra) are $\\delta\\approx 0.55-0.60$ and $\\alpha \\approx 2.14-2.17$, depending on the dataset used and the number of free parameters in the fit. All source-spectrum parameters (slope and abundances) are positively correlated among themselves and with the reacceleration strength, but are negatively correlated with the other propagation parameters.} {The MCMC is a practical and powerful tool for cosmic-ray physic analyses. It can be used to confirm hypotheses concerning source spectra (e.g., whether $\\alpha_i\\neq \\alpha_j$) and/or determine whether different datasets are compatible. A forthcoming study will extend our analysis to more physical diffusion models.} ",
        "introduction": "One issue of cosmic-ray (CR) physics is the determination of the transport parameters in the Galaxy. This determination is based on the analysis of the secondary-to-primary ratio (e.g., B/C, sub-Fe/Fe), for which the dependence on the source spectra is negligible, and the ratio remains instead mainly sensitive to the propagation processes (e.g., \\citealt{2001ApJ...555..585M} and references therein). For almost 20 years, the determination of these parameters relied mostly on the most constraining data, namely the HEAO-3 data, taken in 1979, which covered the $\\sim 1-35$~GeV/n range \\citep{1990A&A...233...96E}.  For the first time since HEAO-3, several satellite or balloon-borne experiments (see ICRC~2007 reporter's talk \\citealt{2008arXiv0801.4534B}) have acquired higher quality data in the same energy range or covered a scarcely explored range (in terms of energy, 1~TeV/n$-$PeV/n, or in terms of nucleus): from the balloon-borne side, the ATIC collaboration has presented the B/C ratio at $0.5-50$~GeV/n \\citep{2007arXiv0707.4415P}, and for H to Fe fluxes at 100 GeV$-$100 TeV \\citep{2006astro.ph.12377P}. At higher energy, two long-duration balloon flights will soon provide spectra for Z=1-30 nuclei. The TRACER collaboration has published spectra for oxygen up to iron in the GeV/n-TeV/n range \\citep{2007astro.ph..3707B,2008ApJ...678..262A}. A second long-duration flight took place in summer 2006, during which the instrument was designed to have a wider dynamic-range capability and to measure lighter B, C, and N elements. The CREAM experiment \\citep{2004AdSpR..33.1777S} flew a cumulative duration of 70 days in December 2004 and December 2005 (\\citealt{2006cosp...36.1846S}, and preliminary results in \\citealt{2006cosp...36.3129M} and \\citealt{2006cosp...36.3231W}), and again in December 2007. A fourth flight was scheduled for December 2008\\footnote{\\tiny http://cosmicray.umd.edu/cream/cream.html}. Exciting data will arrive from the PAMELA satellite \\citep{2007APh....27..296P}, which was successfully launched in June 2006 \\citep{2008AdSpR..42..455C}.  With this wealth of new data, it is relevant to question the method used to extract the propagation parameters. The value of these parameters is important to many theoretical and astrophysical questions, because they are linked, amongst others, to the transport in turbulent magnetic fields, sources of CRs, and $\\gamma$-ray diffuse emission (see \\citealt{2007ARNPS..57..285S} for a recent review and references). It also proves to be crucial for indirect dark-matter detection studies (e.g., \\citealt{2004PhRvD..69f3501D}, and \\citealt{2008PhRvD..77f3527D}). The usage in the past has been based mostly on a manual or semi-automated|hence partial|coverage of the parameter space (e.g., \\citealt{1992ApJ...390...96W}, \\citealt{1998ApJ...509..212S}, and \\citealt{2001ApJ...547..264J}). More complete scans were performed in \\citet{2001ApJ...555..585M,2002A&A...394.1039M}, and \\citet{2005JCAP...09..010L}, although in an inefficient manner: the addition of a single new free parameter (as completed for example in \\citealt{2002A&A...394.1039M} compared to \\citealt{2001ApJ...555..585M}) remains prohibitive in terms of computing time. To remedy these shortcomings, we propose to use the Markov Chain Monte Carlo (MCMC) algorithm, which is widely used in cosmological parameter estimates (e.g., \\citealt{2001CQGra..18.2677C}, \\citealt{2002PhRvD..66j3511L}, and \\citealt{2005MNRAS.356..925D}). One goal of the paper is to confirm whether the MCMC algorithm can provide similar benefits in CR physics.  The analysis is performed in the framework of the Leaky-Box Model (LBM), a simple and widely used propagation model. This model contains most of the CR phenomenology and is well adapted to a first implementation of the MCMC tool. In Sect.~\\ref{s:key}, we highlight the appropriateness of the MCMC compared to other algorithms used in the field. In Sect.~\\ref{s:MCMC}, the MCMC algorithm is presented. In Sect.~\\ref{s:implementation}, this algorithm is implemented in the LBM. In Sect.~\\ref{s:results}, we discuss the MCMC advantages and effectiveness in the field of CR physics, and present results for the LBM. We present our conclusions in Sect.~\\ref{s:conclusions}. Application of the MCMC technique to a more up-to-date modelling, such as diffusion models, is left to a forthcoming paper. ",
        "conclusions": "\\label{s:conclusions}  We have implemented a Markov Chain Monte Carlo to extract the posterior distribution functions of the propagation parameters in a LBM. Three trial functions were used, namely a standard Gaussian step, an N-dimensional Gaussian step and its covariance matrix, and a binary-space partitioning. For each method, a large number of chains were processed in parallel to accelerate the PDF calculations. The three trial functions were used sequentially, each method providing some inputs to the next: while the first one was good at identifying the general range of the propagation parameters, it was not as efficient in providing an accurate description of the PDF. The two other methods provided this accuracy, and the final results were based on PDF obtained from the chains processed with the binary-space partitioning.  Taking advantage of the sound statistical properties of the MCMC, confidence intervals for the propagation parameters can be given, as well as confidence contours for all fluxes and other quantities derived from the propagation parameters. The MCMC was also used to compare the impact of choosing different datasets and ascertain the merits of different hypotheses concerning the propagation models. Concerning the first aspect, we have shown that combining different B/C datasets leaves mostly unchanged the propagation parameters, while strongly affecting the assessment of the goodness of a model. We also show that at present, the \\pbar\\ data do not cover a sufficiently large energy range to constrain the propagation parameters, but they could be useful in crosschecking the secondary nature of the flux in the future.   In this first paper, we have focused on the phenomenologically well-understood LBM, to ease and simplify the discussion and implementation of the MCMC. In agreement with previous studies, we confirm that a model with a rigidity cutoff performs more successfully than one without and that reacceleration is preferred over no reacceleration. Such a model can be associated with a diffusion model with wind and reacceleration. As found in \\citet{2001ApJ...555..585M}, the best-fit models demand both a rigidity cutoff (wind) and reacceleration, but do not allow us to reconcile the diffusion slope with a Kolmogorov spectrum for turbulence. An alternative model with two slopes for the diffusion was used, but it is not favoured by the data. In a last stage, we allowed the abundance and slope of the source spectra to be free parameters, as well as the element abundances of C, N, and O. This illustrated a correlation between the propagation and source parameters, potentially biasing the estimates of these parameters. The best-fit model slope for the source abundances was $\\alpha\\approx 2.17$ using HEAO-3 data, compatible with the value $\\alpha\\approx 2.14$ for TRACER data. The MCMC approach allowed us to draw confidence intervals for the propagation parameters, the source parameters, and also for all fluxes.   A wealth of new data on Galactic CR fluxes are expected soon. As illustrated for the LBM, the MCMC is a robust tool in handling complex data and model parameters, where one has to fit simultaneously all source and propagation parameters. The next step is to apply this approach to more realistic diffusion models and larger datasets, on a wider range of nuclear species."
    },
    "0808/0808.0574_arXiv.txt": {
        "abstract": "% The excess of far-ultraviolet (far-UV) radiation in elliptical galaxies has remained one of their most enduring puzzles. In contrast, the origin of old blue stars in the Milky Way, hot subdwarfs, is now reasonably well understood: they are hot stars that have lost their hydrogen envelopes by various binary interactions. Here, we review the main evolutionary channels that produce hot subdwarfs in the Galaxy and present the results of binary population synthesis simulations that reproduce the main properties of the Galactic hot-subdwarf population. Applying the same model to elliptical galaxies, we show how this model can explain the main observational properties of the far-UV excess, including the far-UV spectrum, without the need to invoke {\\em ad hoc} physical processes.  The model implies that the UV excess is not a sign of age, as has been postulated previously, and predicts that it should not be strongly dependent on the metallicity of the population. ",
        "introduction": "One of the first major discoveries soon after the advent of UV astronomy was the discovery of an excess of light in the far-ultraviolet (far-UV) in elliptical galaxies (see the review by O'Connell 1999). This came as a complete surprise since elliptical galaxies were supposed to be entirely composed of old, red stars and not to contain any young stars that radiate in the UV.  Since then it has become clear that the far-UV excess (or upturn) is not a sign of active contemporary star formation, but is mainly caused by an older population of helium-burning stars or their descendants with a characteristic surface temperature of 25,000\\,K (Ferguson et al.\\ 1991), also known as hot subdwarfs or sdB stars. In recent years, it has become increasingly clear that their Galactic counterparts are predominantly produced by binary interactions (Maxted et al.\\ 2001; Morales-Rueda et al.\\ 2003; Lisker et al.\\ 2005; Green 2008).  Here, we first review the main evolutionary channels that produce hot subdwarfs in our Galaxy (see Han et al.\\ 2002, 2003 for further details) and then apply this model to elliptical galaxies to show that these may also be able to account for their far-UV excess (Han, Podsiadlowski, \\& Lynas-Gray 2007). ",
        "conclusions": "\\begin{figure}[t] \\psfig{figure=zhan_fig2.ps,width=0.85\\textwidth,angle=270} \\caption{Evolution of far-UV properties [the slope of the far-UV spectrum, $\\beta_{\\rm FUV}$, versus $(1550-V)$] for a two-population model of elliptical galaxies.  The age of the old population is assumed to be 12\\,Gyr (filled squares, filled triangles, or filled circles) or 5\\,Gyr (open squares, open triangles, or open circles). The mass fraction of the younger population is denoted as $f$ and the time since the formation as $t_{\\rm minor}$ [plotted in steps of $\\Delta \\log (t)=0.025$].  Note that the model for $f=100\\%$ (the dotted curve) shows the evolution of a simple stellar population with age $t_{\\rm minor}$.  The legend is for $b_{\\rm FUV}$, which is the fraction of the UV flux that originates from hot subdwarfs resulting from binary interactions. The effect of internal extinction is indicated in the top-left corner, based on the Calzetti internal extinction model with $E(B-V)=0.1$ (Calzetti et al.\\ 2000).  For comparison, we also plot galaxies with error bars from HUT (Brown et al.\\ 1997) and IUE observations (Burstein 1988).  The galaxies with strong signs of recent star formation are denoted with an asterisk (NGC 205, NGC 4742, NGC 5102).  } \\label{fig2} \\end{figure}  Figure~\\ref{fig1} shows our simulated evolution of the far-UV spectrum of a galaxy in which all stars formed at the same time, where the flux has been scaled relative to the visual flux (between 5000 and 6000\\AA) to reduce the dynamical range. At early times the far-UV flux is dominated by the contribution from single young stars. Binary hot subdwarfs become important after about 1.1\\,Gyr, which corresponds to the evolutionary timescale of a 2\\,$M_{\\odot}$ star and soon start to dominate completely. After a few Gyr the spectrum no longer changes appreciably.  There is increasing evidence that many elliptical galaxies had some recent minor star-formation events (Kaviraj et al.\\ 2007; Schawinski et al.\\ 2007), which also contribute to the far-UV excess.  To model such secondary minor starbursts, we have constructed two-population galaxy models, consisting of one old, dominant population with an assumed age $t_{\\rm old}$ and a younger population of variable age, making up a fraction $f$ of the stellar mass of the system.  In order to illustrate the appearance of the galaxies for different lookback times (redshifts), we adopted two values for $t_{\\rm old}$, of 12 Gyr and 5 Gyr, respectively; these values correspond to the ages of elliptical galaxies at a redshift of 0 and 0.9, respectively, assuming that the initial starburst occurred at a redshift of 5 and adopting a standard $\\Lambda$CDM cosmology with $H_0=72{\\rm km/s/Mpc}$, $\\Omega_{\\rm M}=0.3$ and $\\Omega_\\Lambda=0.7$. Our spectral modelling shows that a recent minor starburst mostly affects the slope in the far-UV SED.  We therefore define a far-UV slope index $\\beta_{\\rm FUV}$ as $f_\\lambda \\sim \\lambda ^{\\beta_{\\rm FUV}}$, where $\\beta_{\\rm FUV}$ is fitted between 1075\\AA~ and 1750\\AA. This parameter was obtained from our theoretical models by fitting the far-UV SEDs and was derived in a similar manner from observed far-UV SEDs of elliptical galaxies (Burstein et al.\\ 1988; Brown et al.\\ 1997), where we excluded the spectral region between 1175\\AA~ and 1250\\AA, the region containing the strong Ly$\\alpha$ line.  In order to assess the importance of binary interactions, we also defined a binary contribution factor $b={F_{\\rm b}/ F_{\\rm total}}$, where $F_{\\rm b}$ is the integrated flux between 900\\AA~ and 1800\\AA~ radiated by hot subdwarfs produced by binary interactions, and $F_{\\rm total}$ is the total integrated flux between 900\\AA ~and 1800\\AA. Figure~\\ref{fig2} shows the far-UV slope as a function of UV excess, a potentially powerful diagnostic diagram which illustrates how the UV properties of elliptical galaxies evolve with time in a dominant old population with a young minor sub-population.  For comparison, we also plot observed elliptical galaxies from various sources.  Overall, the model covers the observed range of properties reasonably well.  Note in particular that the majority of galaxies lie in the part of the diagram where the UV contribution from binaries is expected to dominate (i.e. where $b> 0.5$).  The two-component models presented here are still quite simple and do not take into account, e.g., more complex star-formation histories, possible contributions to the UV from AGN activity, non-solar metallicity or a range of metallicities.  Moreover, the binary population synthesis is sensitive to uncertainties in the binary modelling itself, in particular the mass-ratio distribution and the condition for stable and unstable mass transfer (Han et al.\\ 2003). We have varied these parameters and found that these uncertainties do not change the qualitative picture, but affect some of the quantitative estimates.  Despite its simplicity, our model can successfully reproduce most of the properties of elliptical galaxies with a UV excess: e.g., the range of observed UV excesses, both in $(1550-V)$ and $(2000-V)$ (e.g., Deharveng, Boselli, \\& Donas 2002), and their evolution with redshift.  The model predicts that the UV excess is not a strong function of age, and hence is not a good indicator for the age of the dominant old population, as has been argued previously (Yi et al.\\ 1999), but is very consistent with recent GALEX findings (Rich et al.\\ 2005).  We typically find that the $(1550-V)$ color changes rapidly over the first 1\\,Gyr and only varies slowly thereafter. This also implies that all old galaxies should show a UV excess at some level. Moreover, we expect that the model is not very sensitive to the metallicity of the population since metallicity does not play a significant role in the envelope ejection process (although it may affect the properties of the binary population in more subtle ways).  Our model is sensitive to both low levels and high levels of star formation.  It suggests that elliptical galaxies with the largest UV excess had some star formation activity in the relatively recent past ($\\sim 1\\,$Gyr ago).  AGN and supernova activity may provide supporting evidence for this picture, since the former often appears to be accompanied by active star formation, while supernovae, both core collapse and thermonuclear, tend to occur mainly within 1\\,--\\,2\\,Gyr after a starburst in the most favoured supernova models.  The modelling of the UV excess presented in this study is only a starting point: with refinements in the spectral modelling, including metallicity effects, and more detailed modelling of the global evolution of the stellar population in elliptical galaxies, we suspect that this may become a powerful new tool helping to unravel the complex histories of elliptical galaxies that a long time ago looked so simple and straightforward."
    },
    "0808/0808.0367.txt": {
        "abstract": "{ We review basic constraints on the acceleration of ultra-high-energy (UHE) cosmic rays (CRs) in astrophysical sources, namely the geometrical (Hillas) criterion and restrictions from radiation losses in different acceleration regimes. Using the latest available astrophysical data, we redraw the Hillas plot and figure out potential UHECR accelerators. For the acceleration in central engines of active galactic nuclei, we constrain the maximal UHECR energy for a given black-hole mass. Among active galaxies, only the most powerful ones, radio galaxies and blazars, are able to accelerate protons to UHE, though acceleration of heavier nuclei is possible in much more abundant lower-power Seyfert galaxies. }  %%% PACS numbers %\\pacno{98.70.Sa} ",
        "introduction": "\\label{sec:intro}  \\markboth{Physical conditions in potential UHECR sources I}{ Physical conditions in potential UHECR sources I}  The origin of ultra-high-energy (UHE; energy ${\\cal E} \\gtrsim 10^{19}$~eV) cosmic rays (CRs) remains unknown despite decades of intense studies (see e.g.\\ Ref.~\\cite{NaganoWatson} for a comprehensive review and Ref.~\\cite{Kachelriess:lectures} for a recent pedagogical introduction). Recent studies, notably the observation of the Greisen--Zatsepin--Kusmin~\\cite{G, ZK} cutoff by the HiRes experiment~\\cite{HiRes:cutoff}, further supported by results of the Pierre Auger observatory (PAO)~\\cite{PAOspectrum}, suggest that at least a large fraction of UHECRs is accelerated in cosmologically distant astrophysical sources. The birth of (scientific) UHECR astronomy, however, awaits our firm understanding of energies and primary composition of the observed cosmic rays as well as identification of at least one astrophysical object where these particles are accelerated.  Given  experimental ambiguities, it is important to understand theoretically, which astrophysical objects may serve as UHECR accelerators. It has been understood long ago that the UHECR sources should be distinguished objects with extreme physical conditions. One simple criterion is the geometrical one: the particle should not leave the accelerator before it gains the required energy. Presumably, the particle is accelerated by the electric field and confined by the magnetic one; then the geometrical criterion is expressed in terms of the particle's Larmor radius which should not exceed the linear size of the accelerator (see, e.g., Ref.~\\cite{StellarMagnetospheres}). In the context of UHECRs, this condition is recognized as the Hillas criterion~\\cite{Hillas} and is often presented graphically in terms of the Hillas plot where the accelerator size $R$ and the magnetic field $B$ are plotted. We note that even quite recent reviews use either cut-and-pasted or slightly refurbished versions of the original 25-years-old plot. However, astrophysics experienced enormous progress, if not a revolution, during these decades. One of the aims of this study is to give an updated version of the Hillas plot with references to either -- when possible -- measurements or estimates of the magnetic fields and sizes of potential astrophysical accelerators. The most important update corresponds to a wide variety of active galaxies whose sizes and magnetic fields differ by many orders of magnitude from one object to another so that some of them may, while most of them may not, accelerate particles to UHE.  Another restriction on the cosmic-ray accelerators is posed by the radiation losses which inevitably accompany the acceleration of a charged particle. The corresponding constraints were studied, in particular, in Refs.~\\cite{Hillas, Protheroe, Derishev, Medvedev}. The radiation losses depend on the particular field configuration and the maximal achievable energy of a particle in the loss-limited regime depends on the acceleration mechanism. Restricting to particular mechanisms or particular field configurations may result in would-be contradictory results, cf.\\ Refs.~\\cite{Derishev, Medvedev}. In this work, we review the radiation-loss constraint for different cases; they further limit the acceptable region on the updated Hillas plot. One of possible applications of these general constraints is a study of active galaxies correlated with the Auger events~\\cite{paper2}.  One should always keep in mind that even if both geometric and radiation constraints are satisfied, they do not yet guarantee particle acceleration to the corresponding energy. Each particular source should be discussed in the context of an acceleration mechanism operating there.  The rest of the paper is organised as follows. In Sec.~\\ref{sec:accel:general}, we review constraints on potential UHE accelerators, that is model-independent Hillas geometrical constraint and limitations due to radiation losses for particular acceleration mechanisms. In Sec.~\\ref{sec:accel:mf-measurements}, we take advantage of the modern astrophysical data and redraw the Hillas plot supplemented by the radiation-loss constraints. Our results are summarized  and discussed in Sec.~\\ref{sec:accel:summary} while brief conclusions are given in Sec.~\\ref{sec:concl}. Appendix~\\ref{app:ed} contains derivation of some formulae. ",
        "conclusions": "\\label{sec:concl}  We reviewed constraints on astrophysical UHE accelerators and presented the Hillas plot supplemented with radiation-loss constraints and updated with recent astrophysical data. Contrary to previous studies, we emphasised that active galaxies span a large area on the plot, and only the most powerful ones (radio galaxies, quasars and BL Lac type objects) are capable of acceleration of protons to UHE. If UHECR particles are accelerated close to the supermassive black holes in AGN, then most likely the mechanism is ``one-shot'' with energy losses dominated by the curvature radiation. Other potential UHE acceleration sites are jets, lobes, knots and hot spots of {\\em powerful} active galaxies, starburst galaxies and shocks in galaxy clusters. Acceleration of particles in supercluster-scale shocks, gamma-ray bursts and inner part of AGN is subject to additional constraints from $p \\gamma$ interactions which are not discussed here.  Unlike protons, heavy nuclei can be accelerated to UHE in circumnuclear regions of low-power active galaxies. Since these galaxies are abundant, this contribution to UHECR flux may be important, leading to a mixed primary cosmic-ray composition at highest energies.   The authors are indebted to D.~Gorbunov, S.~Gureev, A.M.~Hillas, V.~Lukash, A.~Neronov, S.~Popov and D.~Semikoz for interesting discussions. We acknowledge the use of online tools~\\cite{LEDA, NED}. This work was supported in part by the grants RFBR 07-02-00820 09-07-08388 (ST), NS-1616.2008.2 (ST) and by FASI government contracts 02.740.11.0244 (ST) and 02.740.11.5092 (KP, ST) and by the Dynasty Foundation (KP, ST). .  \\appendix"
    },
    "0808/0808.0268_arXiv.txt": {
        "abstract": "Quasar feedback has most likely a substantial but only partially understood impact on the formation of structure in the universe. A potential direct probe of this feedback mechanism is the Sunyaev-Zeldovich effect: energy emitted from quasar heats the surrounding intergalactic medium and induce a distortion in the microwave background radiation passing through the region. Here we examine the formation of such hot quasar bubbles using a cosmological hydrodynamic simulation which includes a self-consistent treatment of black hole growth and associated feedback, along with radiative gas cooling and star formation. From this simulation, we construct microwave maps of the resulting Sunyaev-Zeldovich effect around black holes with a range of masses and redshifts. The size of the temperature distortion scales approximately with black hole mass and accretion rate, with a typical amplitude up to a few micro-Kelvin on angular scales around 10 arcseconds. We discuss prospects for the direct detection of this signal with current and future single-dish and interferometric observations, including ALMA and CCAT. These measurements will be challenging, but will allow us to characterize the evolution and growth of supermassive black holes and the role of their energy feedback on galaxy formation. ",
        "introduction": "The temperature fluctuations in the cosmic microwave background, as measured by the Wilkinson Microwave Anisotropy Probe (WMAP) satellite (Bennett et al.\\ 2003) and numerous other microwave experiments (e.g., Dawson et al.\\ 2002; Rajguru et al.\\ 2005; Reichardt et al. 2008) have proven to be the single most powerful tool in constraining cosmology (Spergel et al.\\ 2007). The temperature anisotropy has been mapped with large statistical significance on angular scales down to around a quarter degree, where the dominant physical mechanisms contributing to the fluctuations arise from density perturbations at the epoch of recombination. Attention is now turning to arcminute angular scales, where temperature fluctuations arise due to interaction of the microwave photons with matter in the low-redshift universe (for a brief review, see Kosowsky 2003). These low-redshift and small-angle anisotropies are collectively known as ``secondary anisotropies'' in the microwave background. The most prominent among them is the Sunyaev-Zeldovich (SZ) effect (Sunyaev \\& Zeldovich 1972) from the inverse Compton scattering of the microwave photons due to hot electrons. The SZ effect provides a powerful method for detecting accumulations of hot gas in the universe (Carlstrom Holder \\& Reese 2002).  Galaxy clusters, which contain the majority of the thermal energy in the universe, provide the largest SZ signal; clusters were first detected this way through pioneering measurements over the past decade (e.g, Marshall et al.\\ 2001) and thousands of them will be detected by the upcoming SZ surveys like the Atacama Cosmology Telescope (ACT) (Kosowsky et al. 2006) and the South Pole Telescope (SPT) (Ruhl et al. 2004). However, a number of other astrophysical processes will also create SZ distortions. This includes SZ distortion from peculiar velocities during reionization (McQuinn et al. 2005, Illiev et al. 2006), supernova-driven galactic winds (Majumdar, Nath, \\& Chiba 2001), black hole seeded proto-galaxies (Aghanim, Ballad \\& Silk 2000), kinetic SZ from Lyman Break Galaxy outflow (Babich \\& Loeb 2007), effervescent heating in groups and clusters of galaxies (Roychowdhury, Ruszkowski \\& Nath 2005) and supernova from first generation of stars (Oh, Cooray, \\& Kamionkowski 2003). The SZ distortion in galactic scales (hot proto galactic gas) have been studied by different authors (e.g, de Zotti et al. 2004, Rosa-Gonz'alez et al. 2004, Massardi et al. 2008 ). Here we investigate one generic class of SZ signals: the hot bubble surrounding a quasar powered by a supermassive black hole. Probing black hole energy feedback via SZ distortions is one direct observational route to understanding the growth and evolution of supermassive black holes and their role in structure formation.  Analytic studies of this signal have been done by several authors (e.g., Natarajan \\& Sigurdsson 1999; Yamada, Sugiyama \\& Silk 1999; Lapi, Cavaliere \\& De Zotti 2003; Platania et al.\\ 2002; Chatterjee \\& Kosowsky 2007); the current numerical work complements a similar study by Scannapieco Thacker and Couchman 2008. \\begin{table*} \\begin{tabular}[t]{c|c|c|c|c|c|c} \\hline \\multicolumn{1}{c|}{Run}& \\multicolumn{1}{c|}{Boxsize}& \\multicolumn{1}{c|}{$N_{P}$}& \\multicolumn{1}{c|}{$m_{DM}$}& \\multicolumn{1}{c|}{$m_{gas}$}& \\multicolumn{1}{c|}{$\\epsilon$}& \\multicolumn{1}{c}{$z_{end}$}\\\\ \\multicolumn{1}{c|}{}& \\multicolumn{1}{c|}{($h^{-1}$ Mpc)}& \\multicolumn{1}{c|}{}& \\multicolumn{1}{c|}{($h^{-1}M\\odot$)}& \\multicolumn{1}{c|}{($h^{-1}M\\odot$)}& \\multicolumn{1}{c|}{($h^{-1}$ kpc)}& \\multicolumn{1}{c}{}\\\\ \\hline D4 & 33.75 & $2\\times216^{3}$ & $2.75\\times10^{8}$ & $4.24\\times10^{7}$ & 6.25 & 0.00 \\\\ D6 (BHCosmo) & 33.75 & $2\\times486^{3}$ & $2.75\\times10^{7}$ & $4.24\\times10^{6}$ & 2.73 & 1.00 \\\\ \\hline \\end{tabular} \\caption{The numerical parameters in the simulation. For the current study we have used the low-resolution version because we have a matching simulation with no black holes; resolution effects are discussed in Sec.~4.3. $N_{p}$, $m_{DM}$, $m_{gas}$, $ \\epsilon $ and $z_{end}$ are defined as the total number of particles, mass of the dark matter particles, mass of the gas particles, gravitational softening length and final redshift run respectively.}  \\end{table*}  Analytic models and numerical simulations of galaxy cluster formation indicate that the temperature and the X-ray luminosity relation should be related as $L_{x} \\simeq T^{2}$ in the absence of gas cooling and heating (Peterson \\& Fabian 2006 and references therein). Observations show instead that $L_{x} \\simeq T^{3}$ over the temperature range 2 to 8 kev with a wide dispersion at lower temperature, and a possible flattening above (Markevitch 1998; Arnaud \\& Evrard 1999; Peterson \\& Fabian 2006). The simplest explanation for this result is that the gas had an additional heating of 2 to 3 keV per particle (Wu, Fabian \\& Nulsen 2000; Voit et al.\\ 2003). Several nongravitational heating sources have been discussed in this context (Peterson \\& Fabian 2006; Morandi Ettori \\& Moscardini 2007); quasar feedback (e.g., Binney \\& Tabor 1995; Silk \\& Rees 1998; Ciotti \\& Ostriker 2001; Nath \\& Roychowdhury 2002; Kaiser \\& Binney 2003; Nulsen et al. 2004) is perhaps the most realistic possibility. The effect of this feedback mechanism on different scales of structure formation have been addressed by several authors (e.g., Mo \\& Mao 2002; Oh \\& Benson 2003; Granato et al. 2004). The mechanism of quasar heating in cluster cores has been observationally motivated by studies from McNamara et al. 2005, Voit \\& Donahue 2005 and Sanderson, Ponman \\& O'Sullivan 2006 (see McNamara \\& Nulsen 2007 for a recent review). The impact of this nongravitational heating in galaxy groups, which have shallower potential wells and thus smaller intrinsic thermal energy than galaxy clusters, can also be substantial (Arnaud \\& Evrard 1999; Helsdon \\& Ponman 2000; Lapi, Cavaliere \\& Menci 2005). Observational efforts to detect the impact of this additional heating source in the context of quasar feedback have been carried out using galaxy groups in the Sloan Digital Sky Survey (SDSS) by Weinmann et al.\\ 2006, and with a Chandra group sample by Sanderson, Ponman \\& O'Sullivan 2006. Detailed theoretical studies of galaxy groups using simulations which include quasar feedback have been undertaken by, e.g., Zanni et al.\\ 2005, Sijacki et al.\\ 2007, and Bhattacharya, Di Matteo \\& Kosowsky 2007. At smaller scales, the impact of quasar feedback has been investigated by Schawinski et al.\\ 2007 with early-type galaxies in SDSS, and has also been studied in several theoretical models of galaxy evolution (e.g, Kawata \\& Gibson 2005; Bower et al. 2006; Cattaneo et al.\\ 2007).  Growing observational evidence points to a close connection between the formation and evolution of galaxies, their central supermassive black holes (e.g., Magorrian et al.\\ 1998, Ferrarese \\& Merritt 2000, Tremaine et al.\\ 2002) and their host dark matter halos (Merritt \\& Ferrarese 2001; Tremaine et al.\\ 2002).  Several different groups have now investigated black hole growth and the effects of quasar feedback in the cosmological context (e.g., Scannapieco \\& Oh 2004; Di Matteo, Springel \\& Hernquist 2005; Lapi et al. 2006; Croton et al.\\ 2006; Thacker, Scannapieco \\& Couchman 2006, Sijacki et al.\\ 2007). Recently Di Matteo et al.\\ (2008) carried out a hydrodynamic cosmological simulation following in detail the growth of supermassive black holes, using a simple but realistic model of gas accretion and associated feedback. Using this simulation, we construct maps of the SZ distortion around black holes of different masses at various redshifts. We demonstrate that the SZ signal around quasars scales with black hole mass and accretion rate. A similar approach has been taken by Scannapieco, Thacker \\& Couchman (2008) from a cosmological simulation with a different implementation of quasar feedback. Predictions for the SZ distortion due to a phenomenological treatment of galactic winds from simulations have been obtained previously by White, Hernquist \\& Springel (2002).  \\begin{figure*} \\begin{center} \\begin{tabular}{cc}  \\resizebox{55mm}{!}{\\includegraphics{bh8582500feed1.eps}} \\resizebox{55mm}{!}{\\includegraphics{bh11152500feed1.eps}} \\resizebox{55mm}{!}{\\includegraphics{bh12662500feed1.eps}}\\\\ \\resizebox{55mm}{!}{\\includegraphics{bh858100feed1.eps}} \\resizebox{55mm}{!}{\\includegraphics{bh1115100feed1.eps}} \\resizebox{55mm}{!}{\\includegraphics{bh1266100feed1.eps}}\\\\  \\resizebox{55mm}{!}{\\includegraphics{bh3892500feed1.eps}} \\resizebox{55mm}{!}{\\includegraphics{bh19692500feed1.eps}} \\resizebox{55mm}{!}{\\includegraphics{bh22552500feed1.eps}}\\\\ \\resizebox{55mm}{!}{\\includegraphics{bh389100feed1.eps}} \\resizebox{55mm}{!}{\\includegraphics{bh1969100feed1.eps}} \\resizebox{55mm}{!}{\\includegraphics{bh2255100feed1.eps}}\\\\  \\end{tabular} \\caption{Simulated $y$-distortion maps around two massive black holes, at three different redshifts; left, middle, and right columns are for $z=3$, $z=2$, and $z=1$. The top two rows are for the most massive black hole in the simulation, with mass  $7.35 \\times 10^{8} M_{\\odot}$, $2.76 \\times 10^{9} M_{\\odot}$ and $4.26 \\times 10^{9}M_{\\odot}$ at redshifts 3, 2, and 1 respectively. The top row shows $y$ in a 5 Mpc square region centered on the black hole; the second row shows zooms in to a smaller region 200 kpc square. Note the two rows have different color scales. The third and fourth rows are the same as the first two rows, for a different black hole (the second most massive black hole at redshift 3.0) with mass $7.15 \\times 10^{8} M_{\\odot}$, $8.2 \\times 10^{8} M_{\\odot}$ and $2.11 \\times 10^{9}M_{\\odot}$ at redshifts 3, 2, and 1. For both black holes, the peak value of $y$ is between $10^{-7}$ and $10^{-6}$, corresponding to an effective maximum temperature distortion between a few tenths of a $\\mu$K to a few $\\mu$K.} \\end{center} \\end{figure*}   Typically, the SZ distortions from quasar feedback result in effective temperature distortions at the micro-Kelvin level on arcminute angular scales. Recent advances in millimeter-wave detector technology and the construction of several single-dish and interferometric experiments (see Birkinshaw \\& Lancaster 2007 for a recent review on SZ observations), including the Sub Millimeter Array (SMA), the combined array for research in Millimeter wave Astronomy (CARMA), the Cornell Caltech Atacama Telescope (CCAT), the Atacama Large Millimeter Array (ALMA), and the Large Millimeter Telescope (LMT), along with SZ surveys like ACT and SPT, have brought detection of this signal into the realm of possibility. Although the direct detection of this signal from current SZ surveys seems unlikely, since the amplitude of fluctuation observed is at or below the noise threshold of ACT or SPT (Chatterjee \\& Kosowsky 2007), proposed submillimeter facilities offer some possibility for direct detection of this signal.  An additional route for detection is cross-correlation of optically-selected quasar with microwave maps (Chatterjee \\& Kosowsky 2007; Scannapieco Thacker \\& Couchman 2008). However, this approach likely requires multifrequency observations to discriminate the SZ effect from intrinsic quasar emission or infrared sources.  The paper is organized as follows. Section 2 describes the simulation that has been used in this work. In Section 3 we give a brief review of the SZ distortion and display the SZ maps derived from the simulation. Astrophysical results from the maps are presented in Section 4, including radial profiles around individual black holes and the correlation between black hole mass and SZ distortion. The concluding Section estimates detectability of these signals and summarizes future prospects.  Throughout we use units with $c=k_B=1$.  \\begin{table*} \\begin{tabular}[t]{c|c|c|c|c|c} \\hline \\hline \\multicolumn{1}{c}{}& \\multicolumn{1}{c}{Most massive black hole}& \\multicolumn{1}{c|}{}& \\multicolumn{1}{c}{}& \\multicolumn{1}{c}{Second black hole}& \\multicolumn{1}{c|}{}\\\\ \\hline  \\multicolumn{1}{c|}{Redshift}& \\multicolumn{1}{c|}{$N_{BH}$}& \\multicolumn{1}{c|}{$dM/dt$, $M_\\odot$/yr}&  \\multicolumn{1}{c|}{Redshift}& \\multicolumn{1}{c|}{$N_{BH}$}& \\multicolumn{1}{c|}{$dM/dt$, $M_\\odot$/yr}  \\\\ \\multicolumn{1}{c|}{}& \\multicolumn{1}{c|}{}& \\multicolumn{1}{c|}{}& \\multicolumn{1}{c|}{}& \\multicolumn{1}{c|}{}& \\multicolumn{1}{c|}{}   \\\\ \\hline \\hline 3.0 & 0 & 0.034 & 3.0 & 0 & 0.240  \\\\ 2.0 & 3 & 0.003 & 2.0 & 2 & 0.013 \\\\ 1.0 & 4 & 0.013 &  1.0 & 1 & 0.005\\\\ \\hline \\end{tabular}  \\caption{The accretion rates and the number of neighboring black holes within a radius of 100 kpc, for the two black holes in Fig.~1.} \\end{table*}    \\begin{figure*} \\begin{center} \\begin{tabular}{cc}  \\resizebox{60mm}{!}{\\includegraphics{bh858diffszlog1.eps}} \\resizebox{60mm}{!}{\\includegraphics{bh858difftemplog1.eps}} \\resizebox{60mm}{!}{\\includegraphics{bh858diffdenlog1.eps}}\\\\ \\resizebox{60mm}{!}{\\includegraphics{bh1266diffszlog1.eps}} \\resizebox{60mm}{!}{\\includegraphics{bh1266difftemplog1.eps}} \\resizebox{60mm}{!}{\\includegraphics{bh1266diffdenlog1.eps}}\\\\ \\resizebox{60mm}{!}{\\includegraphics{bh389diffszlog1.eps}} \\resizebox{60mm}{!}{\\includegraphics{bh389difftemplog1.eps}} \\resizebox{60mm}{!}{\\includegraphics{bh389diffdenlog1.eps}}\\\\ \\resizebox{60mm}{!}{\\includegraphics{bh2255diffszlog1.eps}} \\resizebox{60mm}{!}{\\includegraphics{bh2255difftemplog1.eps}} \\resizebox{60mm}{!}{\\includegraphics{bh2255diffdenlog1.eps}}\\\\  \\end{tabular} \\caption{The difference in $y$-distortion between a simulation with black hole feedback and a simulation without, for the same region of space shown in Fig.~1.  The two simulations have identical resolution and initial conditions. The first row corresponds to the most massive black hole at $z=3$, the second row at $z=1$; the third row corresponds to the other black hole (second most massive black hole at redshift 3.0) in Fig.~1 at $z=3$, the fourth row at $z=1$. The left column shows $y$, the middle shows the log of the mass-weighted average temperature in units of Kelvin. The right column shows the log of the electron number surface density in units of cm$^{-2}$ .} \\end{center} \\end{figure*}    \\begin{figure*}  \\begin{tabular}{cc} \\resizebox{90mm}{!}{\\includegraphics{diffangprof1.eps}} \\resizebox{90mm}{!}{\\includegraphics{diffangprof2.eps}}\\\\  \\end{tabular} \\caption{The angular profiles of the $y$ distortion for the two black holes shown in Fig.~1 at three different redshifts. The solid, dot-dashed and dashed lines are for redshifts 3, 2, and 1 respectively. The left panel shows the most massive black hole and the right panel is the other black hole (second most massive black hole at redshift 3.0).  Although the central amplitude at a particular time depends strongly on the instantaneous state of the black hole accretion, the distortion amplitude increases monotonically with time at a 20 arcsecond angular distance. } \\end{figure*} ",
        "conclusions": "Observationally, quasar feedback is directly detectable by resolving Sunyaev-Zeldovich peaks on small angular scales of tens of arcseconds with amplitudes of up to a few $\\mu$K above the immediately surrounding region. The combination of angular scale and small amplitude make detecting this effect very challenging, at the margins of currently planned experiments.  The necessary sensitivity requires large collecting areas, while the angular resolution needed points to an interferometer in a compact configuration, or a large single-dish experiment.  Since the SZ signal is manifested as a peak over the surrounding background level, a region substantially larger than the SZ peak must be imaged. This requires a telescope having sufficient resolution to resolve the central peak in the SZ distortion in an SZ image and enough field of view so that the peak could be identified. An example is the compact ALMA subarray known as the Atacama Compact Array (ACA), composed of 12 7-meter dishes. The ALMA sensitivity calculator gives that the synthesized beam for this array is about 14 arcseconds, and the integration time required to attain 1 $\\mu$K sensitivity per beam at a frequency of 145 GHz and a maximum band width of 16 GHz is on the order of 1000 hours (ALMA sensitivity calculator).  A very deep survey with this instrument could detect the SZ effect from individual black holes. The 50-meter Large Millimeter-Wave Telescope instrumented with the AzTEC bolometer array detector will have a somewhat similar sensitivity but detectibility would require a very deep (thousands of hours) integration time. The Cornell-Caltech Atacama Telescope (CCAT), a 25-meter telescope, estimates a possible pixel sensitivity for SZ detection at 150 GHz of 310 $\\mu{\\rm K}\\,{\\rm s}^{1/2}$ for 26 arcsecond pixels, so a 30 hour observation could give 1$\\mu$K pixel noise. These pixels would not be small enough to resolve the hot halo around a black hole, but might be able to detect the difference in a single pixel due to black hole emission compared to the surrounding pixels. Aside from raw sensitivity and angular resolution, a serious difficulty with direct detection is the confusion limit from infrared point source emission; these sources are generally high-redshift star forming galaxies with a high dust emission. CCAT estimates show that their one-source-per-beam confusion limit will be around 6 $\\mu$K at 150 GHz (Golwala 2006). This will present substantial difficulties for detecting a 1 $\\mu$K temperature distortion if accurate. It is noted that the observations in the sub-millimeter band is limited by confusion noise and so another possibility of direct detection of the signal through radio frequency telescopes could be considered. Massardi et al. 2008 shows that the confusion due to dusty galaxies is lower at 10 GHz then at 100 GHz. The authors show that for galactic scale SZ effect the optimal frequency range for detection is between 10 to 35 GHz. However substantial confusion from radio galaxies at these low frequency observations would still be a challenging issue in the direct detection of the signal.   Given these substantial difficulties associated with direct detection, an alternate route may be necessary. Cross-correlation of arcminute-resolution microwave maps with optically selected quasar or massive galaxies is a second possible detection strategy (Chatterjee and Kosowsky 2007, Scannapieco, Thacker, and Couchman 2008). By averaging over large numbers of objects, we can have an estimate of a small mean black hole distortion signal from the noise in the maps.  The primary challenge with this technique is the direct emission from quasar in the microwave band. It may be possible to select a sample of quasar which is sufficiently radio-quiet that the cross-correlation is not dominated by the intrinsic emission. Another possibility is to select massive field galaxies under the assumption that they harbor a central massive black hole which at one time was active; the hot bubble produced has a cooling time comparable to the Hubble time, so formerly active galaxies should still have an SZ signature. Finally, the SZ effect from black holes is in addition to the SZ emission from any hot gas in which the black hole's host galaxy is embedded. Massive galaxies trace large-scale structure, and any cross-correlation will also detect this signal. Although the observational requirements for the cross-correlation method are plausible the scopes for detectibility with this method is still limited by confusion noise. Stacking microwave (SZ) maps in the direction of known quasars would also serve as an independent route in detecting the signal (Chatterjee \\& Kosowsky 2007). This can improve the signal to noise by a substantial amount although this method would still be limited by the uncertainties described above. Quantifying in detail the observable signal (which will need to be disentangled from other confusions such as dusty galaxies, radio galaxies etc.) for the possible direct detections methods or from cross-correlation analysis that we have discussed is beyond the scope of this paper and we defer it to a future work. The simulations and maps presented here provide a basis for further modeling of all these effects.  The main conclusions drawn from this work are summarized as follows. We have used the first cosmological simulations to incorporate realistic black hole growth and feedback to produce simulated maps of the Sunyaev-Zeldovich distortion of the microwave background due to the feedback energy from accretion onto supermassive black holes. These simulations address the rapid accretion phases of black holes: periods of strong emission are typically short-lived and require galaxy mergers to produce strong gravitational tidal forcing necessary for sufficient nuclear gas inflow rates (Hopkins, Narayan \\& Hernquist 2006; Di Matteo et al. 2008). The result is heating of the gas surrounding the black hole, so that the largest black holes produce a surrounding hot region which induces a $y$-distortion (related to a temperature distortion) with a characteristic amplitude of a few $\\mu$K. We have obtained a scaling relation between the black hole mass and their SZ temperature decrement, which in turn is a measure of the amount of feedback energy output. The correspondence between the y distortion and the accretion rates is not exact but there is a close association which shows the correlation between feedback output and black hole activity. From our results we have shown that with the turn on of AGN feedback the signal gets enhanced largely and the enhancement is predominant at angular scales of 5 arcseconds. Finally we have shown that there is a fair probability of detecting this signal even from the planned sub millimeter missions.  \\par The role of energy feedback from quasars and from star formation is known to have substantial impact on the process of galaxy formation and evolution of the intergalactic medium, but the details of this process are not well understood. Probes based on Sunyaev-Zeldovich distortions are challenging, but an eventual detection can be used to put useful constraints and checks on models of AGN feedback."
    },
    "0808/0808.0652_arXiv.txt": {
        "abstract": "This paper presents the results of a Fresnel Interferometric Array testbed. This new concept of imager involves diffraction focussing by a thin foil, in which many thousands of punched subapertures form a pattern related to a Fresnel zone plate. This kind of array is intended for use in space, as a way to realizing lightweight large apertures for high angular resolution and high dynamic range observations. The chromaticity due to diffraction focussing is corrected by a small diffractive achromatizer placed close to the focal plane of the array.   The laboratory test results presented here are obtained with an 8 centimeter side orthogonal array, yielding a 23 meter focal length at 600 nm wavelength. The primary array and the focal optics have been designed and assembled in our lab. This system forms an achromatic image. Test targets of various shapes, sizes, dynamic ranges and intensities have been imaged. We present the first images, the achieved dynamic range, and the angular resolution. ",
        "introduction": "\\label{sec:intro} Fresnel arrays are interferometric devices involving many hundreds or thousands of \"basic\" subapertures: mere holes punched in a large and thin foil. For each of these individual subapertures, diffraction occurs. The subapertures positioning law is designed so that the resulting interference produces the desired wavefront shape: for instance a spherical wavefront. The circular Fresnel Zone Plate (also called Soret zone plate), constituted of alternatively opaque and transparent concentric rings, is an example of Fresnel array: the distance from one transmissive ring to the other is adjusted so that the optical path to a point named \"focus\" differs of one $\\lambda$ from one ring to the other. As a result constructive interference occurs at this point, resulting in an equivalent of focalisation: concentration of light.  In our approach, the orthogonal geometry has been preferred to Soret's concentric rings for two main reasons: its allowance of vacuum rather than a transparent material for transmissive zones, and the higher dynamic range allowed in most of the image field. This optical scheme strongly reduces manufacturing and positioning constraints compared to those of standard interferometric arrays and monolithic mirrors. It also yields high dynamic range as well as larger field-resolution ratios than classical imaging interferometers. We propose this concept for high angular resolution imaging in future space projects (Koechlin et al 2005 \\cite{Koechlin_aa_2005}) .   An orthogonal Fresnel array of side $c$ involving $N$ Fresnel zones results in a focal length $f$ at its first order of interference, which can be written \\cite{Koechlin_aa_2005}: \\begin{equation} \\centerline{$f = \\frac{c^2}{8N}\\,\\frac{1}{\\lambda}$ .} \\label{eq:lambda_f} \\end{equation} The unavoidable chromaticity of this kind of focussing element can be corrected using an optical scheme based on Schupmann studies \\cite{Schupmann_1899} , and emerging from the chromatic correction of holograms developed by Faklis \\& Morris in 1989 \\cite{Faklis_oe_1989} . The principle is to place in a pupil plane an optical device whose chromatic dispersion is opposite to that of the Fresnel array. That means in our case to use a diffractive optical device used at the order -1: a diverging Fresnel zone lens whose focal length varies in $-\\frac{1}{\\lambda}$. Fig.\\ref{fig:achromatisation_scheme} presents the optical scheme.  \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[width=0.7\\linewidth]{figs/schema_acchrom.png} \\end{tabular} \\end{center} \\caption[Achromatisation scheme] { \\label{fig:achromatisation_scheme} An orthogonal Fresnel array is placed on plane (1) and used at its first order of interference. Different wavelengths (dashed and dotted lines) are focussed at different distances from the Fresnel array. A field optics (2) images the Fresnel array onto a pupil plane (3), where a diverging Fresnel zone lens is placed. The emerging beam is perfectly achromatic, but needs a classical re-imaging optics to be convergent (4). The final achromatic image is formed on plane (5).} \\end{figure}  Following [Koechlin et al 2004 \\cite{Koechlin_spie_2004}] and [Serre et al 2005 \\cite{Serre_jithd_2006}] , and thanks to a CNES{\\footnote {Centre National d'Etudes Spatiales% }} R\\&T funding, we developed in the past two years at Universit\\'e Paul Sabatier a Fresnel Interferometric Array breadboard testbed. The main focussing element is an 8 cm side square array and it is associated to optimized focal optics, leading to achromatic images and allowing diffraction limited performances. We present in this paper: firstly the constitutive elements of this prototype, and secondly the measured performances. ",
        "conclusions": "We have presented in this paper the first results of an orthogonal Fresnel interferometric array breadboard testbed. We are able to confirm the efficiency of the achromatisation scheme retained, as well as the diffraction-limited performances and the \"wide\" field imaging capabilities. The measured rejection rates are also in good agreement with computer generated previsions.\\\\ The next steps using this testbed are to evaluate rejection rates with two sources of different luminosity in the same field, to test different manufacturing processes for the blazed diverging Fresnel zone lens, and to assess Fresnel arrays subaperture shapes and layout, allowing higher transmission rates and dynamic ranges.\\\\ In the next few years, on one hand members of our team (and we hope new people !) aim to evaluate a 2$^{\\text{nd}}$ generation of ground-based testbed: its main evolutions will be an increase in dimensions (typically 20 cm side for the Fresnel array), higher number of Fresnel zones, and last but not least: real sky observation. On the other hand, studies will keep on being done to evaluate the performances of this kind of instrument and the potential astrophysical targets of a space-based Fresnel interferometric array concept. This space observatory could have a front array dimension of a few meters or more, and acquire data in the UV, visible or/and IR domains.\\\\  \\footnotesize{This work was supported by Universit\\'{e} Paul Sabatier Toulouse III, CNRS, CNES, Thal\\`{e}s Alenia Space and the Fonds Social Europ\\'{e}en.}"
    },
    "0808/0808.0839_arXiv.txt": {
        "abstract": "In this work, we consider time-dependent dark energy star models, with an evolving parameter $\\omega$ crossing the phantom divide, $\\omega=-1$. Once in the phantom regime, the null energy condition is violated, which physically implies that the negative radial pressure exceeds the energy density. Therefore, an enormous negative pressure in the center may, in principle, imply a topology change, consequently opening up a tunnel and converting the dark energy star into a wormhole. The criteria for this topology change are discussed and, in particular, we consider a Casimir energy approach involving quasi-local energy difference calculations that may reflect or measure the occurrence of a topology change. We denote these exotic geometries consisting of dark energy stars (in the phantom regime) and phantom wormholes as \\textit{phantom stars}. The final product of this topological change, namely, phantom wormholes, have far-reaching physical and cosmological implications, as in addition to being used for interstellar shortcuts, an absurdly advanced civilization may manipulate these geometries to induce closed timelike curves, consequently violating causality. ",
        "introduction": "Recent high-precision observational data have confirmed that the Universe is undergoing a phase of accelerated expansion \\cite{expansion}. Several candidates, responsible for this expansion, have been proposed in the literature, in particular, dark energy models (see Ref. \\cite{Copeland} for a review) and modified gravity (e.g., see Refs. \\cite{modgravity} for recent reviews). In particular, the former models are fundamental candidates, in which a simple way to parameterize the dark energy is by an equation of state of the form $\\omega\\equiv p/\\rho$, where $p$ is the spatially homogeneous pressure and $\\rho$ is the dark energy density. A value of $\\omega<-1/3$ is required for cosmic expansion, and $\\omega=-1$ corresponds to a cosmological constant. A specific exotic form of dark energy denoted phantom energy, with $\\omega<-1$, has also been proposed \\cite{Caldwell}, and possesses peculiar properties, such as the violation of the null energy condition (NEC) and the energy density increases to infinity in a finite time \\cite{Caldwell}, at which point the size of the Universe blows up in a finite time, which is known as the Big Rip. In this context, the violation of the NEC presents us with a natural scenario for the existence of traversable wormholes, and indeed it has been shown that these exotic geometries can be supported by phantom energy \\cite{Sushkov,Lobo-phantom}. It is also interesting to note that recent fits to supernovae, CMB and weak gravitational lensing data probably favor an evolving equation of state, with the parameter crossing the phantom divide $\\omega=-1$ \\cite{Vikman}.  Despite the fact that the dark energy equation of state represents a spatially homogeneous cosmic fluid and is assumed not to cluster, it is possible that inhomogeneities may arise due to gravitational instabilities. More precisely, although the equation of state leading to the acceleration of the Universe on large scales is an average equation of state corresponding to a background fluid, it is possible that dark energy condensates may possibly originate from density fluctuations in the cosmological background, resulting in the nucleation through the respective density perturbations. Despite the fact that once in the dark energy regime the material system becomes gravitationally repulsive, we may consider the possibility of the formation of a matter system that originally obeys all the energy conditions. Cosmological observations do not rule out, and in some studies favor, an evolving equation of state for the dark energy. It is therefore quite possible that what we know as dark energy today has evolved from a more benign fluid. An over-density of this fluid could in principle commence a collapse into a star which. Such a model is presented in section \\ref{sec:IIIC}. We also point out that even in the case of a dark-energy fluid, there is no definite resolution to the debate of clustering scales. This is mainly due to non-linearity, especially in the vein of dark energy interacting with ordinary fluids. It may also be possible to glean some information on the cosmological dark matter by studying certain properties of such gravitational condensates. (See \\cite{ref:aandaandijtp} and references therein for comments on these issues.) In this context, a number of inhomogeneous solutions have been the object of analysis, such as the phantom wormholes \\cite{Sushkov,Lobo-phantom} mentioned above, dark energy stars \\cite{Lobo:2005uf}, and other structures such as condensates supported by the generalized Chaplygin gas \\cite{Chapgas} which possibly arise from density fluctuations in the generalized Chaplygin gas background, and condensed structures supported by the van der Waals equation of state \\cite{Lobo:2006ue}% . In a recent paper \\cite{Dzhunushaliev:2008bq}, it was also shown that the $4D$ Einstein-Klein-Gordon equations with a phantom scalar field possess non-singular, spherically symmetry solutions, although a stability analysis on these solutions indicates they are unstable.  The dark energy star models are also a generalization of a new emerging picture for an alternative final state of gravitational collapse, namely, the gravastar (\\textit{grav}itational \\textit{va}cuum \\textit{star}) models. The latter proposed by Mazur and Mottola \\cite{Mazur}, has an effective phase transition at/near where the event horizon is expected to form, and the interior is replaced by a de Sitter condensate. The latter is then matched to a thick layer, with an equation of state given by $p=\\rho$, which is in turn matched to an exterior Schwarzschild solution. The issue of gravastars has been extensively analyzed in the literature, and we refer the reader to Refs. \\cite{gravastar2,ref:contgrava}. The generalization of the gravastar picture is considered by matching an interior solution governed by the dark energy equation of state, $\\omega\\equiv p/ \\rho<-1/3$, to an exterior Schwarzschild vacuum solution at a junction interface \\cite{Lobo:2005uf}. The dynamical stability of the transition layer was also explored, and it was found that large stability regions exist that are sufficiently close to where the event horizon is expected to form, so that it was argued that it would be difficult to distinguish the exterior geometry of the dark energy stars from an astrophysical black hole. Thus, these alternative models do not possess a singularity at the origin and have no event horizon, as its rigid surface is located at a radius slightly greater than the Schwarzschild radius. This restriction arises from the observed lack of energy emission due to surface collisions of infalling material in suspected black hole systems. In fact, although evidence for the existence of black holes is very convincing, a certain amount of scepticism regarding the physical reality of event horizons is still encountered, and it has been argued that despite the fact that observational data do indeed provide strong arguments in favor of event horizons, they cannot fundamentally prove their existence \\cite{AKL}.  As mentioned above, recent fits to observational data probably favor an evolving equation of state, with the dark energy parameter crossing the phantom divide $\\omega=-1$ \\cite{Vikman}. Motivated by this fact, in a rather speculative scenario one may theoretically consider the existence of a dark energy star, with an evolving parameter starting out in the range $-1<\\omega<-1/3$, and crossing the phantom divide, $\\omega=-1$. Once in the phantom regime, the null energy condition is violated, which physically implies that the negative radial pressure exceeds the energy density. Therefore, an enormous negative pressure in the center may, in principle, imply a topology change, consequently opening up a tunnel, and converting the dark energy star into a wormhole \\cite{Lobo:2005uf,Morris}. One may assume that the topology change may occur at approximately the Planck length scales, and once created may be self-sustained as shown in Ref. \\cite{Garattini:2007ff}. In fact, the change in topology is an extremely subtle issue, as in general relativity these changes probably entail spacetime singularities. However, at the Planck length scales quantum gravity effects dominate and spacetime undergoes a deep and rapid transformation in its structure, probably producing a multiply-connected quantum foam structure \\cite{Wheeler,geons}. It was suggested in Ref. \\cite{Morris} that one could imagine an absurdly advanced civilization \\cite{Lemos:2003jb} pulling a wormhole from this submicroscopic spacetime quantum foam and enlarging it to macroscopic dimensions. However, in a more plausible scenario, the possibility that inflation might provide a natural mechanism for the enlargement of such wormholes to macroscopic size was explored \\cite{Roman:1992xj}. In this work, we outline the theoretical difficulties associated to the change in topology and present a method based on the Casimir energy approach. Although it is still unsure if this method produces a topology change, it is extremely useful as the quasi-local energy difference calculation may reflect or measure the occurrence of a change in topology. Other concepts of topology changing spacetimes have been studied, for instance: Using semi-classical and Morse-index methods in Refs. \\cite{Morseindex}; higher order back-reaction terms due to fluctuations of gauge fields in the vicinity of a black hole may result in the formation of a wormhole-like object \\cite{Hochberg:1992rd}; and more recently an approach based on a Ricci flow may result in quantum wormholes \\cite{Dzhunushaliev:2008cz}.  Once the topology change has occurred, with the respective opening of a tunnel, then the dark energy star has been converted into wormhole supported by phantom energy. As mentioned above, it has recently been shown that traversable wormholes may, in principle, be supported by phantom energy \\cite{Sushkov,Lobo-phantom}, which apart from being used as interstellar shortcuts, may induce closed timelike curves with the associated causality violations \\cite{mty,Visser}. Particularly interesting solutions were found \\cite{Lobo-phantom}, and by using the ``volume integral quantifier'', it was found that these wormhole geometries are, in principle, sustained by arbitrarily small amounts of averaged null energy condition (ANEC) violating phantom energy. A complementary approach was traced out in \\cite{Sushkov}, by considering specific choices for the distribution of the energy density threading the wormhole. Recently, $4D$ static wormhole solutions supported by two interacting phantom fields were found as well \\cite{Dzhunushaliev:2007cs}. Despite the fact that traversable wormholes violate the NEC in general relativity (see Ref. \\cite{Lobo:2007zb} for a recent review), it has been shown that the stress energy tensor profile may satisfy the energy conditions in the throat neighborhood in dynamic wormholes (see Ref. \\cite{Arellano:2006ex} and references therein) and in certain alternative theories to general relativity \\cite{satisfyNEC}. Perhaps not so appealing, one could denote these exotic geometries consisting of dark energy stars (in the phantom regime) and phantom wormholes as \\textit{phantom stars}. We would like to state our agnostic position relatively to the existence of dark energy stars and phantom wormholes, or for that matter of \\textit{phantom stars}. However, it is important to understand their general properties and characteristics, and we emphasize that the presence of a dark energy fluid permeating the universe makes the study of dark energy condensates a physically relevant endeavor.  This paper is organized in the following manner: In section \\ref{secII}, we briefly review static dark energy stars, followed by a deduction of general solutions of time-dependent spacetimes. In Section \\ref{secIII}, specific time-dependent dark energy solutions are outlined, in particular, we present the specific cases of a constant energy density, the Tolman-Matese-Whitman mass function solution, and a class of models with a non-zero energy flux term, which form from gravitational collapse. In Section \\ref{secIV}, we describe the theoretical difficulties associated with changes in topology and present in some detail specific methods used in the literature, namely, a Casimir energy approach involving quasi-local energy difference calculations that may reflect or measure the occurrence of a topology change. We also briefly review the Morse index analysis. In section \\ref{sec:conclusion}, we conclude. ",
        "conclusions": "\\label{sec:conclusion}   In this work, we have considered time-dependent dark energy star models, with an evolving parameter $\\omega$ crossing the phantom divide, $\\omega=-1$. In particular, we briefly reviewed static and spherically symmetric dark energy stars, and further analyzed general solutions of time-dependent spacetimes in detail. Specific time-dependent solutions were extensively explored, in particular, the specific cases of a constant energy density, the Tolman-Matese-Whitman mass function solution, and a class of models with a non-zero energy flux term, which form from gravitational collapse. Once the parameter $\\omega$ evolves into the phantom regime, the null energy condition is violated, which physically implies that the negative radial pressure exceeds the energy density. Therefore, an enormous negative pressure at the center may, in principle, imply a topology change, consequently opening up a tunnel and converting the dark energy star into a wormhole. The theoretical difficulties and criteria for this topology change were discussed in detail, where in particular we considered a Casimir energy approach involving quasi-local energy difference calculations that may reflect or measure the occurrence of a topology change. Once the topology change has occurred, it is possible that the resulting wormhole structures, supported by phantom energy, be self-sustained. As mentioned in the Introduction, recent fits to observational data probably favor an evolving equation of state, with the dark energy parameter crossing the phantom divide $\\omega=-1$ \\cite{Vikman}. However, in a cosmological setting the transition into the phantom regime is physically implausible for a single scalar field \\cite{Vikman}, so that a possible approach would be to consider a mixture of interacting non-ideal fluids. One may consider that the time-dependent dark energy star model outlined in this work, is a simplification of this possible approach. In fact, recently, static models with two interacting phantom and ghost scalar fields were considered, and it was shown that regular solutions exist \\cite{Dzhunushaliev:2007cs}. It would be interesting to generalize the latter study to time-dependent solutions, extending the analysis considered in this work.  It is interesting to note that the topology change at the center should influence the surface stresses at the thin shell, as there is a redistribution of the stress-energy tensor components of the interior solution during the change in topology. That this is so may be verified through the conservation identity given by $S^{i}_{j|i}=[T_{\\mu\\nu}e^{\\mu}_{(j)}n^{\\nu}]^{+}_{-}\\,$, where $[X]^{+}_{-}$ denotes the discontinuity across the surface interface $\\Sigma$, i.e., $[X]^{+}_{-}=X^{+}|_{\\Sigma}-X^{-}|_{\\Sigma}$. The quantity $S^{i}_{j}$ is the surface stress-energy tensor at the junction surface $\\Sigma$; $n^{\\mu}$ is the unit normal $4-$vector to $\\Sigma$; and $e^{\\mu }_{(i)}$ are the components of the holonomic basis vectors tangent to $\\Sigma$ (see Refs. \\cite{WHshell} for details). Note the dependency of the conservation identity on the stress-energy tensor $T_{\\mu\\nu}$, and the right hand side of the conservation identity may also be written as $S^{i}_{\\tau |i}=-\\left[  \\dot{\\sigma}+2\\dot{a}(\\sigma+\\mathcal{P} )/a \\right]  $. The momentum flux term, i.e., $[T_{\\mu\\nu}e^{\\mu}_{(j)}n^{\\nu}]^{+}_{-}\\,$, corresponds to the net discontinuity in the momentum flux $F_{\\mu}=T_{\\mu\\nu }\\,U^{\\nu}$ which impinges on the shell. The conservation identity is a statement that all energy and momentum that plunges into the thin shell, gets caught by the latter and converts into conserved energy and momentum of the surface stresses of the junction. Now, it may be that the topology change is sufficiently violent to disrupture the thin shell. On the other hand, one may also assume that it is sufficiently mild as not to significantly affect the stability of the surface layer.  In analogy to the case outlined in Ref. \\cite{Roman:1992xj}, where the possibility that inflation might provide a natural mechanism for the enlargement of wormholes to macroscopic size was explored, one could imagine that microscopic wormholes originated through a topology change, and due to the accelerated expansion of the Universe, these submicroscopic constructions could naturally be grown to macroscopic dimensions. For instance, in Ref. \\cite{gonzalez2} the evolution of wormholes and ringholes embedded in a background accelerating Universe driven by dark energy, was analyzed. It was shown that the wormhole's size increases by a factor proportional to the scale factor of the Universe, and still increases significantly if the cosmic expansion is driven by phantom energy. The accretion of dark and phantom energy onto Morris-Thorne wormholes~\\cite{diaz-phantom3,diaz-phantom4}, was further explored, and it was shown that this accretion gradually increases the wormhole throat which eventually overtakes the accelerated expansion of the universe, consequently engulfing the entire Universe, and becomes infinite at a time in the future before the Big Rip. This process was dubbed the ``Big Trip'' \\cite{diaz-phantom3,diaz-phantom4}. However, in the context of the generalized Chaplygin gas, it was shown that the Big Rip may be avoided altogether \\cite{diaz-phantom,Madrid}. We refer the reader to Refs. \\cite{Yurov:2006we} for more recent details on these issues. In summary, we denote these exotic geometries consisting of dark energy stars (in the phantom regime) and phantom wormholes as \\textit{phantom stars}. The final product of this topological change, namely, phantom wormholes, have far-reaching physical and cosmological implications, as in addition to being used for interstellar shortcuts, an absurdly advanced civilization may manipulate these geometries to induce closed timelike curves, consequently violating causality.  Relative to the topology change issue, a few words are in order. We emphasize that it is still uncertain how to obtain this change in topology, and if possible, it is riddled with technical and physical difficulties, such as the appearance of closed timelike curves. Nevertheless, it is likely that enormous negative pressures at the center is accompanied by a large quantum fluctuation of the metric. The geometry of spacetime undergoing quantum fluctuations does not seem to be a source of disagreement in the literature, but the question of whether or not the topology of spacetime undergoes quantum fluctuations is more subtle. In the latter, Wheeler conjectured that spacetime could be subjected to a topology fluctuation at the Planck scale, where spacetime undergoes a deep and rapid transformation in its structure, resulting in a spacetime quantum foam, which can be taken as a model for the quantum gravitational vacuum. Nevertheless, how to realize such a foam-like space and also whether this represents the real quantum gravitational vacuum are still unknown. From the quantum point of view we can separate the problem of topology change generated by a canonical quantization approach and a functional integral quantization approach. As mentioned above, the Hawking topology change theorem is thus enough to show that the topology of space cannot change in canonically quantized gravity. In the Feynman functional integral quantization of gravitation things are different, where an approach to spacetime foam is possible where fluctuations of topology become an important phenomenon at least at the Planck scale."
    },
    "0808/0808.2711_arXiv.txt": {
        "abstract": "The Ap star HD\\,3980 appears to be a \\changea{promising roAp candidate}  based on its fundamental parameters, leading us to search for rapid pulsations with the VLT UV-Visual Echelle Spectrograph (UVES). A precise {\\it Hipparcos} parallax and estimated temperature of 8100\\,K place HD\\,3980 in the middle of the theoretical instability strip for rapidly oscillating Ap stars, about halfway through its main sequence evolution stage. The star has a strong, variable magnetic field, as is typical of the cool magnetic Ap stars.  Dipole model parameters were determined from \\changea{VLT observations using FORS1}. From Doppler shift measurements for individual spectral lines of rare earth elements and the H$\\alpha$ line core, we find no pulsations above $20 - 30$\\,m\\,s$^{-1}$. This result is corroborated by inspection of lines of several other chemical elements, as well as with cross- correlation for long spectral regions with the average spectrum as a template. Abundances of chemical elements were determined and show larger than solar abundances of rare earth elements. Further, ionisation disequilibria for the first two ionised states of Nd and Pr are detected. We also find that the star has a strong overabundance of manganese, which is typical for much hotter HgMn and other Bp stars. Line profile variability with the rotation period was detected for the majority of chemical species. ",
        "introduction": "After the discovery of low amplitude pulsation for the bright and well-studied cool, magnetic chemically peculiar A (Ap) star $\\beta$\\,CrB \\citep{hatzesetal04,kurtzetal07}, the question arose whether all cool Ap stars (with effective temperatures below $T_{\\rm eff}\\sim 8200$\\,K) are rapid oscillators. This still unanswered question could be a pivotal point for theoretical modelling and understanding of the driving mechanism in rapidly oscillating (roAp) stars. Several searches for rapid oscillations in Ap stars have been made. \\citet{elkinetal08a} found that 24 known roAp stars that exhibit pulsations photometrically also show rapid radial velocity variations for the corresponding pulsation periods. However, several Ap stars with photometric indices typical for known roAp stars were tested photometrically for pulsations by \\citet{martinezetal94} and found to be stable to high precision; these stars are called non-oscillating Ap stars, or noAp stars.  About a decade later, \\citet{elkinetal05a} used fast UVES spectroscopy to discover low amplitude radial velocity pulsation for one of these stars, the former noAp star HD\\,116114. \\citet{lorenzetal05} reported a new photometric null result for the star, but noted a low amplitude pulsation peak in the amplitude spectrum of their data slightly above the noise level for the same frequency. Spectroscopy therefore is superior to photometry as a tool for detection of rapid low amplitude pulsation, and we expect that more Ap stars previously identified as non- oscillating (noAp) may exhibit rapid radial velocity variations. One such promising example is HD\\,965, which has physical parameters corresponding to those of known roAp stars. The fast photometry by \\citet{kurtzetal03} and high time resolution spectroscopy by \\citet{elkinetal05b}, however, found no rapid variability. Both of these studies mentioned the known roAp problem of amplitude modulation with the rotation phase as a possible explanation for the null result and proposed to re-observe the star at a different rotation phase when one of the magnetic poles is at a more favourable aspect.  As seen in, e.g., the astrometric HR-diagram by \\citeauthor{hubrigetal05}\\, (\\citeyear{hubrigetal05}, their figure 2) for roAp and noAp stars, the apparent noAp stars occupy essentially the same regions as the roAp stars. However, the noAp stars appear to be systematically more evolved than the roAp stars \\citep{northetal97,handleretal99,hubrigetal00}. Nevertheless, the theoretical roAp instability strip \\citep{cunha02} also predicts rapid pulsations for the more evolved Ap stars (near the terminal age the main sequence). HD\\,116114 \\citep{elkinetal05b} was indeed detected in this region of the HR-diagram, oscillating with the predicted frequency 0.79\\,mHz (the lowest frequency known for the roAp stars). Still the only case of a luminous roAp star, HD\\,116114 shows extremely low radial velocity pulsation amplitude and only for small number of chemical elements such as europium and lanthanum.  Using the same instrument (UVES) and procedure as \\citet{elkinetal05b}, \\citet{freyhammeretal08} searched for rapid pulsations among a group of nine evolved Ap stars inside the roAp instability region in the HR-diagram. Surprisingly they did not detect any radial velocity pulsation for these stars, but showed that only $3-5$ of the stars may have magnetic field strengths considerably in \\changea{excess of 2\\,kG}. More evolved stars are theoretically expected to require relatively strong magnetic fields to suppress local surface convection and facilitate observable amplitudes of rapid pulsations (see \\citealt{cunha02}).  The bright Ap star HD\\,3980 ($\\xi$\\,Phe, $V=5.719$ mag) was proposed by one of us (Hubrig) to be a good roAp candidate because of its many properties in common with those of known roAp stars, such as a magnetic field and peculiar abundances of rare earth elements, as we show in Fig.\\,\\ref{Fig1}. A photometric search for rapid pulsations had already been performed for HD\\,3980 by \\citet{martinezetal94} who failed to detect any pulsational variability. Based on our success with high time resolution spectroscopy (e.g. \\citealt{kurtzetal07,mathysetal07}), we then decided to use fast spectroscopy to test this promising roAp candidate for rapid pulsations. The following sections discuss the collection of data, their analysis and the null result for rapid oscillations.  \\begin{figure} \\begin{center} \\hfil \\epsfxsize 8.5cm\\epsfbox{sp_pl1.eps} \\caption{\\label{Fig1}A small spectral region of HD\\,3980, compared with those of the two roAp stars, HD\\,128898 ($\\alpha$\\,Cir) and HD\\,176232 (10\\,Aql). The spectra of the latter two stars are shifted up and down, respectively, in intensity by 0.2. Selected atomic lines are indicated. HD\\,3980 shows strong lines for the rare earth elements Nd and Pr, a common characteristic of known roAp stars. } \\end{center} \\end{figure} ",
        "conclusions": "The cool Ap star HD\\,3980 shows no rapid oscillations in radial velocity above a few tens of m\\,s$^{-1}$ at two rotational phases separated by half a rotation period. We have shown that at least one of these coincides with a rotation phase near a magnetic extremum when pulsation amplitude for an oblique pulsator is expected to be highest. The methods applied, and the lengths of the data sets acquired, would be sufficient for detecting pulsations in all known roAp stars. We therefore now reconsider whether HD\\,3980 {\\it is} a good roAp candidate, when comparing its characteristics to those of known roAp stars. If indeed it is a good candidate, it either does not pulsate (thus is a noAp star), or we failed to detect the pulsations, or, alternatively, the roAp class characteristics are still too poorly established to physically discern noAp stars from roAp stars.  The blue edge of the roAp instability strip where it crosses the main sequence is not firmly established. \\citet{hubrigetal00} showed that it is around 8500\\,K, while the theoretical roAp instability strip by \\citet{cunha02} extends up to around 9500\\,K. \\citet{ryabchikovaetal04} suggest a transition region (noAp/roAp) around 8100\\,K where cooler Ap stars may be rapid oscillators. HD\\,3980 has a precise parallax  ($\\pi=14.91\\pm0.35$, \\citealt{vanleeuwen07}) that for \\changea{$A_V=0.054$ (NASA/IPAC IRSA maps) gives a luminosity of $\\log L$/L$_\\odot=1.27\\pm0.03$ (including all uncertainties). We found an effective temperature of $8100\\pm200$\\,K (estimated error) which places the star well inside the theoretical roAp instability strip. The luminosity was derived using a bolometric correction, $BC=0.024\\pm0.009$, from the relations by \\citet{flower96} for the range of $T_{\\rm eff}$ within its error. From this luminosity and $T_{\\rm  eff}$ we derive a radius of ${\\rm R} = 2.19^{+0.19}_{-0.16}$\\,R$_{\\odot}$. Eq.~5 then gives the magnetic obliquity to be $\\beta = 88\\fdg7^{+0.8}_{-4.0}$. Bolometric corrections are notoriously difficult to determine for A stars with peculiar abundances. If one similarly uses the bolometric calibration by \\citet{landstreetetal07}, accounting for a 0.1\\,mag intrinsic error in the calibration, the corresponding luminosity and radius become: $\\log L$/L$_\\odot=1.25\\pm0.07$ and ${\\rm R} = 2.14^{+0.30}_{-0.26}$\\,R$_{\\odot}$, resulting in $\\beta = 88\\fdg6^{+0.9}_{-28.4}$. }  \\begin{figure} \\begin{center} \\epsfxsize 8.5cm\\epsfbox{evol_3980.eps} \\caption{\\label{Fig6} A theoretical HR-Diagram for a sample of roAp stars. Luminosities were calculated using Hipparcos trigonometric parallaxes \\citep{perrymanetal97} \\changea{except for HD\\,3980, for which the revised parallax of \\citet{vanleeuwen07} was used}. The \\citet{moonetal85} calibration of Str\\\"omgren photometric indices was employed for estimating the effective temperature. The position of HD\\,3980 is shown by the filled diamond which is the size of the error bar in \\changea{$\\log L$/L$_\\odot$. Solid lines are evolutionary} tracks for stars with 1.5, 1.7 and 2.0 M$_{\\odot}$ taken from \\citet{schalleretal92}.} \\end{center} \\end{figure}  The roAp stars have strong magnetic fields that range from several hundred Gauss to 24.5\\,kG in HD\\,154708 \\citep{hubrigetal05}. HD\\,3980 has a longitudinal field, which is within this range. The evolutionary tracks in Fig.~\\ref{Fig6}, as well as those in figure 1 of \\citet{cunha02}, indicate that HD\\,3980 in the midst of its life on the main sequence, halfway evolved from the ZAMS to the TAMS. \\citet{cunha02} argued that more evolved roAp stars require stronger magnetic fields to suppress the upper envelope convection and enable the rapid oscillations to reach observable amplitudes. She further speculated that the magnetic field intensities needed for suppressing the convection in roAp stars more often are found in roAp stars with magnetically resolved lines. That typically requires a magnetic field modulus of $\\sim$$3$\\,kG \\changea{and a slowly rotating Ap star} ($v \\sin i=1- 3$\\,km\\,s$^{-1}$). Based on the longitudinal field measurements of HD\\,3980, it could have a magnetic field modulus more then 3\\,kG. The rotational broadening dominates any magnetic splitting of lines at this field intensity and direct measurement is not possible.  \\citet{freyhammeretal08} searched for pulsations in 9 evolved cool Ap stars located inside the theoretical instability strip, the majority having estimated temperatures below $8100$\\,K. With similar precision and upper limits on radial velocity amplitudes as in this study, these authors found 9 null results. Out of 7 stars, only 3 had magnetic fields significantly stronger than 2\\,kG, which possibly explains most of their null results for such evolved stars. The magnetic field of HD\\,3980 is strong enough to suppress local convection and enable pulsational driving of rapid pulsations to reach observable amplitudes. The star does not appear to be near the terminal end of its main sequence lifetime, in which case this explanation may be less likely. \\citet{hubrigetal00} pointed out the apparent deficiency of close binaries among roAp stars (although a `handful' of exceptions are known). In this context, it is interesting to note that HD\\,3980 is component of a common proper-motion visual binary (see, e.g.,  \\citealt{perrymanetal97}).  Our abundance analysis found that HD\\,3980 has strong enhancement of rare earth elements, ionisation disequilibria for Nd and Pr and inhomogeneous surface distributions of these elements. In these respects, HD\\,3980 strongly resembles known roAp stars. Only a relatively high temperature, compared to the typical roAp stars, may explain why no pulsations are found. We cannot exclude extremely low amplitude pulsations of about $5-10$\\,m\\,s$^{-1}$. Low amplitude oscillations in roAp stars are known from, e.g., $\\beta$\\,CrB \\citep{kurtzetal07}, HD\\,154708 \\citep{kurtzetal06a} and HD\\,116114 \\citep{elkinetal05a}. For such low pulsation amplitudes, m\\,s$^{-1}$ or even sub-m\\,s$^{-1}$ high precision measurements are required, such as \\citet{kochukhovetal07} obtained with HARPS for HD\\,75445. This in turn requires the stars to be slow rotators, although cross-correlation for long spectral regions partly compensates for that. In the case of HD\\,3980, even the high precision of cross-correlation analysis did not detect low amplitude pulsations.  It is important to find roAp stars with rotation periods of a few days, such as for HD\\,3980, as they can be subjected to 3D pulsation and abundance studies, such as HR\\,3831 \\citep{kochukhovetal07}. But if the pulsation amplitudes are very low, such a study will be limited to very few individual lines or elements and will not be very informative.  We have found that HD\\,3980 is a bright, nearby, Li-spotted Ap star with a short rotation period which makes it an excellent observing target for \\changea{Doppler imaging.} We demonstrated that the surface distribution is spotted for several elements, such as Nd, Eu and Li. We have shown that the star has similar characteristics to the known roAp stars, but that it is stable to pulsations with an upper limit to the radial velocity amplitudes of individual spectral lines of $30$\\,m\\,s$^{-1}$, less than the amplitudes of the known roAp stars."
    },
    "0808/0808.0714_arXiv.txt": {
        "abstract": "We present CCD imaging observations of early-type galaxies with dark lanes obtained with the Southern African Large Telescope (SALT) during its performance-verification phase. The observations were performed in six spectral bands that span the spectral range from the near-ultraviolet atmospheric cutoff to the near-infrared. We derive the extinction law by the extragalactic dust in the dark lanes in the spectral range $ 1.11 \\mu m^{-1} <\\lambda^{-1}<2.94 \\mu m^{-1}$ by fitting model galaxies to the unextinguished parts of the image, and subtracting from these the actual images. This procedure allows the derivation, with reasonably high signal-to-noise, of the extinction in each spectral band we used for each resolution element of the image. We also introduce an alternative method to derive the extinction values by comparing various colour-indices maps under the assumption of negligible intrinsic colour gradients in these galaxies. We than compare the results obtained using these two methods.  We compare the total-to-selective extinction derived for these galaxies with previously obtained results and with similar extinction values of Milky Way dust to derive conclusions about the properties of extragalactic dust in different objects and conditions.  We find that the extinction curves run parallel to the Galactic extinction curve, which implies that the properties of dust in the extragalactic enviroment are similar to those of the Milky Way, despite our original expectations. The ratio of the total V band extinction to the selective extinction between the V and B bands is derived for each galaxy with an average of $2.82\\pm0.38$, compared to a canonical value of 3.1 for the Milky Way. The similar values imply that galaxies with well-defined dark lanes have characteristic dust grain sizes similar to those of Galactic dust. We use total optical extinction values to estimate the dust mass for each galaxy, compare these with dust masses derived from IRAS measurements, and find them in the range $10^4$ to $10^7M_{\\odot}$. ",
        "introduction": "\\label{txt:Obs_and_Red} The Southern African Large Telescope (SALT) was described by Buckley, Swart \\& Meiring (2006), and its CCD camera (SALTICAM) was described by O'Donoghue et al.\\ (2006). We used the SALT and SALTICAM to observe dark-lane early-type galaxies.  It may seem superfluous to observe bright galaxies with an (effectively) nine-meter telescope since smaller telescopes served admirably this purpose in the past. These observations, however, were obtained during the Performance Verification (PV) phase of the SALT telescope when only the SALTICAM was available, and even this with the instrument before its upgrade to auto-guiding. The telescope was still affected by the limited image quality, which restricted the size of the workable field to some three arcmin instead of the nominal eight arcmin region. The specific advantage in performing this project with SALT and SALTICAM is the high efficiency of this combination in the near-ultraviolet, and the availability of the U1 and U2 filters (see below) among the SALTICAM filter complement. \\begin{figure} \\includegraphics[width=84mm]{transmission.eps} \\caption{SALTICAM short wavelength filters: 340/45 nm (U1- dashed line) and 380/40 nm (U2- solid line) transmission.} \\label{fig:transmission} \\end{figure}  The SALTICAM is a CCD mosaic of two 2048x4102 pixels CCDs and therefore, for consistency, program objects were imaged with the same CCD throughout the observations. The full SALTICAM image is about 9'.6 by 9'.6 with pixels of $\\sim$0\".28 (after binning on-chip by a factor of two), but the nominal science field is eight arcmin.  Observations of nine dust lane early-type galaxies were performed in service mode from February to May 2007 with SALT, whereas only for eight objects we derived the extinction values (see \\S~\\ref{txt:results}). The nine galaxies are listed in Table \\ref{t:Obs} with their coordinates, morphological classification and optical size as taken from the RC3 catalog (de Vaucouleurs et al.\\ 1992) or from LEDA (http://leda.univ-lyon1.fr). CCD imaging observations were performed with SALTICAM using the standard B, V, R \\& I filters, as well as two short-wavelength interference filters, U1 and U2, with transmission bands peaking at 340 nm (FWHM 35 nm) and 380 nm (FWHM 40 nm) respectively, as shown in Fig.\\ \\ref{fig:transmission}. The plotted transmission profile is logarithmic, to emphasize the reasonably good blocking of out-of-band wavelengths. U1 has its strongest out-of-band transmission band above 1 $\\mu$m where the SALTICAM CCD already has a fairly low quantum efficiency. U2 is slightly more affected by a red ``leak'' given its prominent transmissions at $\\sim$500 nm and at $\\sim$700 nm. However, the first out-of-band peak at 480 nm is at less than $1\\%$ level, and the second one, a blue leak near 310nm, is at $\\sim 0.1\\%$ level.  The objects were selected from catalogs for objects with dust lanes by Hawarden et al.\\ (1981), Ebneter \\& Balick (1985) and V\\'{e}ron-Cetty \\& V\\'{e}ron (1988). The observed objects were chosen according to visibility conditions and availability of observing nights requiring that the galaxies would fit in the good image quality field of SALTICAM ($\\sim$ three arcmin), would be in the sky region accessible to SALT at the date of observation, and would not present reduction complications such as dust lanes with bright edges that might signify internally-produced light within the dust lanes. Galaxies from Patil et al.\\ (2007) and from Goudfrooij et al.\\ (1994) need the full observation set even though they were already observed by these authors in some of the bands. This is because the images to be used for the derivation of the extinction should all have the same seeing size. This is ensured by obtaining the six images quasi-simultaneously, with the same camera and telescope. The observations consist of short-duration integrations while the filter wheel cycles between the six filters, thus ensuring quasi-similiar seeing size for all bands observed in a single filter cycle. Exposure times are determined to be sufficiently short to prevent image trailing when operating in the SALTICAM unguided mode (but with SALT tracking) or overexposure of the brightest part of each galaxy. The six-image cycles were repeated while the pointing was slightly shifted to allow for a dithering effect. The cycles were then accumulated to reach better exposure depth. Exposure times in seconds for B, V, R, I, U2 \\& U1 bands, as well as the typical FWHM of stars near target in the different bands, are listed in Table \\ref{t:exposure}. The numbers in parenthesis represent the number of exposures acquired in a given band. Since each image consists of a combination of a number of short-duration exposures, the FWHM value represents the value to which the different images in each band were convolved. The size of the resolution elements is a result of both seeing and image quality (IQ) effects.  The SALT pipeline includes bias subtraction, overscan subtraction and ``cross-talk'' correction, while other standard preprocessing steps were done using standard tasks within IRAF\\footnote{IRAF is distributed by the National Optical Astronomy Observatories (NOAO), which is operated by the Association of Universities, Inc. (AURA) under co-operative agreement with the National Science Foundation}. Such steps include geometric alignment of multiple frames taken in each filter by measuring centeroids of several common stars in the galaxy frames. The frames were then combined with median scaling to improve the S/N ratio. This alignment procedure involves IRAF tasks for scaling, translation and rotation of the images, so that a small amount of smearing is introduced affecting the accuracy to be better than half a pixel. Median combination was also useful in removing cosmic rays events while the CCD hot pixels were removed with the CCDMASK task in IRAF using an appropriate mask. We emphasize that no flatfield correction was made during this  reduction process since at the time of observations the telescope did not yet have a moving baffle to simulate the effect of its continuously changing pupil on the flatfield during observations.  Since no flatfield correction was applied, the sky background cannot be estimated naivly by measuring at various locations in the frame away from the galaxy, or by fitting a tilted plane on the basis of small areas far from the galaxy images. In order to determine the sky background we took advantage of its statistical nature; we first created a count-frequency histogram for the galaxy surroundings (that is, in a 3' radius from the galaxy center) and then fitted a Gaussian distribution curve to the peak of the histogram, where the most frequent count value was assumed to represent the sky background. This value was subtracted from the galaxy frame. Due to the nature of the unflatfielded frames we avoided combining frames from different observing nights. Furthermore, only observations with similar seeing values within the same observing night were selected to produce the combined image after  convolution with a Gaussian to match the image seeing. ",
        "conclusions": "\\label{txt:summ}  We presented SALT observations using the SALTICAM of early-type galaxies with dust lanes obtained in order to derive the extragalactic extinction law down to the ultraviolet atmospheric cutoff. Our work also extends that done by previous authors since our derived extinction laws reach to $\\lambda^{-1}\\approx2.94$. This extension demonstrates the necessity of performing such studies at as short a wavelength as possible.  Extinction curves derived for the sample galaxies run parallel to the canonical MW curve, implying similiar properties between MW canonical grains and dust in the extragalactic enviroment. We derived the ratio between the total V band extinction and the selective B and V extinction $R_V$ for each galaxy,  and obtained an avergage of $2.82\\pm0.38$. These results are in agreement with Patil et al.\\ (2007), indicating that galaxies with well-defind dust lanes tend to have slightly smaller $R_V$ values with respect to the Galactic value, but within the error bars. This suggests that the characteristic grain sizes responsible for the optical extinction are similar in size to the canonical MW grains. We verified our results by using an alternative method to derive the extinction values, which is based on comparing colour-index maps.  The dust content derived by the optical extinction method for the sample galaxies is in the range $\\sim 10^4$ to $10^7M_{\\odot}$; this also is in agreement with Patil et al.\\ (2007) and with previous studies.  \\subsection*"
    },
    "0808/0808.0843_arXiv.txt": {
        "abstract": "We introduce a simple and efficient algorithm for diffusion in smoothed particle hydrodynamics (SPH) simulations and apply it to the problem of chemical mixing. Based on the concept of turbulent diffusion, we link the diffusivity of a pollutant to the local physical conditions and can thus resolve mixing in space and time. We apply our prescription to the evolution of an idealized supernova remnant and find that we can model the distribution of heavy elements without having to explicitly resolve hydrodynamic instabilities in the post-shock gas. Instead, the dispersal of the pollutant is implicitly modelled through its dependence on the local velocity dispersion. Our method can thus be used in any SPH simulation that investigates chemical mixing but lacks the necessary resolution on small scales. Potential applications include the enrichment of the interstellar medium in present-day galaxies, as well as the intergalactic medium at high redshifts. ",
        "introduction": "Understanding chemical enrichment and the dispersal of heavy elements in the wake of energetic supernovae (SNe) has become essential to a number of fields in astrophysics. The details of how enriched material mixes with ambient gas are not only relevant for the cooling and fragmentation properties of the interstellar medium (ISM), but also manifest themselves in the composition and dynamics of the resulting stars. As massive stars return metals to the ISM, mixing plays a key role in the overall matter cycle in galaxies. Tracing this loop back in time, one encounters the initial enrichment of the pristine, pure H/He gas at the end of the cosmic dark ages, when the very first stars, termed Population III (or Pop~III), ended their lives in violent SN explosions and expelled a significant fraction of their mass in metals \\citep{heger03,iwamoto05,greif07,wa08b}. The resulting mixing in the early Universe has two aspects. The first concerns the enrichment inside the first galaxies, when the metals ejected by Pop~III stars recollapsed into more massive haloes that cooled by atomic hydrogen lines and vigorously mixed with pristine material in the presence of a highly turbulent medium \\citep{wa07b,greif08a}. The second relates to the enrichment of the intergalactic medium (IGM) at redshifts $z>5$ \\citep{mfr01}. The resulting metallicity distribution is a crucial ingredient in modeling the reionization history of the Universe \\citep{ybh04,fl05} and in determining when cosmic star formation shifted from a predominantly high-mass, Pop~III mode, to a more normal Pop~II mode \\citep{schneider02,mbh03,venkatesan06,tfs07}.  Unfortunately, an accurate treatment of mixing, while simultaneously simulating the larger-scale environment, is presently not feasible, as the turbulent motions responsible for mixing typically cascade down to very small scales. A frequently encountered approach to this problem is to assume that the products of stellar nucleosynthesis are distributed within a fixed volume \\citep[e.g.][]{nop04,scannapieco05b,ksw07,tfs07}. Significantly better results can be achieved in grid-based simulations by relating the mass flux between cells to the mixing efficiency, even though it remains unclear how much of this mixing is numerical instead of physical \\citep[e.g.][]{wa08b}. Such a direct approach is difficult to implement in SPH simulations, due to their Lagrangian nature, and instead chemical mixing has been modelled as a diffusion process \\citep[e.g.][]{martinez-serrano08}. This is somewhat more accurate since the rms displacement of a fluid element in a homogeneously and isotropically driven turbulent velocity field can be described by the diffusion equation, with the diffusion coefficient set by the velocity dispersion and the typical shock travel distance \\citep{kl03}.  Prescriptions with a constant diffusion coefficient have been applied to models of the chemical evolution of the Milky Way \\citep{karlsson05,kg05}, galaxies in a cosmological context \\citep{martinez-serrano08}, and the environment of the first galaxies \\citep{kjb08}. In the present paper, we introduce an improved method that resolves chemical mixing in space and time based on the velocity dispersion within the SPH smoothing kernel. This yields results in agreement with more detailed investigations of mixing in the early phases of SN remnants, without having to explicitly resolve the hydrodynamic instabilities in the post-shock gas. In future work, we plan to use our method to investigate the long-term evolution of energetic SNe in a cosmological environment, particularly at high redshift when the Universe was enriched with the first metals. However, since our algorithm is quite generic, we hope that it will prove a valuable tool for any SPH simulation that attempts to follow the mixing of pollutants.  The structure of our work is as follows. In Section~2, we describe our model for diffusion and its numerical implementation in the cosmological simulation code {\\sc GADGET}-2 \\citep{springel05}, followed by a series of idealized test simulations (Section~3). We then discuss the relation between the diffusion coefficient and the velocity dispersion and apply our prescription to the evolution of an idealized SN remnant (Section~4). Finally, in Section~5 we summarize our results and assess their implications. For consistency, all quoted distances are physical, unless noted otherwise. ",
        "conclusions": "We have introduced a simple and efficient algorithm for diffusion in SPH simulations and investigated its accuracy with a number of idealized test cases. Although the algorithm is quite generic and can be applied to most problems involving diffusion, we have here focused on a model for the dispersal of enriched material in supernova explosions. Adopting the mixing length approach discussed by \\citet{kl03}, we link the diffusivity of the pollutant to the local physical conditions and can thus describe the space- and time-dependent mixing process. We have applied our prescription to an idealized SN explosion and found that we can properly describe the mixing process without having to explicitly resolve the hydrodynamic instabilities in the post-shock gas. Instead, this process is implicitly governed by the dependence on the local velocity dispersion. Our method can thus be used in any SPH simulation that attempts to follow the mixing of a pollutant but lacks the necessary resolution on small scales. This will be relevant, in particular, for simulations that aim at simultaneously representing the large-scale environment of the SN explosion and the fine-grained mixing. A crucial assumption of our model is that resolved motions cascade down to subresolution scales which then homogeneously mix the gas. For this reason our method yields an upper limit to the mixing efficiency and the resulting degree of homogeneity, and cannot capture the details of a real turbulent cascade. However, since turbulent motions on scales larger than the smoothing length are explicitly resolved, this approximation should generally be valid.  Our new algorithm ties in well with the current generation of radiation-hydrodynamical simulations of star formation, both at high redshifts and in the local Universe. The ultimate goal is to describe the interrelated cycle of gas fragmentation and stellar feedback. To achieve adequate realism, radiative effects have to be considered together with the mechanical and chemical feedback due to SN explosions. The methodology developed here allows us to include chemical feedback, suitable even for extremely large simulations. We will report on specific applications in future studies."
    },
    "0808/0808.0591_arXiv.txt": {
        "abstract": "We compute the luminosity function (LF) and the formation rate of long gamma ray bursts (GRBs) in three different scenarios:  i) GRBs follow the cosmic star formation and their LF is constant in time; ii) GRBs follow the cosmic star formation but the LF varies with redshift; iii) GRBs form preferentially in low--metallicity environments. We then test model predictions against the {\\it Swift} 3-year data, showing that scenario i) is robustly ruled out. Moreover, we show that the number of bright GRBs detected by {\\it Swift} suggests that GRBs should have experienced some sort of luminosity evolution with redshift, being more luminous in the past. Finally we propose to use the observations of the afterglow spectrum of GRBs at $z\\ge 5.5$  to constrain the reionization history and we applied our method to the case of GRB~050904. ",
        "introduction": "Long Gamma Ray Bursts (GRBs) are powerful flashes of high--energy photons occurring at an average rate of a few per day throughout the Universe up to very high redshift (the current record is $z=6.29$). The energy source of a long GRB is believed to be associated to the collapse of the core of a massive star (see M\\'esz\\'aros 2006 for a review). One of the main goals of the {\\it Swift}  satellite (Gehrels et al. 2004) is to tackle the key issue of the GRB luminosity function (LF). Unfortunately, although the number of GRBs with good redshift determination has been largely increased by {\\it Swift}, the sample is still too poor (and bias dominated) to allow a direct measurement of the LF. We use here the {\\it Swift} 3-year data to constrain the GRB LF and its evolution (Salvaterra \\& Chincarini 2007, Salvaterra et al. 2008b). Moreover, we show a possible use of GRBs detected at $z\\ge 5.5$ to study the history of reionization (Gallerani et al. 2008). ",
        "conclusions": "We have tested different formation and evolution scenarios for long GRB against the 3-year {\\it Swift} dataset. We found that {\\it Swift} data strongly rule out models in which GRBs follow the cosmic star formation and their LF is constant in time. In particular, the number of bright GRBs suggests that GRBs should have experienced some sort of luminosity evolution with cosmic time, being more luminous in the past. Finally we have shown that GRBs at $z\\ge 5.5$ can be use to constrain the reionization history and we applied our method to the case of GRB~050904 at $z=6.29$."
    },
    "0808/0808.0601.txt": {
        "abstract": "% context heading (optional) {} % aims heading (mandatory) {We present a general method to solve radiative transfer problems including scattering in the continuum as well as in lines in 3D configurations with periodic boundary conditions. } % methods heading (mandatory) {The scattering problem for line transfer is solved via means of an operator splitting (OS) technique. The formal solution is based on a full characteristics method. The approximate $\\Lambda$ operator is constructed considering nearest neighbors exactly. The code is parallelized over both wavelength and solid angle using the MPI library. } % results heading (mandatory) {We present the results of several test cases with different values of the thermalization parameter and two choices for the temperature structure. The results are directly compared to 1D plane parallel tests. The 3D results agree very well with the well-tested 1D calculations. } %conclusions heading (optional)%,leave it empty if necessary {Advances in modern computers will make realistic 3D radiative transfer calculations possible in the near future. Our current code scales to very large numbers of processors, but requires larger memory per processor at high spatial resolution.} ",
        "introduction": "Interest in 3-D radiative transfer in stellar atmospheres has grown with the calculations of Asplund and collaborators \\citep{ANTS99,ANTS00,AspIII00,AGS02,GAS07}. This work has indicated that the solar oxygen abundance needs to be revised downward. However, the revised abundances are difficult to reconcile with helioseismological results \\citep[see][and references therein]{BAPR08}. The work of Asplund et al.\\ is based on comparisons of synthetic spectra produced by formal solutions of hydrodynamical models of solar convection. We present a framework for solving the full scattering problem that is applicable to hydrodynamical calculations of stellar atmospheres. \\citet[][hereafter: Paper I]{hb06} and \\citet[][hereafter: Paper II]{bh07}  described a framework for the solution of the radiative transfer equation for scattering continua and lines in 3D (when we say 3D we mean three spatial dimensions, plus three momentum dimensions) for the time independent, static case. In the 3rd paper of this series we apply these methods to problems with period boundary conditions which typically arise in radiation-hydrodynamical simulations of convective atmospheres. In such calculations the radiation transport has to be simplified compared to the full problem in order to keep the calculations tractable. However, a full solution of the scattering line problem is needed for comparison and post-processing of the structures.  {We describe our method, its rate of convergence, and present comparisons to our well-tested 1-D calculations.} ",
        "conclusions": "Using rather difficult plane parallel test problems, we have shown that our 3D full-characteristics method gives very good results when compared to our well-tested 1D code. The periodic boundary conditions method discussed here is particularly well-suited to 3-D hydrodynamical simulations of convection in stellar atmospheres and in future work we will compare our results to observations as well as to previous calculations.  The results for the computed 3D structure show that $\\Lstar$ leads also to good convergence for a true 3D structure, with convergence rates that are comparable to the simple test cases (see also Papers I and II)."
    },
    "0808/0808.3137_arXiv.txt": {
        "abstract": "We study flavor oscillations in the early universe, assuming primordial neutrino-antineutrino asymmetries. Including collisions and pair processes in the kinetic equations, we not only estimate the degree of flavor equilibration, but for the first time also kinetic equilibration among neutrinos and with the ambient plasma. Typically the restrictive BBN bound on the $\\nu_e\\bar\\nu_e$ asymmetry indeed applies to all flavors as claimed in the previous literature, but fine-tuned initial asymmetries always allow for a large surviving neutrino excess radiation that may show up in precision cosmological~data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.2313_arXiv.txt": {
        "abstract": "We investigate the effect of dust on the scaling properties of galaxy clusters based on hydrodynamic $N$-body simulations of structure formation. We have simulated five dust models plus a radiative cooling and adiabatic models using the same initial conditions for all runs. The numerical implementation of dust was based on the analytical computations of \\citet{montier04}.  We set up dust simulations to cover different combinations of dust parameters that put in evidence the effects of size and abundance of dust grains. Comparing our radiative {\\it plus} dust cooling runs to a purely radiative cooling simulation we find that dust has an impact on cluster scaling relations. It mainly affects the normalisation of the scalings (and their evolution), whereas it introduces no significant differences on their slopes.  The strength of the effect critically depends on the dust abundance and grain size parameters as well as on the cluster scaling.  Indeed, cooling due to dust is effective at the cluster regime and has a stronger effect on the ``baryon driven'' statistical properties of clusters such as $L_{\\rm X}-M$, $Y- M$, $S-M$ scaling relations. Major differences, relative to the radiative cooling model, are as high as 25\\% for the $L_{\\rm X}-M$ normalisation, and about 10\\% for the $Y- M$ and $S-M$ normalisations at redshift zero.  On the other hand, we find that dust has almost no impact on the ``dark matter driven'' $T_{\\rm mw}-M$ scaling relation.  The effects are found to be dependent in equal parts on both dust abundances and grain sizes distributions for the scalings investigated in this paper. Higher dust abundances and smaller grain sizes cause larger departures from the radiative cooling (i.e. with no dust) model. ",
        "introduction": "From the first stages of star and galaxy formation, non-gravitational processes drive together with gravitation the formation and the evolution of structures. The complex physics they involve rule the baryonic component within clusters of galaxies, and in a more general context within the intergalactic medium (IGM hereafter -- see the review by \\citet{voit05} and references therein). The study of these processes is the key to our understanding of the formation and the evolution of large-scale structure of the Universe. Indeed, understanding how their heating and cooling abilities affect the thermodynamics of the IGM at large scales and high redshifts, and thus that of the intra-cluster medium (hereafter ICM) once the gas get accreted onto massive halos is a major question still to be answered. The continuous accretion and the merger events through which a halo assembled lead to a constant interaction of the IGM gas with the evolving galactic component. Within denser environments, like clusters, feedback provided by AGN balances the gas cooling (see for instance \\citet{cattaneo07,conroy07}, and \\citet{mcnamara07} for a review). Also, from high redshift, the rate of supernovae drives the strength of the galactic winds and thus the amount of material that ends ejected within the IGM and the ICM (see \\citet{loewenstein06}). These ejecta are then mixed in the environment by the action of the surrounding gravitational potential and the dynamics of cluster galaxies within.   Since long, X-ray observations have shown the abundant presence of heavy elements within the ICM (see for instance review works by \\citep{sarazin88,arnaud05}). Physical processes like ram-pressure stripping, AGN interaction with the ICM, galaxy-galaxy interaction or mergers are scrutinized within analytical models and numerical simulations in order to explain the presence of metals (see for instance works by \\citep{kapferer06, domainko06,moll07}). Moreover, it is obvious that the process of tearing of material from galaxies leads not only to the enrichment of the ICM/IGM in metal, but in gas, stars and dust as well. Recent work on numerical simulations \\citep{murante04,murante07,conroy07} have stressed the role of hierarchical buiding of structures in enriching the ICM with stars in a consistent way with the observed amount of ICM globular clusters, and ICM light. Indeed, the overall light coming from stars in between cluster galaxies represent an important fraction of the total cluster light: for instance \\citep{krick07} measured 6 to 22\\% from a sample of 10 clusters.  The effect of a diffuse dust component within the IGM/ICM, and its effect is less known. A few observational studies with the ISO and the Spitzer satellites have tried without frank success to detect the signature of such a component \\citep{stickel98, stickel02, bai06, bai07}. More successfully, \\citep{montier05} have obtained a statistical detection, via a stacking analysis, of the overall infrared emission coming from clusters of galaxies. However, they were not able to disentangle the IR signal from dusty cluster galaxies from a possible ICM dust component. On the other hand, from a theoretical point of view a few works have looked at the effect of dust on the ICM \\citep{popescu00} or in conjunction with the enrichment of the ICM in metals \\citep{aguirre01}. However, the effect of dust on a ICM/IGM-type thermalized plasma has been formalized by \\citep{montier04}. These authors have computed the cooling function of dust taking into account the energetic budget for dust. They have shown the ability of dust to be a non negligible cooling/heating vector depending on the physical properties of the environment.  Dust thus comes, within the ICM/IGM, as an added source of non-gravitational physics that can potentially influence the formation and the evolution of large scale structure in a significant way. Indeed, since redshift of $z\\simeq 2-5$ during which the star formation activity reached its maximum in the cosmic history, large amounts of dust has been produced and thus ejected out of the galaxies due to violent galactic winds into the IGM \\citep{springel03}. As this material is then accreted by the forming halos, one can wonder about the impact produced by dust on the overall properties of clusters of galaxies once assembled and thermalized. In a hierarchical Universe, the population of clusters is self-similar, thus is expected to present well defined structural and scaling properties. However, to date, it is common knowledge that the observed properties deviate form the prediction by a purely gravitational model (see \\citep{voit05, arnaud05} for review works). It is thus important to address the issue of the impact of dust on the statistical properties of structures such as clusters of galaxies, the same way it is done for AGNs, supernovae, stripping or mergers.  In order to tackle this question, we have put into place the first N-body numerical simulations of hierarchical structure formation implementing the cooling effect of dust according to the dust nature and abundance. In this paper, we present the first results of this work focusing at the scale of galaxy clusters, and more specifically on their scaling properties. The paper is organized as follows: we start by presenting the physical dust model and how it is implemented in the numerical simulation code. In Sec .~\\ref{sec:sinum}, we describe the numerical simulations and the various runs (i.e. model) that have been tested. From these simulations our analysis concerns the galaxy cluster scale, and focus on the impact of the presence of dust on the scaling relation of clusters. In Sec.~\\ref{sec:scal}, we present our results on the $M-T$, the $S-T$, the $Y-T$ and the $L_X-T$ relations. The derived results are presented in Sec.~\\ref{sec:res} and discussed in Sec.\\ref{sec:dis}. ",
        "conclusions": "In this work, we have presented the first simulations of structure formation investigating the effect of dust cooling on the properties of the intra-cluster medium. We have compared simulations with radiative $plus$ dust cooling with respect to a purely radiative cooling simulation. We have shown that: \\begin{itemize} \\item The cooling due to dust is effective at the cluster regime and has a significant effect on the ``baryon driven'' statistical properties of cluster such as $L_{\\rm X}-M$, $Y- M$,$S-M$ scaling relations. As an added non-gravitational cooling process dust changes the normalisation of these laws by a factor up to 27\\% for the $L_{\\rm X}-M$ relation, and up to 10\\% for the $Y-M$ and $S-M$ relations. On the contrary, dust has almost no effect on a ``dark matter driven'' scaling relation such as the $T_{\\rm mw}-M$ relation.  \\item The inclusion of cooling by dust does not change significantly the slopes of the cluster scaling laws investigated in this paper. They compare with the results obtained in the radiative cooling simulation model.  \\item Through the implementation of our different dust models, we have demonstrated that the dust cooling effect at the scale of clusters depends strongly on the dust abundance in the ICM, but also on a similar proportion on the size distribution of dust grains. Therefore the dust efficiency is strongly dependent on the nature of the stripped and ejected galactic material, as well as the history of these injection and destruction processes along the cluster history. Indeed the early enrichment of dust might provide an already modified thermodynamical setup for the ``to-be-accreted'' gas at lower redshifts.  \\end{itemize}  The setup of our simulations and the limitation of our dust implementation can be considered at a ``zero order'' test with which we demonstrated the active effect of dust on structure formation and especially at the cluster scale. In order to go one step further, a perspective of this work will be needed to couple the radiative cooling function of dust with a physical and dynamical implementation of the creation and destruction processes of dust in the IGM/ICM."
    },
    "0808/0808.0911_arXiv.txt": {
        "abstract": "Recent studies have shown that distant red galaxies (DRGs), which dominate the high-mass end of the galaxy population at $z \\sim 2.5$, are more strongly clustered than the population of blue star-forming galaxies at similar redshifts.  However these studies have been severely hampered by the small sizes of fields having deep near-infrared imaging.  Here we use the large UKIDSS Ultra Deep Survey to study the clustering of DRGs.  The size and depth of this survey allows for an unprecedented measurement of the angular clustering of DRGs at $2 < z_{\\rm phot} < 3$ and $K<21$.  The correlation function shows the expected power law behavior, but with an apparent upturn at $\\theta \\lesssim 10\\arcsec$.  We deproject the angular clustering to infer the spatial correlation length, finding $10.6 \\pm 1.6 h^{-1} \\textrm{Mpc}$.  We use the halo occupation distribution framework to demonstrate that the observed strong clustering of DRGs is not consistent with standard models of galaxy clustering, confirming previous suggestions that were based on smaller samples.  Inaccurate photometric redshifts could artificially enhance the observed clustering, however significant systematic redshift errors would be required to bring the measurements into agreement with the models.  Another possibility is that the underlying assumption that galaxies interact with their large-scale environment only through halo mass is not valid, and that other factors drive the evolution of the oldest, most massive galaxies at $z \\sim 2$. ",
        "introduction": "\\label{sec:introduction}  Thus far, the precise measurements of galaxy clustering at $z \\gtrsim 2$ that are necessary for meaningful physical interpretation have only been possible for the relatively blue, star-forming galaxies that dominate optical surveys \\citep{adelberger05, lee06, ouchi05}. However the most massive galaxies at these redshifts tend to be faint in the optical, and are more appropriately selected in the near-infrared \\citep[e.g.][]{vandokkum06}.  In particular, galaxies meeting the $J-K > 2.3$ criterion for distant red galaxies \\citep[DRGs;][]{franx03} have been shown to dominate the high-mass end of the galaxy mass function \\citep{vandokkum06, marchesini07a, rudnick06}.  Measurements of DRG clustering have been severely hampered by small fields and by the near-complete reliance on photometric redshifts.  Using the largest DRG sample then available, \\citet{quadri07a} confirmed the very strong clustering found by previous authors \\citep{daddi03, grazian06}.  They also showed that strong clustering implies that DRGs occupy dark matter halos with $M \\gtrsim 10^{13} \\rm{M_\\odot}$.  However, the number density of these dark matter halos is at least an order of magnitude smaller than the number density of DRGs, suggesting that the DRG clustering measurements are incompatible with models of dark matter clustering.  Here we use a $\\sim$0.65$\\rm{deg}^2$ field from the UKIRT Infrared Deep Sky Survey (UKIDSS), which is $\\sim$8 times larger than the area used by \\citet{quadri07a} but with similar NIR depth, to determine whether the strong observed clustering of DRGs found in previous studies was an artifact due to limited field sizes or whether some other explanation must be found.  We use $(\\Omega_M, \\Omega_\\Lambda, \\sigma_8, h) = (0.3, 0.7, 0.9, 0.7)$.  Small changes in these parameters do not affect our basic conclusions.  Magnitudes are given in the Vega system, except where noted. ",
        "conclusions": "\\label{sec:discussion}  We have used the UKIDSS-UDS to perform the first precise measurement of the clustering of red, $K$-selected galaxies at $2 < z_{\\rm phot} < 3$.  These DRGs show strong angular clustering that is well-described by a power law, but with an excess at small scales.  We use photometric redshifts to deproject the angular clustering, finding the spatial correlation length $r_0 = 10.6 \\pm 1.6 h^{-1} \\textrm{Mpc}$. This value is comparable to that measured for luminous red galaxies in the local universe \\citep{zehavi05}, however DRGs are significantly more numerous.  We show that standard models of halo occupation statistics are unable to simultaneously reproduce the observed clustering and number density, because DRGs outnumber their inferred host dark matter halos by too large a margin.  The most obvious explanation is that we have used the incorrect redshift distribution in deprojecting the angular correlation function.  A narrower distribution would reduce the correlation length (while a moderate shift in the overall distribution makes a relatively smaller difference).  However, a narrower redshift distribution would also decrease the effective volume probed by our sample, thereby increasing $n_g$.  We illustrate these effects in Figure \\ref{fig:r0_n}, which shows the observed $r_0$ and $n_g$ compared to the range of values for a typical HOD model.  It also shows how estimating $N(z)$ directly from the unperturbed photometric redshifts --- which, as mentioned in \\S~\\ref{sec:zphots}, represents the extreme assumption of no random photometric redshift errors --- affects the results.  It may still be the case that we are subject to \\emph{systematic} redshift errors, but we also note that a significantly narrower $N(z)$ would adversely affect the reasonably good agreement between our estimate of $n_g$ and the luminosity functions derived by \\citet{marchesini07a}.  Finally, we have verified that our basic results hold when using a different photometric redshift code \\citep[HYPERZ][]{bolzonella00} and with a different template set \\citep{bruzual03}.  \\begin{figure} \\epsscale{1.1} \\plotone{f3.eps} \\caption{Correlation length versus number density.  The black data point shows the values calculated using the preferred redshift distribution (\\emph{inset:} black histogram), while the red data point shows how the results change when using a narrower redshift distribution (\\emph{inset:} red hatched histogram).  For the purposes of this figure we have fixed the power law index of the correlation function to $\\gamma=1.5$, and assigned conservative $50\\%$ uncertainties to the observed galaxy number density.  The dashed line shows the approximate relation for HOD models of mass-selected galaxies based on the simulations of \\citet{zheng05} and \\citet{kravtsov04}.  The filled circles show the approximate relation for dark matter halos; the larger open circles are labeled with $\\log (M_h/\\rm{M_\\odot})$.  } \\label{fig:r0_n} \\end{figure}   Given the apparently high quality of our photometric redshifts, as well as the consistency with previous results for the clustering of DRGs, it is worth considering alternative explanations.  One possibility is that current HOD models are too simplistic, and that massive red galaxies occupy halos in unexpected ways.  The fundamental assumption underlying these models is that galaxy observables depend on halo mass, and not on the larger-scale environment or on halo properties such as structure or age.  But environment may play a role, for instance via its effect on mass accretion rates \\citep{scannapieco03,furlanetto06}.  Additionally, halo clustering varies with several halo properties, even at fixed mass; this phenomenon is generically referred to as ``assembly bias'' \\citep[e.g.][]{gao07}.  To the extent that these properties affect galaxy observables, they will also affect galaxy clustering measurements.  However, we note that current estimates for the strength of the assembly bias are too small to account for the observed discrepancies.  As an example, \\citet{gao07} have found that, at $z \\sim 2-3$ and in the relevant halo mass range, halos in the upper 20$\\%$ tail of the distribution of halo spins have a $\\sim$20$\\%$ larger bias than the mean value.  It might then be supposed that DRGs occupy less massive and more numerous halos with higher spin. But the increased number density of these low mass halos is approximately cancelled by the requirement that DRGs can only occupy 20$\\%$ of them, so the discrepancy in number densities is unchanged.  The existence of massive red galaxies at $z \\gtrsim 2$ was not predicted by models of galaxy formation, although progress has been made on this front \\citep[e.g.][but see \\citealt{marchesini07b}]{croton06,delucia07}.  The results shown here suggest that the conflict between models and observations extends to the relationship between galaxies and dark matter halos.  However, clustering measurements are sensitive to a number of systematic effects, so our conclusions remain tentative.  The most obvious source of error comes from our use of photometric redshifts.  Ongoing medium-band NIR observations will significantly reduce the photometric redshift uncertainties \\citep{vandokkum08}, and in the longer term multi-object NIR spectrographs will also improve the situation.  If future work confirms our results, clustering measurements such as those presented here may provide a new way to understand the detailed relationship between galaxy and halo properties."
    },
    "0808/0808.0857_arXiv.txt": {
        "abstract": "We study the 0.5--10\\,keV emission of a sample of five of the broadest double-peaked Balmer-line emitters with \\emph{Chandra}. The Balmer lines of these objects originate close (within a few hundred gravitational radii) to the central black holes of the Active Galactic Nuclei (AGNs), and their double-peaked profiles suggest an origin in the AGN accretion disk. We find that four of the five targets can be modeled by simple power-law continua with photon indices ($1.6\\leq\\Gamma\\leq1.8$) typical of similar luminosity AGNs. One object, SDSS\\,J0132$-$0952, shows evidence of ionized intrinsic absorption. The most-luminous SDSS double-peaked emitter, SDSS\\,J2125$-$0813, has either an unusual flat spectrum ($\\Gamma\\sim1$) or is also highly absorbed. It is the only double-peaked emitter for which no external illumination is necessary to account for the Balmer line emission. The strength of the Balmer-line emission in the remaining four objects suggests that the total line flux likely exceeds the viscous energy that can be extracted locally from the accretion disk and external illumination is necessary. All five double-peaked emitters have unusually strong X-ray emission relative to their UV/optical emission, which is the likely source of the external illumination necessary for the production of the observed strong broad lines. On average about 30\\% of their bolometric luminosities are emitted between 0.5--10\\,keV. The spectral energy distributions of the five double-peaked emitters show the big blue bumps characteristic of radiatively efficient accretion flows. The Balmer line profiles, as well as the optical and X-ray fluxes of the double-peaked emitters, are highly variable on timescales of months to years in the AGN rest frame. ",
        "introduction": "\\label{intro}  Double-peaked emitters are Active Galactic Nuclei (AGNs) whose  broad Balmer-line profiles have a characteristic ``double-humped'' shape, with one peak blueward of the rest-frame wavelength of the line and another one redward of it. The best explanation for the origin of the double-peaked shape involves emission from a region in the AGN accretion disk. Since the first examples of double-peaked emission lines in AGNs were published in the 1980s \\citep[e.g.,][]{Oke87,CH89a}, more than 100 double-peaked emitters have been found. Some of these objects have also been studied in the X-ray band \\citep[e.g.,][]{2Pxray1,2Pxray2,2Pxray3,2Pxray4,S06}, but most of these studies focused on individual objects or on serendipitous X-ray detections of limited quality.  \\begin{figure} \\plotone{f1.eps} \\caption{The redshift-luminosity distribution of the five-object \\emph{Chandra} sample with respect to the full S03 sample, the \\emph{ROSAT} sample studied by S06, and the X-ray observed E03 sample. Two of the E03 objects, discussed in \\S~\\ref{intro}, are denoted by letters:  A stands for Pictor\\,A, and B stands for Arp\\,102B. All monochromatic luminosities are corrected for host-galaxy contamination. \\label{lzplot}} \\end{figure}  \\begin{deluxetable}{ccccc} \\tablewidth{0pt} \\tablecaption{\\emph{Chandra} Sample} \\tablehead{\\colhead{SDSS Name} &\\colhead{$\\log R$} &\\colhead{$z$} &\\colhead{FWHM} &\\colhead{$M_i$} \\\\ \\colhead{(1)} &\\colhead{(2)} &\\colhead{(3)} &\\colhead{(4)} &\\colhead{(5)} \\\\ \\colhead{HHMMSS.ss$\\pm$DDMMSS.s} &\\colhead{} &\\colhead{} &\\colhead{km\\,s$^{-1}$}&\\colhead{mag}} \\startdata J013253.31$-$095239.5  & $<0.7$ & 0.260 & 15200 & $-22.9$\\\\ J100443.44$+$480156.5 & $<0.6$ & 0.199 & 16000 & $-22.8$\\\\ J123807.77$+$532556.0 & $2.3$ & 0.348  & 15400 & $-24.4$\\\\ J140720.70$+$023553.1 & $<1$  & 0.309  & 20100 & $-22.9$\\\\ J212501.21$-$081328.6  & $<0.2$ & 0.625 & 14000 & $-26.8$ \\\\ \\enddata \\tablecomments{(1) SDSS name given in the J2000 epoch RA and Dec form; (2) Radio loudness parameter; (3) Redshift; (4) Full-Width at Half Maximum of the H$\\alpha$ (objects \\# 1 through 4) or H$\\beta$ (object \\# 5) line; (5) Extinction- and K-corrected \\citep[assuming a power-law continuum with index $\\alpha_{\\nu}=-2.45$ for \\hbox{$\\lambda>5000$\\,\\AA};][]{Berk} SDSS $i$-band model magnitude (including the host-galaxy contribution).} \\label{tabS} \\end{deluxetable}  \\begin{deluxetable*}{ccccccccc} \\tablewidth{0pt} \\tabletypesize{\\small} \\tablecaption{X-ray Observations} \\tablehead{\\colhead{Object} &\\colhead{ObsID} &\\colhead{Instrument} &\\colhead{ObsDate} &\\colhead{$t_{\\textrm{exp}}$} &\\colhead{Counts} & \\colhead{Count Rate} &\\colhead{PSF FWHM}&\\colhead{\\emph{ROSAT}}\\\\ \\colhead{(1)} &\\colhead{(2)} &\\colhead{(3)} &\\colhead{(4)} &\\colhead{(5)} &\\colhead{(6)} &\\colhead{(7)} &\\colhead{(8)}&\\colhead{(9)}\\\\ \\colhead{HHMM$\\pm$DDMM} &\\colhead{} &\\colhead{} &\\colhead{yyyy-mm-dd} &\\colhead{sec} &\\colhead{}&\\colhead{counts\\,s$^{-1}$} &\\colhead{$''$} &\\colhead{counts\\,s$^{-1}$}} \\startdata J0132$-$0952& 6770 & ACIS-7             & 2006-06-24 & 8151 & 727 & $0.089\\pm0.003$ & $0.85\\pm0.03$ & $\\lesssim0.011$\\\\ J1004$+$4801 & 6771 & ACIS-235678 & 2006-01-31 & 8628 & 359 & $0.042\\pm0.002$ & $0.80\\pm0.03$ & $0.04\\pm0.01$\\\\ J1238$+$5325 & 6773 & ACIS-7           & 2006-05-31 & 4061 & 1133 & $0.279\\pm0.008$ & $0.75\\pm0.02$ & $0.10\\pm0.02$\\\\ J1407$+$0235 & 6772 & ACIS-235678 & 2005-12-07 & 5471 & 1348* & $0.246\\pm0.007$* & $1.1\\pm0.1$ & $\\lesssim0.014$\\\\ J2125$-$0813 & 6774 & ACIS-7             & 2006-04-02 & 4110 & 155   & $0.038\\pm0.003$ & $1.1\\pm0.1$ & $0.06\\pm0.01$\\\\ \\enddata \\tablecomments{The asterisk (*) denotes observed values, which are not corrected for pile-up. (1) Short object name; (2) \\emph{Chandra} Observation ID; (3) ACIS chips used; all observations used the 1/8 subarray except ObsID 6772, which used the 1/2 subarray; (4) Date of the observation; (5) Effective exposure time, $t_{\\textrm{exp}}$; (6) Total counts in the \\hbox{0.3--10\\,keV} spectrum; (7) Count rate for the ACIS-S3 chip; (8) The Full Width at Half Maximum (FWHM) of the ACIS image Point Spread Function (PSF); (9) The \\hbox{0.5--2\\,keV} \\emph{ROSAT} count rate.} \\label{tab1} \\end{deluxetable*}   The small sample of double-peaked emitters studied in detail in the X-ray band to date consists of predominantly low-luminosity radio-loud\\footnote{Following \\citet{rl} we define an AGN as radio loud if $\\log_{10} R=\\log[F\\si{20\\,cm}/F_i]=0.4(i-t)>1$, where $i$ is the Sloan Digital Sky Survey $i$-band magnitude and $t=-2.5\\log(F\\si{20\\,cm}/3136\\,\\textrm{Jy})$ is the 20\\,cm AB radio magnitude.} (RL) objects -- broad line radio galaxies (BLRGs, Eracleous et al.~1994, 2003, hereafter E03). Based on the prototype disk emitter (Arp~102B) and the small original BLRG sample, a model was developed to explain the existence of double-peaked emission in AGNs. \\citet{CH89b} proposed that BLRGs and other low-accretion-rate, low-luminosity\\footnote{$L\\si{bol}<1.3\\times10^{43}\\textrm{\\,erg\\,s}^{-1}=10^{-3} \\eta M_8^{-1}L\\si{Edd}$, where $M_8$ is the black hole mass in units of $10^8$\\,$M_{\\sun}$, $\\eta$ is the accretion efficiency, and $L\\si{Edd}$ is the Eddington luminosity. Efficient accretors can convert up to $\\eta\\lesssim0.3$ \\citep[for a maximally rotating black hole;][]{Thorne74} of the gravitational energy into radiation.} AGNs -- which do not accrete through a standard thin disk -- could exhibit low-ionization lines with double-peaked profiles. The line profiles are produced as the outer portions of the disk are illuminated by an X-ray-emitting elevated structure in the center -- a structure that is characteristic of radiatively inefficient accretion flows \\citep[RIAFs; e.g,][]{Rees82,adaf,adios,cdaf} and replaces the standard thin disk. The illumination provided by the central elevated structure in AGNs with RIAFs is necessary, but not sufficient, for the production of double-peaked lines of the observed strength: all BLRGs with double-peaked lines require external illumination (since the viscous power available locally in the disk is insufficient to explain the line strength), but not all BLRGs or other RIAF AGNs have double-peaked lines. The association of RL AGNs with RIAFs could explain the higher frequency of double-peaked emitters observed among BLRGs (20\\%; E03), compared to 3\\% of low-redshift Sloan Digital Sky Survey (SDSS) AGNs. Recent work by \\citet{Lewis}, however, has shown that not all BLRG double-peaked emitters with low optical luminosities are inefficient emitters: Pictor~A, for example, emits a bolometric luminosity of more than 8\\% of  the Eddington luminosity, which is too high to support a RIAF \\citep[RIAF solutions exist for $\\dot M\\si{crit}=\\alpha^2\\dot M\\si{Edd}<0.09\\dot M\\si{Edd}$, for a viscosity parameter $\\alpha<0.3$, and the radiated energy is further reduced by the accretion efficiency;][]{NY95}.  \\begin{figure*} \\plotone{f2.eps} \\caption{X-ray spectra of the five double-peaked emitters overlaid with the best-fit models. The panels are arranged in order of increasing RA, with the RA increasing to the right and down. The second panel displays the $1-3\\sigma$ confidence contours for the gas column density and ionization parameter for the SDSS\\,J0132$-$0952 fit shown in the first panel. The SDSS\\,J1407$-$0235 fit takes into account the pile-up estimated to be 12\\%. The last panel gives the SDSS\\,J2125$-$0813 Cash-statistic fit which has no $\\Delta\\chi^2$ residuals; $\\Delta\\chi$ refers to the residuals in units of $\\sigma$ with errorbars of size 1. \\label{Xspec}} \\end{figure*}  The systematic search for AGNs with double-peaked Balmer lines in the SDSS revealed that as many as 3\\% of $z<0.33$ objects are double-peaked emitters, regardless of their radio loudness (Strateva et al.~2003; hereafter, S03). Moreover, a handful of luminous higher-redshift double-peaked emitters were discovered serendipitously based on their H$\\beta$ and \\ion{Mg}{2} line profiles. The expansion of the parameter space of the double-peaked emission-line AGNs to include radio-quiet (RQ, $\\log R\\leq 1$) and higher luminosity objects, as well as the discovery of low optical luminosity objects like Pictor~A that are efficient accretors,  require a critical evaluation of the existing models of illumination and the conditions necessary for the existence of double-peaked emission lines.  X-ray observations of the new extended sample are particularly important, since they allow us to obtain more accurate estimates of the bolometric luminosities of the sources and to gauge the need for extra illumination of the accretion disk. Double-peaked emitters which are efficient accretors likely require an alternative mode of illumination. The development of new scenarios can be guided by observations in the X-ray band. Basic X-ray absorption constraints, for example, provide sensitive probes of the nuclear conditions: (1) they allow us to see any intervening material right up to the AGN core, and (2) they are sensitive to molecular, atomic, or ionized gas and dust. The \\hbox{X-ray} absorption constraints  could provide sensitive tests of specific models proposed to explain the expanded population of double-peaked emitters. A recent model by  \\citet{jet}, for example, aims to explain the disk illumination necessary for the existence of the double-peaked lines with the presence of jets in RL, and slower, less-collimated outflows in RQ double-peaked emitters. The predictions of this model can be tested directly by obtaining X-ray spectroscopy of double-peaked emitters and looking for signs of these outflows in absorption.  In addition to testing specific models of illumination, comparison of the X-ray emission of double-peaked emitters to that of typical AGNs is essential for uncovering the necessary and sufficient conditions for the dominance of accretion-disk Balmer-line emission in this type of AGN. Such findings, which will also elucidate the physical conditions in AGNs with the broadest  observed Balmer lines, will have implications for the (still uncertain) structure of the broad line region (BLR) in AGNs in general. It is conceivable that the accretion disk is part of the BLR in all AGNs, but its contribution dominates the Balmer-line emission only in the small subsample of AGNs with the broadest lines \\citep[see also][for some tantalizing ideas about the structure of the BLR]{Laor07}.  In this paper we present the first results of an exploratory campaign to map the X-ray properties of a representative sample of AGNs with the broadest double-peaked Balmer-lines. Using the \\emph{Chandra} X-ray Observatory we obtained high-quality X-ray images and spectra for five sources, complemented by contemporaneous optical spectra from the ground. Our five targets span the redshift and optical luminosity plane of the SDSS double-peaked sample presented in S03. The paper is organized as follows: in \\S\\ref{sample} we present the current sample and its selection criteria, followed by the analysis of the X-ray and optical observations of each individual source in \\S\\S\\ref{analysis} and \\ref{optspec}, and the discussion and conclusions in \\S\\ref{conclusions}.  Throughout this paper we use the \\citet{Spergel07} concordance cosmology and quote all flux and luminosity errors as 90\\% confidence limits; all other errors quoted are 1$\\sigma$ Gaussian equivalent.  \\begin{figure*} \\plotone{f3.eps} \\caption{The H$\\alpha$ (first four) and H$\\beta$ (middle right panel) line profiles of the five double-peaked emitters observed with \\emph{Chandra} overlaid with the best-fit accretion-disk emission models. The panels are arranged in order of increasing RA, with the RA increasing down and to the right. Each panel shows the fit residuals, including the indicated narrow lines, displaced to negative values for clarity. The bottom right panel shows the H$\\alpha$ variability of one of our targets, SDSS\\,J1407$-$0235. The top curves denote the HET and SDSS spectra, while the bottom two curves show the accretion-disk fits; for clarity, the continuum levels have been displaced by arbitrary values in all cases. \\label{Ospec}} \\end{figure*} ",
        "conclusions": "\\label{conclusions}  \\subsection{X-ray Continua}  Four of the five \\emph{Chandra} double-peaked emitters are well fit by power-law models with spectral indices in the range $1.6<\\Gamma<1.8$, in good agreement with the continua of similar luminosity AGNs, $\\left< \\Gamma \\right> = 1.74\\pm0.09$, studied by \\citet{piconcelli}. A well-known correlation exists between $\\Gamma$ and the FWHM the H$\\beta$ line in AGNs \\citep[e.g.,][]{Brandt97}. \\citet{Shemmer06,Shemmer08} show that the $\\Gamma$--FWHM$\\si{H$\\beta$}$ relation is driven by the relation between the hard X-ray band photon index and the accretion rate (parameterized by $L/L\\si{Edd}$), since FWHM$\\si{H$\\beta$}$ is an accretion-rate indicator for unobscured AGNs. There is no evidence that the Balmer-line widths can be used as an accretion-rate indicator for double-peaked emitters, especially for the double-peaked emitters with extremely large line widths considered here. However, if the X-ray power-law photon index can be used as an indicator of the accretion rate, the double-peaked emitters studied here would have $L/L\\si{Edd}\\approx 0.1$ (with a $1\\sigma$ range of $0.005<L/L\\si{Edd}<1$) and at least four of the five would be considered efficient accretors, a result that is supported further by the existence of ``big blue bumps\" -- a characteristic feature in the spectral energy distributions of quasars (see \\S~\\ref{bollum}).  \\subsection{X-ray and Optical Variability}  Four of the five double-peaked emitters observed with \\emph{Chandra} are X-ray variable on long timescales. Both SDSS\\,J0132$-$0952 and SDSS\\,J1407$+$0235 have RASS upper limits lower than the observed ACIS-S fluxes in the \\hbox{0.5--2\\,keV} band, suggesting that these sources have brightened on decade timescales by at least a factor of $\\gtrsim1.6$ and $\\gtrsim4$, respectively. SDSS\\,J1238$+$5325 has the same X-ray flux during the two observations separated by 12 years in the AGN rest-frame. The remaining two sources are observed in a fainter state with \\emph{Chandra} than their corresponding early 1990s \\emph{ROSAT} PSPC  detections by factors of  about 3 (SDSS\\,J1004$+$4801) and 6 (SDSS\\,J2125$-$0813). For comparison, \\citet{Shinozaki} find that the typical flux-variability amplitude of similar luminosity AGNs on timescales of decades is $\\sim2.5$. \\citet{Mateos} compute the mean flux-variability amplitude, $\\sigma\\si{intr}$ \\citep[for a description of this quantity see][]{Almani}, of the unobscured AGNs from the Lockman Hole sample. They find that the average AGN has a $\\sigma\\si{intr}\\sim0.3$, with a range of $0.2<\\sigma\\si{intr}<0.5$ for the AGNs with confirmed variability. The corresponding $\\sigma\\si{intr}$ values for our four variable sources are $>$0.2, 0.5, $>$0.6, and 0.7 for SDSS\\,J0132$-$0952, SDSS\\,J1004$+$4801, SDSS\\,J1407$-$0235, and SDSS\\,J2125$-$0813, respectively. Based on this small sample, it appears that most double-peaked emitters are more variable on decade timescales than the typical AGN.  We tested the \\emph{Chandra} light curves for short-term variability using the XRONOS package. It enables us to compute the RMS fractional variation  ($\\sigma\\si{RMS}$) of the count rate within each exposure after binning the data in intervals ranging from half a minute to about 15--30 minutes (allowing for at least 5 bins in each case). No short-term variability was detected. The $3\\sigma$ Gaussian limits on the variability range between $\\sigma\\si{RMS}<0.1$--0.3 for the higher signal-to-noise low-redshift objects and $\\sigma\\si{RMS}<0.4$--0.6 for SDSS\\,J2125$-$0813, depending on the length of the time interval probed.  The non-detection of short-term variability is expected, as our objects are relatively luminous in the \\hbox{0.5--10\\,keV} band (with luminosities ranging between a few times $10^{43}$\\,erg\\,s$^{-1}$ and $10^{45}$\\,erg\\,s$^{-1}$), and AGN variability amplitudes on short timescales are known to decrease with increasing luminosities \\citep[e.g.,][]{Nandra97}. \\citet{LP93} show that for AGNs with luminosities in the $10^{43}-10^{45}$\\,erg\\,s$^{-1}$ range, timescales longer than $10^4-10^6$\\,s are typically necessary to detect RMS variability of more than 10\\%; these timescales are 2--10 times longer than the exposure times in our snapshot survey.  The optical variability of the five double-peaked emitters is also significant on long timescales (see Table~\\ref{tabVar}), with flux-density changes of up to a factor of 3 over 3--4\\,yrs in the AGN rest frame in our sample, accompanied by H$\\alpha$-line flux changes of up to a factor of 2.5. Figure~\\ref{OspecFit} shows the HET and SDSS spectra of three of the five \\emph{Chandra} sources. Only two of the double-peaked emitters (SDSS\\,J0132$-$0952 and SDSS\\,J1407$+$0235) show clear H$\\alpha$ line-profile variations which are similar to those observed in double-peaked emitters with long-term variability monitoring \\citep[e.g.,][]{var1,var2}. The three higher  redshift AGNs in our sample (SDSS\\,J1238$+$5325, SDSS\\,J1407$+$0235, and SDSS\\,J2125$-$0813) vary by factors of \\hbox{$\\sim$2--3} in the 4000\\,\\AA\\ continuum in a few years. For comparison, the expected AGN variability at 4000\\,\\AA\\ for a similar time lag in the SDSS is $\\sim0.2$\\,mag, or about 60\\% \\citep{sdssvar}.  We conclude that the double-peaked emitters studied here are typically more variable on long timescales (years) in the optical (a variability range of 10--300\\%, see Table~\\ref{tabVar}) and the X-ray (factors of $\\gtrsim1.6$, 3, 1, $\\gtrsim4$, and 6, in order of increasing RA) than the average AGN. In S06 we compared the  width of the $\\alpha\\si{ox}$ residual distribution to that of the luminosity-matched \\citet{St06} subsample and found that it was broader at the 1$\\sigma$ level. This finding, which could be interpreted as an indication of larger X-ray and/or optical/UV variability for double-peaked emitters, is confirmed with the small sample studied here.  \\begin{figure} \\plotone{f5.eps} \\caption{The UV-to-X-ray spectral index vs.\\ the 2500\\,\\AA\\ monochromatic luminosity. The solid line is the \\citet{St06} $\\alpha\\si{ox}$-$l\\si{2500\\,\\AA}$ relation; the dotted part is an extrapolation. \\label{aox}} \\end{figure}  \\subsection{Distribution of the UV-to-X-ray Indices} \\label{Saox}  We use the UV-to-X-ray index, $\\alpha\\si{ox}$, defined in terms of the logarithm of the ratio of the rest-frame 2500\\,\\AA\\ and 2\\,keV monochromatic fluxes, $\\alpha\\si{ox}=-0.3838\\log[F_{\\nu}(\\textrm{2500\\,\\AA})/F_{\\nu}(\\textrm{2\\,keV})]$, to compare the X-ray emission of our sample to that of more typical AGNs with similar optical/UV monochromatic luminosities. The UV-to-X-ray indices of the five \\emph{Chandra} targets are shown as a function of their 2500\\,\\AA\\ monochromatic luminosities in Figure~\\ref{aox} and listed in Table~\\ref{tab3}. There is a tendency for the double-peaked emitters to have flatter $\\alpha\\si{ox}$ values than typical RQ AGNs. This is expected for the one RL double-peaked emitter, SDSS\\,J1238$-$5325, but is also apparent in the four RQ targets. We can characterize this discrepancy by subtracting from each observed value of $\\alpha\\si{ox}$ the value expected based on the $\\alpha\\si{ox}(l\\si{2500\\,\\AA})$ relation found by \\citet{St06} for RQ AGNs, $\\alpha\\si{ox}(l\\si{2500\\,\\AA})=-0.137l\\si{2500\\,\\AA}+2.638$, and compare this to the spread, $\\sigma\\si{$\\alpha\\si{ox}$}$, observed around the relation for normal RQ AGNs. The $\\alpha\\si{ox}$ residual values ($\\Delta\\alpha\\si{ox}$), together with the values of $\\sigma\\si{$\\alpha\\si{ox}$}$ for similar luminosity RQ AGNs are given in columns 7 and 8 of Table~\\ref{tab3}.  The $\\alpha\\si{ox}$ residuals of the four RQ double-peaked emitters are larger than or about equal to the spread observed around the $\\alpha\\si{ox}(l\\si{2500\\,\\AA})$ relation for normal RQ AGNs. The $\\alpha\\si{ox}$ values are $\\sim 2 \\sigma\\si{$\\alpha\\si{ox}$}$ too flat in the cases of SDSS\\,J0132$-$0952 and SDSS\\,J1407$+$0235, and $\\sim 1 \\sigma\\si{$\\alpha\\si{ox}$}$ too flat in the cases of SDSS\\,J1004$+$4801 and SDSS\\,J2125$-$0813.  The remaining double-peaked emitter, SDSS\\,J1238$+$5325, has  an UV-to-X-ray index which is 0.41\\,dex flatter than the value expected for similar-luminosity RQ AGN. This difference is typical for RL AGNs, which are known to have 2\\,keV monochromatic luminosities of a factor of $\\gtrsim3$ higher than those of RQ AGNs with similar UV emission \\citep[e.g.,][]{rlagn}. The combined optical and X-ray variability of the double-peaked emitters together with the uncertainties in the determination of $l\\si{2500\\,\\AA}$ is large enough to account at least partially for the $\\Delta\\alpha\\si{ox}$ values observed for SDSS\\,J1004$+$4801 (up to 0.2\\,dex), SDSS\\,J1407$+$0235 (up to 0.3\\,dex), and SDSS\\,J2125$-$0813 (up to 0.3\\,dex). For X-ray or optical variability to be the cause of the anomalous $\\alpha\\si{ox}$ values observed for these three RQ AGNs, the variability amplitude over the 2--4 months separating the \\emph{Chandra} and HET observations would have to reach the amplitude observed over years (optical) or decades (X-ray). Since the amplitude of AGN variability typically increases with time, this is not likely.  Using alternative methods to estimate $l\\si{2500\\,\\AA}$, by (1) extrapolating the measured flux density at 3700\\AA\\ to 2500\\,\\AA\\  using a uniform power-law of the form, $f_{\\lambda} \\propto \\lambda^{\\alpha_{\\lambda}}$,  with slope $\\alpha_{\\lambda}=-1.5$ \\citep{Berk}, or (2) using the \\citet{Richards06} SED template scaled to the observed fluxes in the extreme UV (see \\S~\\ref{bollum}), we find results consistent with those reported in Table~\\ref{tab3} within 0.2\\,dex (typically $<$0.1\\,dex)  in all cases. This is equivalent to an uncertainty of $<$0.1\\,dex in $\\alpha\\si{ox}$, which is inadequate to explain the observed flattening of the UV-to-X-ray index.  In addition, although variability and measurement uncertainty could be responsible for the flattening of the $\\alpha\\si{ox}$ values in individual cases, it is unlikely that all four RQ double-peaked emitters would be affected in the same way. The mean value of the $\\alpha\\si{ox}$ residuals is  $\\left<\\alpha\\si{ox}-\\alpha\\si{ox}(l\\si{2500\\,\\AA})\\right>=0.25\\pm0.04$ for the four RQ double-peaked emitters observed with \\emph{Chandra}, compared to  $\\left<\\alpha\\si{ox}-\\alpha\\si{ox}(l\\si{2500\\,\\AA})\\right>=-0.02\\pm0.01$, for the luminosity-matched subsample of RQ AGNs from \\citet{St06}. Both the Gehan and logrank tests ($T=3.2$, $P=0.1$\\% and $T=7.7$, $P<0.1$\\%, respectively; see S06 for the use of these sample-comparison tests) confirm that the current sample of 4 RQ double-peaked emitters has a statistically different distribution of $\\alpha\\si{ox}$ residuals from the one found for all AGNs with similar UV luminosities and single-peaked lines presented in \\citet{St06}. This tendency was discernible in the S06 sample (left panel of Figure~9 in S06), but not statistically significant; it was also observed in the preliminary analysis of the five \\emph{Chandra} targets in combination with a sample of  double-peaked emitters observed with \\emph{Swift} \\citep{S07}.  \\subsection{Bolometric Luminosities} \\label{bollum}  Figure~\\ref{sed} shows the spectral energy distributions (SEDs) of the five \\emph{Chandra} sources from the near-infrared (NIR) to the X-ray band. The X-ray spectra are represented as power laws with photon indices given in Table~\\ref{tab2}. For the two AGN with negligible host-galaxy contamination (SDSS\\,J1238$+$5325 and SDSS\\,J2125$-$0813) we show the HET-epoch spectra in the optical region; for the remaining three objects we plot the SDSS AGN components scaled to the flux level during the HET epoch. In all five cases  we also plot the SDSS-epoch AGN components. Comparison of the HET-epoch and SDSS-epoch spectra gives an idea of the optical variability on long timescales (4--5\\,years). The near-UV (NUV, $\\lambda\\si{eff}=2267$\\,\\AA) and far-UV (FUV, $\\lambda\\si{eff}=1516$\\,\\AA) fluxes, obtained by the GALEX mission (see Table~\\ref{tabGalex}) and corrected for Galactic extinction, are also shown. The extinction corrections, $A_{FUV}=7.9E(B-V)$ and $A_{NUV}=8.0E(B-V)$, are from \\citet{dePaz}, who used the  \\citet{Cardelli} parameterization of the Galactic extinction law. The NIR fluxes (obtained from the 2\\,MASS magnitudes reported in Table~\\ref{tab2MASS}) contain the host-galaxy contributions in all cases, and are well above the template SEDs in Figure~\\ref{sed}. If we subtract an elliptical-galaxy SED (scaled to the galaxy contribution in the SDSS spectra and corrected for the SDSS fiber apertures; see Fioc \\& Rocca-Volmerange 1997 for the galaxy SED) from the NIR fluxes, we can account for part of the discrepancy between the observed NIR fluxes and the SED templates for SDSS\\,J0132$-$0952 and SDSS\\,J1004$+$4801. In both cases a standard elliptical-galaxy template can account for $\\lesssim$0.1\\,dex in the NIR; a dusty starburst would be necessary to explain the high observed NIR luminosities. In the cases of SDSS\\,J1238$+$5325 and SDSS\\,J2125$-$0813, the NIR colors are more consistent with an AGN SED and the high observed luminosities are in line with the higher optical/UV fluxes. \\begin{figure} \\plotone{f6.eps} \\caption{NIR through X-ray SEDs of the five double-peaked emitters from the \\emph{Chandra} sample in comparison with the SED of Arp\\,102B (see \\S~\\ref{bollum} for details). The observations are shown with solid lines and points. The \\citet{Richards06} SED templates, normalized to the optical continua during the HET epoch, are shown with dashed lines in each panel. The SDSS\\,J1407$-$0235 panel also shows the RQ SED of \\citet{Elvis}, with a dashed line that is distinctly higher in the X-ray range compared to the \\citet{Richards06} SED. An SED template was created in each case by extrapolating the observed X-ray power law into the extreme UV, until it meets the optically-normalized \\citet{Richards06} template; the area under the SED template, used to compute the bolometric luminosities, is shaded in each panel. The bottom panel shows the Arp\\,102B SED in comparison to the \\citet{Richards06} SED, normalized at 3200\\,\\AA; the Arp\\,102B SED lacks a Big Blue Bump (BBB) around $\\nu\\sim10^{15.4}$\\,Hz. \\label{sed}} \\end{figure}  The average quasar SED from \\citet{Richards06}, scaled to the AGN continuum flux in the optical in each case, is also shown in Figure~\\ref{sed}. We used the SDSS AGN component, scaled to the HET-epoch flux in the three cases with significant host-galaxy contribution, and the HET spectra in the remaining two cases. The \\citet{Richards06} average SEDs are similar from the infrared (IR) to the UV region to the standard \\citet{Elvis} median SEDs (both RQ SED templates are shown in the second panel of Figure~\\ref{sed} for one of our sources, SDSS\\,J1407$+$0235), but they differ substantially in the X-ray band. The \\citet{Elvis} sample was biased toward AGNs with good S/N \\emph{Einstein} data (hence X-ray brighter than the typical AGN) and bright in the extreme UV, and \\citet{Richards06} argue that typical AGNs are redder and have lower X-ray-to-optical ratios. The \\citet{Richards06} RQ SED has an $\\alpha\\si{ox}=-1.53$ and the \\citet{Elvis} RQ SED an $\\alpha\\si{ox}=-1.38$, compared to $\\alpha\\si{ox}=-1.55$ and  $\\alpha\\si{ox}=-1.54$ for their respective $l\\si{2500\\,\\AA}$ from the \\citet{St06} relation; both SEDs have total (including the IR bump) bolometric luminosities of $2\\times10^{46}$\\,erg\\,s$^{-1}$. Note that the GALEX observations agree well with the optically-scaled \\citet{Richards06} template SEDs for the four RQ \\emph{Chandra} AGN (within 0.11\\,dex, or $<30$\\%).  The \\emph{NUV} band and the optical bracket rest-frame 2500\\,\\AA\\ in all four cases, suggesting that the $l\\si{2500\\,\\AA}$ estimates based on the optically-scaled teplate SEDs and the UV-to-X-ray indices are fairly robust (changes in $\\alpha\\si{ox}$ of less than 0.05). The RL AGN, SDSS\\,J1238$+$5325, is 0.4\\,dex more luminous in the \\emph{NUV} band during the GALEX epoch relative to the HET-epoch optically-scaled RL template SED; a similar variation at 2500\\,\\AA\\ in the 17 days (13 in the AGN rest frame) between the \\emph{Chandra} and HET observations would result in a change in $\\alpha\\si{ox}$ of 0.16, and cannot fully account for the observed flatter $\\alpha\\si{ox}$.  The Arp\\,102B SED, discussed in more detail by \\citet{ArpASCA}, is shown in the last panel of Figure~\\ref{sed} for comparison.  We use the \\emph{Advanced Satellite for Cosmology and Astrophysics} (\\emph{ASCA}) X-ray spectrum presented by Eracleous et al.~(2003a; see also Lewis \\& Eracleous 2006), the \\emph{Hubble Space Telescope} (\\emph{HST}) UV spectrum given in \\citet{Halpern96}, and the NIR and  mid-infrared (MIR) spectra published by \\citet{Riffel} and \\citet{ArpSpitzer}, respectively. \\citet{Riffel} note strong stellar features in the NIR spectrum of Arp\\,102B, which was integrated over a 700\\,pc region centered on the active nucleus. The MIR source extent, studied by \\citet{ArpSpitzer}, was around 250\\,pc; for comparison, the \\emph{HST} aperture diameter is $\\sim$400\\,pc. In the same panel of Figure~\\ref{sed}, we show the average quasar SED from \\citet{Richards06} in the infrared-to-UV region (scaled to the reddest UV wavelength of the Arp\\,102B \\emph{HST} spectrum, $\\sim$3200\\,\\AA) together with the \\emph{ASCA} power law in the X-ray region. In addition, the NIR emission of Arp\\,102B includes a strong galaxy contribution, while in the MIR variability by a factor of 2.5 over a five year period was noted; both variability and varying host-galaxy contributions are likely responsible for part of the discrepancy between the quasar SED and the Arp\\,102B IR data. The Arp\\,102B SED is characterized by a lack of a `big blue bump' (BBB), the excess emission between the X-ray and UV bands ($\\nu\\sim10^{15.4}$\\,Hz) normally associated with the innermost parts of the accretion disk (which is likely replaced by a hot radiatively inefficient central structure in the case of Arp\\,102B and other AGNs with RIAFs). In contrast to the Arp\\,102B SED, the GALEX observations of the five double-peaked emitters presented here show that the BBB is present in all five SEDs, with the possible exception of SDSS\\,J1004$+$4801. The presence of BBBs is further confirmation for the relatively high efficiency of the accretion flows in these AGNs, excluding RIAF models for them.  Using the template SEDs from Figure~\\ref{sed}, we can estimate the bolometric luminosities of the double-peaked emitters which we list in the last column of Table~\\ref{tab3}. The \\citet{Richards06} SED does not extend to the radio domain ($10^9\\textrm{\\,Hz}<\\nu<10^{12.5}\\textrm{\\,Hz}$), but the $\\nu<10^{12.5}$\\,Hz region is energetically unimportant and does not affect the bolometric luminosity estimates. The template SED bolometric luminosities will overestimate the true luminosities if the IR bump consists of fully reprocessed nuclear continuum emission and should be excluded when computing the total luminosity.  Such `double-counting' of the IR emission can result in overestimates of the bolometric luminosity of up to 20--30\\%. We can obtain upper limits to the bolometric luminosities and an estimate of possible uncertainties by normalizing the  IR-to-extreme-UV part of the SED template to the NIR (in all cases but SDSS\\,J1004$+$4801), the extreme UV (in the case of SDSS\\,J1238$+$5325), and/or the SDSS-epoch spectra (in the case of SDSS\\,J2125$-$0813, for example, where the SDSS-epoch flux was higher). In all cases, we find that the bolometric luminosities can be at most a factor of 2--3 higher than those reported in Table~\\ref{tab3}.  On average about 30\\% of the bolometric luminosity is emitted between \\hbox{0.5--10}\\,keV in our sample (with a range of 16--36\\%; or between 2\\% and 12\\% in the 0.5--2\\,keV X-ray band), compared to $\\sim3$\\% emitted between \\hbox{0.5--10}\\,keV for the AGN composites of \\citet{Richards06} and $\\sim8$\\% for the median RQ AGN composite of \\citet{Elvis}. This comparison is not entirely fair, since the \\citet{Richards06}  and \\citet{Elvis} SEDs are constructed using more luminous AGNs ($l\\si{2500\\,\\AA}=30.6$ and $l\\si{2500\\,\\AA}=30.5$, respectively) and $\\alpha\\si{ox}$ is luminosity dependent. If we renormalize the \\citet{Richards06} SED to the average value of  $l\\si{2500\\,\\AA}$ in our sample, $l\\si{2500\\,\\AA}=29.2$, and scale the X-ray part of the SEDs to the level inferred from the \\citet{St06}  $\\alpha\\si{ox}(l\\si{2500\\,\\AA})$ relation for typical RQ AGNs,  we find that even in this case only 7--8\\% of the bolometric luminosity is emitted between \\hbox{0.5--10}\\,keV by typical RQ AGNs. This restates the conclusion from \\S~\\ref{Saox}: the average RQ  double-peaked emitter from our \\emph{Chandra} sample has an $\\alpha\\si{ox}$ value flatter by 0.25\\,dex relative to that of a typical RQ AGN and is 0.65\\,dex more luminous at 2\\,keV, which is equivalent to emitting about 30\\% of its bolometric luminosity between \\hbox{0.5--10}\\,keV.   \\subsection{Disk Illumination} \\label{diskP}  To compute theoretically the emerging spectrum of an accretion disk, both an accurate representation of the vertical structure of the disk and a suitable approximation for the radiation transfer are necessary. \\citet{Williams80}, for example, computes the Balmer-line emission from the outer, optically thin parts of the accretion disks of cataclysmic variables, assuming a vertically homogeneous disk and local termodynamic equilibrium (LTE). This paper finds that the outer regions of the disk cool through line emission (with $<20$\\% of the locally available viscous energy emitted as H$\\alpha$), producing the Balmer-line series with small Balmer decrements ($F\\si{H$\\alpha$}/F\\si{H$\\beta$}\\approx1$). For comparison, the observed Balmer-line decrements range between 3 and 4 for the four low-redshift \\emph{Chandra} targets (for which both H$\\alpha$ and H$\\beta$ fall within the SDSS wavelength coverage) similar to the Balmer decrement of 3.2 which is considered typical for the BLR gas.  Assuming a homogeneous photoionized disk, \\citet{CS87} and \\citet{dumont90} argued that standard accretion disks in AGNs can produce low-ionization lines with characteristics similar to those observed in the BLRs of AGNs (including steeper Balmer-line decrements) if the disk is externally illuminated. The lines arise in the radiatively excited skin of the inner ($10^2 R\\si{G}<R<10^4 R\\si{G}$) accretion disk or the radiatively heated body of the outer ($R>10^4 R\\si{G}$) disk.  The need for external illumination of the accretion disk is an important constraint on models aiming to explain the existence of double-peaked Balmer lines in AGNs. We can use the ratio of the H$\\alpha$ luminosity to the viscous power available locally in the line-emitting region of the disk, $L\\si{H$\\alpha$}/W\\si{d}$, in comparison with the \\citet{Williams80} computations, to estimate the need for external illumination. The viscous power available locally in the line-emitting region of the disk, $W\\si{d}$, can be estimated using an estimate of the bolometric luminosity and the extent of the emission-line region as described in, e.g., E03 and S06. In this work we use the bolometric luminosity estimates based on the observed spectral energy distributions (SEDs) presented in \\S~\\ref{bollum}. \\begin{figure} \\plotone{f7.eps} \\caption{The ratio of the X-ray flux to the H$\\alpha$ line luminosity, $L\\si{H$\\alpha$}$,  vs. the ratio of  $L\\si{H$\\alpha$}$ to the viscous power released locally by the accretion disk, $W\\si{d}$. The open squares denote the RL sample of E03, and the solid (open) circles indicate the RQ (RL) \\emph{Chandra} double-peaked emitters. The dashed line, at $L\\si{H$\\alpha$}/W\\si{d}=0.2$, corresponds to the \\citet{Williams80} prediction for H$\\alpha$-line cooling in a homogeneous, optically-thin LTE disk (see \\S~\\ref{diskP} for details); for objects lying above this line external illumination is necessary to produce H$\\alpha$ emission of the observed strength. The apparent anti-correlation seen in this figure is just a consequence of the fact that  $W\\si{d}$ is derived from the bolometric luminosity which correlates with the X-ray flux. \\label{Power}} \\end{figure}  E03 and S06 have found that for most double-peaked emitters $L\\si{H$\\alpha$}/W\\si{d}$ ranges between 0.1 and 3. The $L\\si{H$\\alpha$}/W\\si{d}$ ratios for the five double-peaked emitters studied here are shown in Figure~\\ref{Power} in comparison to the 20 double-peaked emitters from Table~7 of E03.\\footnote{We do not show the corresponding ratios for the S06 objects as the estimates were based on limits of the inner and our radii of the emitting regions rather than exact profile fits.}  Note that the apparent anti-correlation seen in this figure is just a consequence of the fact that  $W\\si{d}$ is derived from the bolometric luminosity which correlates with the X-ray flux. For $L\\si{H$\\alpha$}/W\\si{d}\\gtrsim1$, the energy released locally by viscous dissipation cannot power the observed line even in principle. This is the case for one of our \\emph{Chandra} targets, SDSS\\,J1004$+$4801, and 7 out of the 20 (35\\%) double-peaked emitters presented by E03.  Three of the remaining four \\emph{Chandra} double-peaked emitters (as well as 9 out of the 20 E03 sources) have $0.2 \\lesssim L\\si{H$\\alpha$}/W\\si{d} \\lesssim 1$. Thus, the observed H$\\alpha$ luminosities  are uncomfortably close to the total energy available locally, and larger than the 20\\% upper limit of \\citet{Williams80}, suggesting that external illumination plays a role. The only exception is the most luminous SDSS double-peaked emitter, SDSS\\,J2125$-$0813; for a standard H$\\alpha$ to H$\\beta$ luminosity ratio of 3.2, the measured  $L\\si{H$\\beta$}/W\\si{d}=0.04$ corresponds to  $L\\si{H$\\alpha$}/W\\si{d}\\approx0.12$, and no external illumination is necessary.  The total soft-band X-ray luminosity is a factor of 10--30 higher than the H$\\alpha$ luminosities for all five double-peaked emitters (and ranges between ~0.7 and 100 for the E03 sources), showing that enough energy is released within a few $R\\si{G}$ of the black hole to illuminate the Balmer-line-emitting portion of the disk and produce line-emission with the observed strength.  Assuming that the disk is photoionized and a thin skin emits the lines, the H$\\alpha$ power depends on the illuminating flux and the reprocessing efficiency (i.e., what fraction of the cooling comes out as H$\\alpha$). \\citet{CD89} show that, in the case of a photoionized disk, about 20\\% of the ``effective continuum power\", defined as $L\\si{E}=\\nu L_{\\nu} (3646\\textrm{\\,\\AA})+\\nu L_{\\nu} (10\\textrm{\\,keV})$, could be emitted as H$\\alpha$. In all five cases studied here, the H$\\alpha$-line luminosities are equal (in the case of SDSS\\,J1238$+$5325) or much smaller (by factors of 7, 6, 16, and 66, respectively, for the remaining four cases in order of increasing RA) than 20\\% of the effective continuum power.  We conclude that external illumination is necessary in at least four of the five cases studied here, since the locally available viscous energy is insufficient and the Balmer-line decrements too steep, compared to the \\citet{Williams80} predictions. A photoionized disk as described in e.g., \\citet{CD89}, can produce H$\\alpha$-line emission of the observed strength with the proper Balmer-line decrements.   \\subsection{Intrinsic Absorption}  Based on the X-ray spectral fits presented in \\S~\\ref{analysis}, large column density neutral absorbers are ruled out in all cases. As a simple test for the presence of ionized absorbers, we added one or more edges at $\\sim0.7-0.8$\\,keV to all model fits; this failed to improve the fits at significance levels $>90$\\% in all cases, with the exception of SDSS\\,J0132$-$0952. SDSS\\,J0132$-$0952 (see \\S~\\ref{spec6770}), requires ionized intrinsic absorption with $N\\si{H,i} \\approx 4\\times10^{22}$\\,cm$^{-2}$ and $\\xi \\approx160$\\,erg\\,s$^{-1}$\\,cm, but both quantities are poorly constrained. Silimarly, although not required by the data, the limited signal-to-noise ratio in the case of SDSS\\,J2125$-$0813 allows for the addition of a $N\\si{H,i}\\sim7\\times10^{22}$\\,cm$^{-2}$, $\\xi \\sim 200$\\,erg\\,s$^{-1}$\\,cm ionized intrinsic absorber.  Recent work by \\citet{jet} casts doubt on the suitability of a RIAF to illuminate efficiently the outer portions of the accretion disk, suggesting instead that the required illumination can be provided by low Lorentz factor ($<2$) jets in RL objects or slow outflows in the case of RQ double-peaked emitters. If slow-moving outflows in the immediate vicinity of the black holes are responsible for the disk illumination, then X-ray observations are expected to show large X-ray absorbing columns: $N\\si{H,i}\\gtrsim10^{23}$\\,cm$^{-2}$, providing a suitable scattering depth of $\\tau\\sim0.2$, according to \\citet{jet}. As an example of the prevalence of suitable ionized absorbers in AGNs, \\citet{jet} cite the work of \\citet{Piconcelli05}, who found that 20 out of the 45 AGNs in their sample show signs of ionized absorption (parameterized by $0.7-0.8$\\,keV absorption edges in 16 of the 20 cases with $\\tau>0.1$, and XSPEC \\emph{absori} fits in the remaining 4).  The \\hbox{0.3--10\\,keV} X-ray spectra of the four RQ double-peaked emitters studied here do not require large column density ionized absorbers, with the exception of SDSS\\,J0132$-$0952, where a lower column-density absorber is present, and SDSS\\,J2125$-$0813, where a high column-density and ionization absorber could not be ruled out. The \\emph{Chandra} snapshot survey presented here has yielded spectra with a limited numbers of counts and was not designed to provide stringent constraints on the column densities and ionization states of possible ionized absorbers, but the presence of such absorbers appears unlikely in at least two of the four RQ objects from the current sample.  If high column density absorbers were common in double-peaked emitters as a class, we might also expect that the UV-to-X-ray indices measured based on X-ray photometry (e.g., S06) should reflect this by showing steeper $\\alpha\\si{ox}$ distributions (since the measured X-ray fluxes would be depressed by absorption). In fact, as argued above,  the double-peaked emitters tend to have flatter $\\alpha\\si{ox}$ values than similar optical/UV luminosity AGNs, suggesting excess X-ray emission relative to the UV emission.  We conclude that the slow-outflow model of \\citet{jet} could provide the illumination necessary to power the strong double-peaked lines observed in some double-peaked emitters, but is unlikely to be responsible for the necessary extra power in all cases."
    },
    "0808/0808.3190_arXiv.txt": {
        "abstract": "We present sensitive Very Large Array observations with an angular resolution of a few arcseconds of the $J= 1 - 0$ line of SiO in the $v$=1 and 2 vibrationally excited states toward a sample of 60 Galactic regions in which stars of high or intermediate mass are currently forming and/or have recently formed.  We report the detection of SiO maser emission in \\textit{both} vibrationally excited transitions toward only three very luminous regions: Orion-KL, W51N and Sgr B2(M). Toward all three, SiO maser emission had previously been reported, in Orion-KL in both lines, in W51N only in the $v=2$ line and in Sgr B2(M) only in the $v=1$ line.  Our work confirms that SiO maser emission in star-forming regions is a rare phenomenon, indeed, that requires special, probably extreme, physical and chemical conditions not commonly found.  In addition to this SiO maser survey, we also present images of the simultaneously observed 7 mm continuum emission from a subset of our sample of star-forming regions where such emission was detected.  This is in most cases likely to be free-free emission from compact- and ultracompact-HII regions. ",
        "introduction": "SiO maser emission has been detected toward more than a thousand asymptotic giant branch (AGB) stars and a few red supergiants \\citep[RSGs; see, e.g., ] []{Benson_etal1990, Habing1996}. These objects have luminosities larger than a few thousand of \\Lsun\\ (AGB stars) to $> 10^5$ \\Lsun\\ (RSGs), which are comparable to the luminosities of intermediate- and high-mass protostellar and young stellar objects.  Nevertheless, toward forming and young stars SiO maser emission seems to be very rarely observable.  {\\it Orion Kleinmann-Low (Orion-KL)} was the first source in the sky toward which SiO maser emission was found (in the $v=1, J=2-1$ line) by \\citet{SnyderBuhl1974}. Subsequently, the first detections of the $J=1-0$ line in the $v=1$ and 2 states, also toward, among others, Orion-KL, were made by \\citet{Thaddeusetal1974} and by \\citet{Buhletal1974}, respectively. Toward the Orion-KL region, maser emission was also found in other SiO lines \\citep[$v=1, J = 3 - 2$,][]{Davisetal1974}, including transitions from the vibrational ground state \\citep{Tsuboietal1996}, but, at least to our knowledge, not from states with $v > 2$, contrary to M-type stars.  Moreover, maser lines from the $^{29}$SiO and $^{30}$SiO isotopomers \\citep{Olofssonetal1981}, and even the very rare $^{28}$Si$^{18}$O were found \\citep{Choetal2005}. Never detected, however, was the $v=2, J= 2 - 1$ line.  The latter line also remains undetected in oxygen-rich (M-type) evolved stars, although it is found in S-type stars (in which O and C abundances are equal). This behavior is explainable by the line's pumping mechanism \\citep{Olofssonetal1981}. Since the discovery of SiO maser emission in the Orion-KL region \\citep{SnyderBuhl1974}, only two more high mass star-forming regions, W51N (= W51-IRS2) and Sgr B2(M) \\citep{Hasegawaetal1986, Ukitaetal1987} have been found to show this kind of emission.  In the past, a number of surveys have been undertaken with the goal of finding SiO masers in a larger number of star-forming regions. The number of sources surveyed for which upper limits have been published is less than two dozen and these upper limits are in the several Jy range \\citep{Jewelletal1985, BarvainisClemens1984}, barely sensitive enough to potentially detect the W51N maser, and much too shallow to detect the $\\sim1$ Jansky-strength maser in Sgr B2(M) \\citep{Hasegawaetal1986, Ukitaetal1987}.  Around luminous AGB stars/RSGs, SiO maser emission arises from a region of density around a few times $10^8$ and $10^9$ \\ccm\\ and temperature above 1000 K, which is within a few stellar radii of the photosphere \\citep{Lockettetal1992,Bujarrabal1994}.  Given these extreme requirements, one expects this emission to pinpoint in star-forming regions the \\textit{exact} location of the embedded high-mass protostar that is exciting it, which frequently is not easy to determine by other means \\citep[see discussion in][]{MentenReid1995}. This is certainly borne out by the best-studied SiO maser associated with a star formation region, that in Orion-KL. \\citet{MentenReid1995}, using simultanenous high resolution VLA observations of the 43.2 GHz Orion-KL SiO maser and weak continuum emission (from source-I), which they accurately register with a 3.8 $\\mu$m speckle image, clearly showed that at the position of ``source-I'' no infrared source is detected. Thus, all that is observed at infrared wavelengths ({\\it i.e}. the famous IRc2) is reprocessed radiation, while the position of the continuum emission, which almost certainly comes from an ionized disk surrounding the protostar, is right at the center of the SiO maser distribution \\citep{Reidetal2007}.  While source-I clearly is self-luminous, as argued by \\citet{Reidetal2007}, it is unlikely to contribute a significant faction of $\\sim 10^5$ \\Lsun\\ of the KL region.  The SiO maser in W51N seems to be different from that in Orion-KL, since it only appeared to show maser emission in the 42.8 GHz $v =2, J= 1 - 0$ line, while that in Orion-KL shows emission in the $v=1$ and $2$ lines at comparable intensities. \\citet{Moritaetal1992} located the W51N SiO maser in a dense compact molecular core mapped in NH$_3$ emission by \\citet{Ho_etal1983} and \\citet{ZhangHo1997}. Very Long Baseline Array (VLBA) radio observations by \\citet{Eisneretal2002} reveal a comparable linear extent to that observed in Orion-KL. It is clear from the latter data that, while H$_2$O masers (also imaged by Eisner et al. 2002) trace a highly luminous region on a few arcsecond (few tenths of a parsec) scales, only the SiO maser marks, as in Orion, the exact location of a self-luminous power source.  Sub-milliarcsecond resolution VLBA observations show that the W51N SiO masers may be tracing the limbs of an accelerating bipolar outflow close to the ``{\\it dominant center}'' (where are located most of the H$_2$O, OH and SiO masers) \\citep{Eisneretal2002}.  No radio continuum emission is detected toward this SiO maser. However, the radio emission from Source-I would be completely undetectable at the distance of W51 ($\\sim$ 7 kpc).  Finally, little is known about the Sgr B2(M) SiO maser, from which only the $v=1$ line had been found, except for the fact that it is located close to (but not coincident with) the radio source F, which at high resolution splits into several compact sub-sources \\citep{Gaumeetal1995}.  As in  W51N, the SiO maser is close to a compact clump of \\nhhh\\ emission imaged by \\citet{Vogeletal1987}.  In this work, we present a sensitive search for SiO maser emission toward a sample of 60 high-mass star-forming regions using the Very Large Array.  We report the detection of such emission in \\textit{both} vibrationally exited transitions ($v$=1 and 2) and \\textit{only} toward Orion-KL, W51N, and Sgr B2(M). This suggests that SiO maser emission  is indeed a very special physical phenomenon possibly only occurring within a short time period during the formation of high-mass stars. Simultaneously, we obtained moderate sensitivity data of those sources at 7 mm.  In \\S \\ref{sample} we introduce our sample and in \\S \\ref{observations} discuss the observations undertook in this study. In \\S \\ref{results} we present and discuss our SiO data and also our 7 mm continuum data. ",
        "conclusions": "We report the detection of SiO maser emission in \\textit{both} vibrationally excited transitions from massive star forming regions only towards three  very luminous regions (Orion-KL, W51N and SgrB2(M)) out of the 60 observed galactic star-forming regions. Toward all three, SiO maser emission had previously been reported, in Orion-KL in both lines, in W51N only in the $v=2$ line and in Sgr B2(M) only in the $v=1$ line.  Since we do not detect either SiO line in 57 other regions of intermediate- and high-mass star formation, our work confirms that such emission in star-forming regions is a rare phenomenon, indeed, that requires special, probably extreme, physical and chemical conditions not commonly found. In addition to this SiO data, we also present images of the simultaneously observed 7 mm continuum emission from a subset of our sample of star-forming regions where such emission was detected.  This is in most cases likely to be free-free emission from compact- and ultracompact-HII regions."
    },
    "0808/0808.3968_arXiv.txt": {
        "abstract": "A multi-target detection system XAX, comprising concentric 10 ton targets of $^{136}Xe$ and $^{129/131}Xe$, together with a geometrically similar or larger target of liquid Ar, is described.  Each is configured as a two-phase scintillation/ionization TPC detector, enhanced by a full $4\\pi$  array of ultra-low radioactivity Quartz Photon Intensifying Detectors (QUPIDs) replacing the conventional photomultipliers for detection of scintillation light.  It is shown that background levels in XAX can be reduced to the level required for dark matter particle (WIMP) mass measurement at a $10^{-10}$ pb WIMP-nucleon cross section, with single-event sensitivity below $10^{-11}$ pb.  The use of multiple target elements allows for confirmation of the $A^2$ dependence of a coherent cross section, and the different Xe isotopes provide information on the spin-dependence of the dark matter interaction. The event rates observed by Xe and Ar would modulate annually with opposite phases from each other for WIMP mass $>\\sim$100 GeV$/c^2$.  The large target mass of $^{136}Xe$ and high degree of background reduction allow neutrinoless double beta decay to be observed with lifetimes of $10^{27}$-$10^{28}$ years, corresponding to the Majorana neutrino mass range 0.01-0.1 eV, the most likely range from observed neutrino mass differences. The use of a $^{136}Xe$-depleted $^{129/131}Xe$ target will also allow measurement of the pp solar neutrino spectrum to a precision of 1-2\\%. ",
        "introduction": "This paper investigates the properties of the proposed multi-ton detector XAX\\footnote{Acronym for ``$^{129/131}Xenon-Argon-^{136}Xenon$'', pronounced ``ZAX''} comprised of three noble liquid targets. XAX is designed to detect nuclear recoils from dark matter, events from neutrinoless double beta decay, and pp solar neutrino scattering events. The target materials are:  \\vspace*{3mm} \\noindent Target 1: Liquid Xenon depleted in $^{136}Xe$ (enriched in $^{129/131}Xe$) for:  \\newenvironment{my_enumerate}{ \\begin{enumerate}[(i)] \\setlength{\\itemsep}{1pt} \\setlength{\\parskip}{0pt} \\setlength{\\parsep}{0pt}}{\\end{enumerate} } \\newenvironment{my_enumerate1}{ \\begin{enumerate}[(a)] \\setlength{\\itemsep}{1pt} \\setlength{\\parskip}{0pt} \\setlength{\\parsep}{0pt}}{\\end{enumerate} }  \\begin{my_enumerate} \\item dark matter interactions with nuclei of non-zero spin. \\item pp solar neutrino interactions (via $\\nu_{e}$ scattering). \\end{my_enumerate}  \\noindent Target 2: Liquid Xenon enriched in $^{136}Xe$ for:  \\begin{my_enumerate} \\item dark matter interactions with zero-spin nuclei. \\item neutrinoless double beta decay. \\end{my_enumerate}  \\noindent Target 3: Liquid Argon depleted in $^{39}Ar$ (or optionally liquid Neon) for: \\begin{my_enumerate} \\item comparison of dark matter interactions in Xenon with lower-A dark matter target. \\item additional pp solar neutrino target (in the case of Ne). \\end{my_enumerate}  \\vspace*{3mm} A possible configuration is sketched in Fig 1, consisting of two cylindrical, active target volumes of 2 m diameter and 2 m height: (a) 10 tons liquid $^{129/131}Xe$, surrounding 10 tons of liquid Xe enriched in $^{136}Xe$, and (b) 10 tons (or larger) liquid Ar in a separate vessel. Two cryostats are located within a single shielding enclosure which provides purified water shielding thicker than 4 meters in all directions.  Events in the water are separately monitored by photomultipliers to provide an active Cerenkov veto against the residual underground muon flux.   \\begin{figure} \\includegraphics[height=55mm]{xax_detector.jpg}% \\caption{\\label{fig:epsart} Conceptual view of XAX detector layout.} \\end{figure}  \\vspace{3mm} The expected energy spectra for the three scientific objectives (dark matter, double beta decay and pp solar neutrinos) are shown in Fig 2. For simplicity, natural Xe is assumed here as a target material. In addition, 100 GeV/$c^{2}$, 1 and 10 TeV/$c^{2}$ WIMP masses with $10^{-8}$ pb $(=10^{-44} cm^{2})$ of nuclear cross section and two-neutrino/neutrinoless double beta decays with life times of $10^{22}$/$10^{27}$ years are assumed and plotted together with $pp/^{7}Be/^{8}B$ solar neutrino spectra.  \\begin{figure*} \\includegraphics[height=90mm]{Natural_Xenon_Backgrounds_3000.jpg}% \\caption{\\label{fig:epsart} Expected energy spectrums of WIMPs, solar neutrinos, and double beta decays in case of natural Xe as a target material. 100 GeV/$c^{2}$, 1 and 10 TeV/$c^{2}$ WIMP masses of WIMP with $10^{-8}$pb nuclear cross section; $2\\nu \\beta \\beta$/$0 \\nu \\beta \\beta$ with life times of $10^{22}$/$10^{27}$ years are assumed. The observed energy is smeared out assuming a resolution of  $\\sigma=\\frac{1.5 \\%}{ \\sqrt{ E \\left( MeV \\right)}}$.} \\end{figure*}  Detection of dark matter events is based on the well-established two-phase TPC (Time Projection Chamber), giving both primary (S1) and secondary (S2) scintillation signals for each event, from which nuclear recoils can be discriminated from gamma background with less than 1\\% overlap \\cite{Aln,Ang,Ben}.  This method currently reaches a sensitivity  of $< 10^{-7}$ pb as shown in Fig 7. The purpose of this paper is to demonstrate that a factor of 1000 improvement on the existing WIMP sensitivity can be achieved in XAX by the use of larger target masses with improved self-shielding and a new ultra-low background photon detection device QUPID surrounding each target volume in a $4\\pi$  array for position sensitivity.  The same scintillation photon detection system will also be capable of recording events at 2479 keV from neutrinoless double beta decay in $^{136}Xe$, as well as the spectrum of 0-250 keV electron recoils from pp solar neutrino scattering in Xe.   The basic detection principles of the double phase TPC have been fully demonstrated on a smaller scale in previous/ongoing projects with both Xe \\cite{XeTPC} and Ar \\cite{ArTPC}. In this paper we describe, for each type of signal, how backgrounds can be reduced in much larger target masses to the new low levels needed. The most important problem, common to the noble liquid detectors, is the reduction of gamma and neutron backgrounds from the photodetector array.  To achieve this, we are developing a new device - a Quartz Photon Intensifying Detector, or QUPID, which we anticipate will have a U/Th impurity level (and hence gamma and neutron emission) a factor of $\\sim100$ lower than that of the current best low-radioactive photomultipliers.  In addition to the ionization/scintillation signal discrimination provided by a double phase TPC, the Argon portion of the detector will also employ pulse shape discrimination in order to improve the efficiency of determining nuclear recoils versus $^{39}Ar$ beta decay events. The difference between the two is based upon the singlet and triplet dimer states of liquid Argon, each of which has a unique lifetime. Thus, the timing information of the pulse provides a method for identifying nuclear and electronic events. Further information can be found in \\cite{PSD}. ",
        "conclusions": "We have studied the capabilities of XAX, a multi-ton, three-target detection system capable of achieving new levels of sensitivity in the detection and measurement of dark matter, neutrinoless double beta decay and pp solar neutrino signals. The three targets are independent but symbiotic. Two liquid Xe targets, one enriched and the other depleted in $^{136}Xe$, are arranged concentrically, allowing the outer $^{129/131}Xe$ to act both as a detector for dark matter and solar neutrinos, while providing additional shielding for the detection of neutrinoless double beta decay in the inner $^{136}Xe$.  A separate vessel of similar or larger total size houses a liquid Ar (or optionally liquid Ne) target which would provide an additional and confirmatory pp solar neutrino detector, but crucially will identify a genuine dark matter WIMP signal through the coherent $A^{2}$ dependence of a low energy cross section, which thus gives an order of magnitude larger signal in Xe than in Ar.  A key feature of the proposed system is the light collection and position sensitivity achievable by means of a $4\\pi$  array of photodetectors, together with the extremely low background levels made possible by the QUPID - a new all-quartz hybrid avalanche photodiode now under development jointly by UCLA and Hamamatsu Photonics.  The system includes both zero spin and non-zero spin Xe dark matter targets, and by inclusion of Ar (optionally Ne), provides two values of atomic number and nuclear mass for signal identification and, with sufficient events, measurement of the WIMP mass.  We have shown that background levels can be reduced or rejected to a level sufficient to identify a dark matter signal at a WIMP-nucleon cross section level as low as $10^{-11}$ pb, and to measure a dark matter energy spectrum of $> 100$ events/year at cross section $>10^{-10}$ pb, thus covering a large fraction of the favored parameter space of supersymmetry theory.  \\vspace{3mm} For the $^{136}Xe$, target, we have shown that background levels can be reduced sufficiently to reach a $0\\nu\\beta\\beta$    half life sensitivity of $\\sim10^{27}$ years if an individual QUPID activity of 1 mBq  is achieved, and  $\\sim10^{28}$ years for an individual QUPID activity of 0.1 mBq.  Thus for the first time the vital goal of positive detection of a Majorana mass level of 0.1 - 0.01 eV, in agreement with the mass level suggested by neutrino mixing results, would be achievable in XAX.  With the background reductions necessary for dark matter and $0\\nu\\beta\\beta$    decay, the detection of the pp solar neutrino flux becomes a by-product of the XAX detection system, provided that the outer Xe target is depleted in $^{136}Xe$ to reduce the $2\\nu\\beta\\beta$    background.  However, again the gamma background is pivotal, and the position sensitivity and ultra-low background of the QUPID array here plays an essential part, since it allows a software-controlled amount of self shielding to be applied to the outer $^{129/131}Xe$ detector, while the latter independently retains its function as a full shield for the inner $^{136}Xe$.  A similar software position-sensitive control of the self-shielding in the Ar or optional Ne target will play a part in optimizing dark matter and pp solar neutrino detection in that target also.  With the foreseeable background reduction, we have shown that the dominant pp solar neutrino flux can be measured to an accuracy of 1-2\\%.  \\vspace{3mm} We have shown that three of the most important particle physics and astrophysics goals can be achieved in XAX, with the component detectors at the 10 ton level.  Clearly, it is also the case that considerable advances in sensitivity can be achieved by a similar configuration with component detectors at the one ton level, although the difference is more than just the factor of 10 in mass, since the self shielding depends on absolute thickness and is less effective at lower masses.  Thus it is at the 10 ton level that the most dramatic advances in detection capability become possible, as demonstrated by the results shown throughout this paper.  \\vspace{5mm} \\noindent {\\bf Acknowledgements}  \\vspace{2mm}  The concept of achieving three goals through the use of Liquid Xe was originally inspired by Y. Suzuki and others in the XMASS collaboration.  We are especially indebted to M. Nakahata in XMASS for providing us with his code for computing the solar neutrino flux and interactions with Xe. QUPID is the result of extensive discussion and collaboration with Hamamatsu Photonics Co., in particular M. Suyama.  We also thank J. Takeuchi, T. Hakamata, S. Muramatsu and A. Fukasawa at Hamamatsu for their continuous support and development. We are thankful for useful discussions with E. Aprile, K. Giboni, K. Ni, R. Cousins, B. Sadoulet, H. Sobel, R. Gaitskell, T. Shutt, D. McKinsey, F. Calaprice and C. Galbiati. This work was supported in part by US DOE grant DE-FG-03-91ER40662 , and by NSF grants PHY-0139065/PHY-0653459 as well as the NSF REU program. This work was also supported in part by discretionary funds from UCLA Vice Chancellor R. Peccei, Dean J. Rudnick and Department Chair F. Coroniti. Lastly, the authors would especially like to thank PK Williams for his strong support and encouragement over the last two decades during his tenure at DOE."
    },
    "0808/0808.4123_arXiv.txt": {
        "abstract": "\\vspace{-0.6cm} Except in a few cases cosmic dust can be studied {\\it in situ} or in terrestrial laboratories, essentially all of our information concerning the nature of cosmic dust depends upon its interaction with electromagnetic radiation. This chapter presents the theoretical basis for describing the optical properties of dust -- how it absorbs and scatters starlight and reradiates the absorbed energy at longer wavelengths. ",
        "introduction": "} \\vspace{-2mm} Dust is everywhere in the Universe: it is a ubiquitous feature of the cosmos, impinging directly or indirectly on most fields of modern astronomy. It occurs in a wide variety of astrophysical regions, ranging from the local environment of the Earth to distant galaxies and quasars: from meteorites originated in the asteroid belt, the most pristine solar system objects -- comets, and stratospherically collected interplanetary dust particles (IDPs) of either cometary or asteroidal origin to external galaxies (both normal and active, nearby and distant) and circumnuclear tori around active galactic nuclei; from circumstellar envelopes around evolved stars (cool red giants, AGB stars) and Wolf-Rayet stars, planetary nebulae, nova and supernova ejecta, and supernova remnants to interstellar clouds and star-forming regions; from the terrestrial zodiacal cloud to protoplanetary disks around young stellar objects and debris disks around main-sequence stars ...  Dust plays an increasingly important role in astrophysics. It has a dramatic effect on the Universe by affecting the physical conditions and processes taking place within the Universe and shaping the appearance of dusty objects (e.g. cometary comae, reflection nebulae, dust disks, and galaxies) (i) as an absorber, scatterer, polarizer, and emitter of electromagnetic radiation; (ii) as a revealer of heavily obscured objects (e.g. IR sources) of which we might otherwise be unaware; (iii) as a driver for the mass loss of evolved stars; (iv) as a sink of heavy elements which if otherwise in the gas phase, would profoundly affect the interstellar gas chemistry; (v) as an efficient catalyst for the formation of H$_2$ and other simple molecules as well as complex organic molecules (and as a protector by shielding them from photodissociating ultraviolet [UV] photons) in the interstellar medium (ISM); (vi) as an efficient agent for heating the interstellar gas by providing photoelectrons; (vii) as an important coolant in dense regions by radiating infrared (IR) photons (which is particularly important for the process of star formation in dense clouds by removing the gravitational energy of collapsing clouds and allowing star formation to take place); (viii) as an active participant in interstellar gas dynamics by communicating radiation pressure from starlight to the gas, and providing coupling of the magnetic field to the gas in regions of low fractional ionization; (ix) as a building block in the formation of stars and planetary bodies and finally, (x) as a diagnosis of the physical conditions (e.g. gas density, temperature, radiation intensity, electron density, magnetic field) of the regions where dust is seen.  The dust in the space between stars -- interstellar dust -- is the most extensively studied cosmic dust type, with circumstellar dust (``stardust''), cometary dust and IDPs coming second. Interstellar dust is an important constituent of the Milky Way and external galaxies. The presence of dust in galaxies limits our ability to interpret the local and distant Universe because dust extinction dims and reddens the galaxy light in the UV-optical-near-IR windows, where the vast majority of the astronomical data have been obtained. In order to infer the stellar content of a galaxy, or the history of star formation in the Universe, it is essential to correct for the effects of interstellar extinction. Dust absorbs starlight and reradiates at longer wavelengths. Nearly half of the bolometric luminosity of the local Universe is reprocessed by dust into the mid- and far-IR.  Stardust, cometary dust and IDPs are directly or indirectly related to interstellar dust: stardust, condensed in the cool atmospheres of evolved stars or supernova ejecta and subsequently injected into the ISM, is considered as a major source of interstellar dust, although the bulk of interstellar dust is not really stardust but must have recondensed in the ISM \\cite{Draine_1990}. Comets, because of their cold formation and cold storage, are considered as the most primitive objects in the solar system and best preserve the composition of the presolar molecular cloud among all solar system bodies. Greenberg \\cite{Greenberg_1982} argued that comets are made of unaltered pristine interstellar materials with only the most volatile components partially evaporated, although it has also been proposed that cometary materials have been subjected to evaporation, recondensation and other reprocessing in the protosolar nebula and therefore have lost all the records of the presolar molecular cloud out of which they have formed. Genuine presolar grains have been identified in IDPs and primitive meteorites based on their isotopic anomalies \\cite{Clayton_Nittler_2004}, indicating that stardust can survive journeys from its birth in stellar outflows and supernova explosions, through the ISM, the formation of the solar system, and its ultimate incorporation into asteroids and comets.  Except in a few cases cosmic dust can be studied {\\it in situ} (e.g. cometary dust \\cite{Kissel_2007}, dust in the local interstellar cloud entering our solar system \\cite{Grun_2005}) or in terrestrial laboratories (e.g. IDPs and meteorites \\cite{Clayton_Nittler_2004}, cometary dust \\cite{Brownlee_2006}), our knowledge about cosmic dust is mainly derived from its interaction with electromagnetic radiation: extinction (scattering, absorption), polarization, and emission. Dust reveals its presence and physical and chemical properties and provides clues about the environment where it is found by scattering and absorbing starlight (or photons from other objects) and reradiating the absorbed energy at longer wavelengths.  This chapter deals with the optical properties of dust, i.e., how light is absorbed, scattered, and reradiated by cosmic dust. The subject of light scattering by small particles is a vast, fast-developing field. In this chapter I restrict myself to astrophysically-relevant topics. In \\S\\ref{sec:Phys_Scat} I present a brief summary of the underlying physics of light scattering. The basic scattering terms are defined in \\S\\ref{sec:Scat_Def}. In \\S\\ref{sec:Diel_Func} I discuss the physical basis of the dielectric functions of dust materials. In \\S\\ref{sec:Scat_Approx} and \\S\\ref{sec:Scat_Tech} I respectively summarize the analytic and numerical solutions for calculating the absorption and scattering parameters of dust.  \\begin{figure} \\centering \\vspace{-2mm} \\includegraphics[height=9.5cm]{li_scat1.cps} \\vspace{-3mm} \\caption{ \\label{fig:scat1} A conceptual illustration of the scattering of light by dust. The dust is conceptually subdivided into many small regions. Upon illuminated by an incident electromagnetic wave, in each small region a dipole moment (of which the amplitude and phase depend on the composition the dust) will be induced. These dipoles oscillate at the frequency of the incident wave and scatter secondary radiation in all directions. The total scattered field of a given direction (e.g. $P_1$, $P_2$) is the sum of the scattered wavelets, where due account is taken of their phase differences. Since the phase relations among the scattered wavelets change with the scattering direction and the size and shape of the dust, the scattering of light by dust depends on the size, shape, and chemical composition of the dust and the direction at which the light is scattered. } \\end{figure}  \\vspace{-4mm} ",
        "conclusions": ""
    },
    "0808/0808.0193_arXiv.txt": {
        "abstract": "We present deep imaging of the star-forming dwarf galaxy IC~2574 in the M81 group taken with the Large Binocular Telescope in order to study in detail the recent star-formation history of this galaxy and to constrain the stellar feedback on its HI gas. We identify the star-forming areas in the galaxy by removing a smooth disk component from the optical images. We construct pixel-by-pixel maps of stellar age and stellar mass surface density in these regions by comparing their observed colors with simple stellar populations synthesized with STARBURST99. We find that an older burst occurred about 100 Myr ago within the inner 4 kpc and that a younger burst happened in the last 10 Myr mostly at galactocentric radii between 4 and 8 kpc. We compare stellar ages and stellar mass surface densities with the HI column densities on subkiloparsec scales. No correlation is evident between star formation and the atomic H gas on local scales, suggesting that star formation in IC~2574 does not locally expel or ionize a significant fraction of HI. Finally, we analyze the stellar populations residing in the known HI holes of IC~2574. Our results indicate that, even at the remarkable photometric depth of the LBT data, there is no  clear one-to-one association between the observed HI holes and the most recent bursts of star formation in IC~2574. This extends earlier findings obtained, on this topic, for other dwarf galaxies to significantly fainter optical flux levels. The stellar populations formed during the younger burst are usually located at the periphery of the HI holes and are seen to be younger than the holes dynamical age. The kinetic energy of the holes expansion is found to be on average 10$\\%$ of the total stellar energy released by the stellar winds and supernova explosions of the young stellar populations within the holes. With the help of control apertures distributed across the galaxy we estimate that the kinetic energy stored in the HI gas in the form of its local velocity dispersion is about 35$\\%$ of the total stellar energy (and 20$\\%$ for the HI non-circular motions), yet no HI hole is detected at the position of these apertures. In order to prevent the HI hole formation by ionization, we estimate an escape fraction of ionizing photons of about 80$\\%$. ",
        "introduction": "It is common use to summarize with the term ``star formation'' how stars form from the neutral gas reservoir of a galaxy, and how they evolve and affect the ambient interstellar medium (ISM). Through their winds and supernova explosions, stars release kinetic energy and freshly synthesized metals to the ISM, with the possible effect of locally suppressing the formation of new stars and triggering it on somewhat larger scales across their host galaxy (i.e. stellar feedback). The global properties of disk galaxies have been extensively studied by now, to find that the average star formation rate (SFR) of a galaxy tightly correlates with its gas surface density (Schmidt 1959, 1963, Kennicutt 1998a,b). This important result alone, however, does not put strong constraints on the physics of star formation itself; it has been indeed shown that different mechanisms can produce this observed correlation (e.g. Silk 1997, Kennicutt 1998b, Elmegreen 2002). Moreover, our knowledge of star formation is mainly based on local disk galaxies, chemically and dynamically different from local, dwarf irregulars which sometimes undergo intense starbursts (Heckman 1998). Because of their size and low metallicity, dwarf galaxies are considered to be the local analogs to the galaxy population at high redshift, where starbursts are observed to be a rather recurrent phenomenon (Steidel et al. 1996, Pettini et al. 2001, Blain et al. 2002, Scott et al. 2002). Therefore, understanding the details of star formation in nearby galaxies of different Hubble types is crucial for boot-strapping theories of galaxy formation and evolution. This is now made possible by a new generation of instruments which provide high angular resolution throughout the whole spectral range, and especially for the observations of atomic H gas at 21 cm (HI). These new data allow us to spatially resolve the star-formation activity and the gas distribution within a galaxy, and to analyze their correlation on small scales and across a wide range of local, physical conditions. This approach has been recently applied by Kennicutt et al. (2007) to M~51a and to a larger sample of galaxies by Bigiel et al. (2008) and Leroy et al. (2008). \\par In this paper, we apply a similar method  to the dwarf galaxy IC~2574 with the aim of describing its star-formation activity (i.e. age and mass surface density of its stellar populations) and comparing those properties with the interstellar medium (i.e. HI) on a $\\simeq$ 100 parsec scale. \\par IC~2574 (aka UGC~5666, DDO~81) is a gas-rich dwarf, actively forming stars in the M81 group. Miller \\& Hodge (1994) observed it in the H$\\alpha$ emission and identified 289 HII regions, each about 50 pc in diameter. The largest complex of star-forming regions is located in the North-East of the galaxy (also known as a HI supergiant shell, Walter et al. 1998, Walter \\& Brinks 1999). The authors estimated a global star-formation rate SFR $\\simeq$ 0.08 M$_{\\odot}$yr$^{-1}$ and concluded that IC~2574 is forming stars at roughly its average rate. Follow-up spectroscopy on some HII regions in the NE complex by Tomita et al. (1998) showed that the H$\\alpha$ velocity field is chaotic in all regions except two (IC2574-I and IC2574-IV), which are characterized by a V-shaped velocity distribution. This V-shaped feature is usually representitative of stellar winds blowing out the interstellar medium. Interestingly enough, Drissen et al. (1993) found three candidate Wolf-Rayet (WR) stars in IC2574-I. WRs are evolved massive stars known to develop stellar winds stronger than O stars of the same luminosity and to make up as much as 50$\\%$ of the total mass and kinetic energy released by all massive stars into the interstellar medium (cf. Leitherer, Robert \\& Drissen  1992). Therefore, their presence in IC2574-I can well explain the bulk motion observed in this HII region, and also date the region to be a few Myrs old, given that the WR  phase is as short as 2 - 5 Myr (Robert et al. 1993).  \\begin{figure*} \\epsscale{.70} \\plotone{f1.eps} \\caption{Color composite of IC~2574 obtained with the LBT LBC Blue camera in the Bessel $U$ (in blue), $B$ (in green) and $V$ (in red) filters, during the Science Demonstration Time of February 2007. North is up and East to the left. The size of the image is 14' $\\times$ 11', and corresponds to the central one fourth of the LBC field of view.} \\end{figure*}  The most intriguing property of IC2574 is, however, the spatial distribution of its neutral interstellar medium. The high-resolution VLA observations of IC~2574 have resulted in a very structured HI map, with 48 holes ranging in size from $\\sim$ 100 pc to $\\sim$ 1000 pc. These features appear to be expanding with a radial velocity of about 10 km~s$^{-1}$, which, in turn, defines dynamical ages and kinetic energies consistent with being driven by stellar winds and supernova explosions (Walter \\& Brinks 1999). In particular, the comparison of the HI map with ultraviolet (UIT) and optical images has revealed that the HII regions located in the NE complex are distributed along the rim of the HI supergiant shell, and the shell itself is devoided of neutral and ionized gas, possibly filled with hot X-ray gas (Walter et al. 1998). Stewart \\& Walter (2000) identified a star cluster within the supergiant shell with no associated H$\\alpha$ emission  and with an age (11 Myr) comparable to the shell dynamical age of 14 Myr. They dated the H$\\alpha$-emitting rim of the shell to be about 3 Myr old. The authors thus concluded that the HI supergiant shell was created by this star cluster and the shell expansion triggered the star formation activity along its rim. As pointed out by Walter \\& Brinks (1999), the occurrence of H$\\alpha$ emission along the rim of a HI feature is common to all the HI holes found in IC~2574, even if no remnant central star cluster has been identified so far. The question is then whether these holes are the result of stellar feedback as in the case of the supergiant shell or are produced by different mechanisms. The data collected so far for the Magellanic Clouds (Kim et al. 1999, Hatzidimitriou et al. 2005), M31 and M33 (cf. van der Hulst 1996) and Holmberg II (Rhode et al. 1999) offer very little support for a one-to-one correlation between HI holes and OB associations, although the kinetic energy of the hole expansion is consistent with the expansion being triggered by the total energy released by a single stellar population. One possible explanation that has been put forward is that at least some HI holes (especially those with no indication of recent star formation) may have been produced by mechanisms unrelated to young stars. For example, Loeb \\& Perna (1998) and Efremov, Elmegreen \\& Hodge (1998) suggest they arise from  Gamma-Ray bursts. HI holes may also be produced by the collision of high-velocity clouds of neutral gas with the galactic disk (cf. Tenorio-Tagle 1980, 1981), or, more simply, by the intrinsic turbulence of the interstellar medium (Elmegreen \\& Efremov 1999, Dib, Bell \\& Burkert 2006). Rhode et al. (1999) proposed the intergalactic UV radiation field as an additional mechanism to create holes in the HI gas by keeping it ionized (at least in the outskirts of the galactic disk where the HI volume density is lower), while Bureau \\& Carignan (2002) suggested that the HI holes detected in Holmberg II may have formed by ram pressure from the intragroup medium. \\par One could also imagine that the lack of a one-to-one correlation between HI holes and OB associations may be due to the limited photometric depth of the data. For example, Rhode et al. (1999) reached a limiting magnitude $B \\simeq$ 23 mag in Holmberg II and may have failed to detect  fainter young stars (i.e. late B spectral types). We have taken advantage of the high sensitivity and wide field of view of the LBC Blue camera mounted on the Large Binocular Telescope to perform deep optical imaging, to study in detail the stellar populations in IC~2574  and to investigate the possible {\\it stellar origin} of the HI holes in this galaxy. Our aim is to identify low-luminosity star clusters within these features, to establish how their stellar content may be different from any other position across the galaxy, and whether the energy budget of the young stars in the HI holes can account for the kinetic energy of these structures. This study allows us also to derive the star-formation history of IC~2574 locally, at different galactocentric radii and in correspondence with different HI column densities. The observations and the data reduction are described in Sect. 2. The high-resolution, optical morphology of IC~2574 is presented in Sect. 3, while the stellar ages and masses resulting from our dating technique are discussed in Sect. 4. We describe the stellar populations associated with the HI holes in Sect. 5. and draw our conclusions in Sect. 6. ",
        "conclusions": "Deep, LBT imaging in the $U$, $B$ and $V$ bands has allowed us to spatially map the star formation history of IC~2574 over the last $\\sim$10$^8$ yrs. We do this by analyzing the pixel-by-pixel colors of all areas in IC~2574 with significant detections in all bands. Overall, the angle averaged surface-brightness profiles indicate a predominance of blue colors [($U - B$) $<$ -0.3 mag and ($B - V$) $<$ 0.4 mag] at galactocentric radii larger that 4 kpc, mostly due to the presence of the large North-East complex of HII regions (associated with a HI supergiant shell) and the South-West extended tail of star formation. At a limiting brightness of 26 mag~arcsec$^{-2}$, the major axes of the HI and $U$-band light distributions are nearly the same, while the minor axis is smaller in the optical. This could be due to differences in the scale height of the $U$-band light and HI distributions.  Using the dating technique of Pasquali et al. (2003), we are  able to construct pixel-by-pixel maps of stellar age and mass surface density from the comparison of the observed colors with simple stellar population models synthesized with STARBURST99. After fitting and subtracting a smooth disk component (whose colors point to a stellar age $>$ 200 Myr), our best-fitting STARBURST99 models indicate the occurrence of two major bursts, one about 100 Myr ago and the other during the last 10 Myr. The older burst appears to be spatially ``confined'' within the inner 4 kpc (in galactocentric radius), while the younger episode of star formation occupies the 4 $< R <$ 8 kpc region. The two episodes differ also in the stellar mass surface density of stars formed: the typical mass density associated with the older burst is about 1 M$_{\\odot}$pc$^{-2}$, while it is $\\simeq$0.04 M$_{\\odot}$pc$^{-2}$ for the younger burst. However, these values are luminosity-weighted and therefore biased. Given that the $U$-band flux emitted by a single stellar population decreases for increasing age and the flux itself is proportional to the stars total mass, the detection of stars above a fixed threshold in the $U$-band (used in this study as a prior) requires stellar populations progressively more massive as they grow older. Quantitatively speaking, a limiting surface brightness of 26 mag~arcsec$^{-2}$ allows us to detect stellar populations older than 10$^8$ yrs with a mass surface density larger than 0.1 M$_{\\odot}$pc$^{-2}$, while those populations few Myr young are detected down to a mass surface density of 0.003 M$_{\\odot}$pc$^{-2}$ (a factor of $\\sim$ 30 lower than at 100 Myr). The fact that we do not detect young stellar populations as massive as 1 M$_{\\odot}$pc$^{-2}$ is not a selection effect of the data; it possibly hints at a constant star formation rate, at least over the last $\\sim$10$^8$ yrs, so that the younger burst has not had time yet to produce the same amount of stellar mass as was assembled during the older burst.  We compare both the stellar ages and the mass surface densities of the recently formed populations  with the HI column density all across IC~2574 and do not find any significant correlation between the star formation activity and the HI reservoir in this galaxy. A similar result was obtained by Wong \\& Blitz (2002), de Blok \\& Walter (2006), Kennicutt et al. (2007) and Bigiel et al. (2008), using different tracers of star formation. The lack of any correlation between star formation and HI is in contrast with the expectation of a relationship between HI and SFR based on both formation of and feedback from young stars. In terms of star formation, the correlation observed between H$_2$ and SFR in gas-rich spirals might indicate that the HI gas reservoir controls the SFR. In terms of stellar feedback, instead, we would expect to see a correlation if HI forms via photodissociation of H$_2$ powered by the stellar UV radiation field (Shaya \\& Federman 1987, Tilanus \\& Allen 1991, Allen et al. 1997) or an anticorrelation if the stellar radiation field ionizes HI into HII. We thus conclude that the star formation in IC~2574 is not regulated by the available HI gas reservoir, nor efficient at creating HI via photodissociation, nor actively expeling or ionizing a significant fraction of the HI.  On a global scale, star formation in IC~2574 appears to preferentially take place in regions where the HI column density is about 6 $\\times$ 10$^{20}$ cm$^{-2}$. This value is comparable with the global critical density derived by Skillman (1987, $\\sim$10$^{21}$ cm$^{-2}$) for a sample of irregular galaxies, by Kennicutt (1989, about 5 $\\times$ 10$^{20}$ cm$^{-2}$) for 33 disk galaxies, and by de Blok \\& Walter for NGC~6822 ((5.6 $\\times$ 10$^{20}$ cm$^{-2}$), and with the theoretical values computed by Schaye (2004).  The lack of correlation between the star formation activity and the HI density in IC~2574 is rather puzzling with respect to the distinct HI holes, as one would expect to systematically detect a central cluster of young stars driving the hole expansion through their stellar winds and supernova explosions. With the possible exception of the supergiant shell (cf. Stewart \\& Walter 2000), we do not observe any concentration of young stellar populations at the center of the HI holes, but rather a mixture of populations formed during the younger and older bursts, with the younger population often in clumps located at the periphery of the HI hole. We determine the fraction of pixels with a stellar age younger than 16 Myr ($f_{\\rm YOUNG}$) for 15 of the 48 holes identified by Walter \\& Brinks (1999) (i.e. those that overlap with the area of the galaxy obtained after the removal of its smooth disk component, see Sect. 4) . Here, the age of 16 Myr is assumed to separate the younger from the older star formation episode identified by us. We also compute the mean age ($<$AGE$>_{\\rm YOUNG}$) and the total energy (E$_{\\rm YOUNG}$, due to stellar winds and supernova explosions) of the stars formed during the younger burst in each hole and compare them with the dynamical age and kinetic energy of the hole (based on its expansion velocity). The relevant results of such a comparison are that: \\par\\noindent {\\it i)} $f_{\\rm YOUNG}$ (the ratio of the number of pixels associated with stellar ages younger than 16 Myr to the total number of pixels within a hole) does not correlate with the hole properties. \\par\\noindent {\\it ii)} The dynamical ages of the holes are generally higher than the younger burst by a factor of 5. The holes dynamical age is though an upper limit as it was computed by Walter \\& Brinks assuming a constant expansion rate. In reality, the holes expansion is expected to have been faster during its earlier phases. \\par\\noindent {\\it iii)} The kinetic energy of the holes is on average 10$\\%$ of the total energy released by the younger burst. This energy ratio (i.e. the efficiency of stellar feedback) is in agreement with the value (10$\\%$ - 20$\\%$) obtained by Cole et al. (1994) from fitting the galaxy luminosity functions and colors, while Bradamante et al. (1998) determined a value of 3$\\%$ for supernovae II and 100$\\%$ for supernovae Ia from the modeling of the chemical evolution of blue compact galaxies. \\par Even at the remarkable photometric depth of the LBT data, we do not find a clear one-to-one association between the observed HI holes and the most recent bursts of star formation. However, the stars formed during the younger burst do in principle release enough energy to power the expansion of the HI holes. In order to investigate better this apparent discrepancy, we map IC~2574 in optical and in HI with a set of control apertures equivalent in size to the beam of the radio observations (136 pc). No observed HI hole is included among the control apertures. For those apertures with $f_{\\rm YOUNG} >$ 70$\\%$ (just to focus on the interplay between the younger stellar component and the HI gas), we measure their average HI velocity dispersion $\\sigma_{\\rm HI}$ (from the second-moment map of Walter \\& Brinks), their average HI column density $\\Sigma_{\\rm HI}$, the total stellar energy of their stellar populations younger than 16 Myr, $E_{\\rm YOUNG}$, and the HI kinetic energy $E_{\\rm HI}$, defined as $E_{\\rm HI}$ = 1.5$\\Sigma_{\\rm HI} \\times Area \\times \\sigma_{\\rm HI}^2$ (where $Area$ is the area of the apetures). For two thirds of the control apertures the ratio $E_{\\rm HI}/E_{\\rm YOUNG}$ is lower than 100$\\%$ and has an average value of 35$\\%$; all the remaining apertures have $E_{\\rm HI}/E_{\\rm YOUNG} >$ 100$\\%$ and are systematically associated with $\\Sigma_{\\rm HI} \\geq$ 8 $\\times$ 10$^{20}$ cm$^{-2}$. Once again, the stellar populations formed during the younger burst could balance the kinetic energy stored in the HI gas, transform it into an expansion motion and thus produce a detectable HI hole for most of the control apertures. Yet, these regions failed to develop a HI hole. These same stellar populations could also sustain the HI non-circular motions, since their average ratio $E_{\\rm HI}^{\\rm ncir}/E_{\\rm YOUNG}$ is 21$\\%$.  Holes in the HI distribution may also be formed by ionization due to the UV emission of the young stars. In the case of those control apertures with $f_{\\rm YOUNG} >$ 70$\\%$, the Str\\\"omgren sphere produced by their stellar populations younger than 16 Myr is 4 times bigger in radius than the apertures and, therefore, would produce a detectable HI hole. Since this is not observed, the fraction of ionizing photons that effectively ionize the HI gas has to be rather small; a fraction $\\leq$ 20$\\%$ of the total number of ionizing photons emitted by the young stellar populations reduces the radius of the Str\\\"omgren sphere to $\\leq$ 2 times the aperture radius, making any HI hole formed via ionization less likely detectable. This result is in agreement with the lack of correlation between star formation activity and HI density seen in Sect. 4.3, indicating that star formation does not appear to significantly affect the properties of the HI gas in IC~2574. From the analysis of the stellar content of the HI holes and the control apertures it is hard to establish whether stellar/supernova feedback is at work in IC~2574 on local scales of few 100 pc; this same result has been found for other dwarf galaxies (LMC and SMC, Kim et al. 1999, Hatzidimitriou et al. 2005; M33, van der Hulst 1996; Holmberg II, Rhode et al. 1999). The ultimate explanation of these findings remains elusive; it may rest either in an overestimate of the stellar mass loss and stellar winds especially at low metallicity, or in our understanding of how stellar energy interacts with the interstellar medium. Another explanation may be found in a stochastic sampling of the IMF in low-mass stellar systems. Stellar-populations synthesis codes usually assume that the IMF is a densely sampled probabilistic distribution function, while the IMF is in reality poorly and discretely sampled in low-mass stellar systems. This assumption produces fluctuactions in the predicted properties of these systems which may mislead comparisons with the observations (Cervi\\~no \\& Valls-Gabaud 2003), by overestimating, for example, the stellar energy realesed to the ISM and thus underestimating the efficiency of stellar feedback. As shown by Cervi\\~no \\& Molla (2002), Cervi\\~no \\& Luridiana (2006) and Carigi \\& Hernandez (2008), stochastic effects on the IMF sampling become significant in low metallicity systems such as dwarf galaxies."
    },
    "0808/0808.3334_arXiv.txt": {
        "abstract": "{}{A novel method of using hard X-rays as a diagnostic for chromospheric density and magnetic structures is developed to infer sub-arcsecond vertical variation of magnetic flux tube size and neutral gas density.}{Using Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI) X-ray data and the newly developed X-ray visibilities forward fitting technique we find the FWHM and centroid positions of hard X-ray sources with sub-arcsecond resolution ($\\sim 0.2''$) for a solar limb flare. We show that the height variations of the chromospheric density and the magnetic flux densities can be found with unprecedented vertical resolution of $\\sim$ 150~km by mapping 18-250~keV X-ray emission of energetic electrons propagating in the loop at chromospheric heights of 400-1500~km.} {Our observations suggest that the density of the neutral gas is in good agreement with hydrostatic models  with a scale height of around $140\\pm 30$~km. FWHM sizes of the X-ray sources decrease with energy suggesting the expansion (fanning out) of magnetic flux tube in the chromosphere with height. The magnetic scale height $B(z)\\left(dB/dz\\right)^{-1}$ is found to be of the order of 300~km and strong horizontal magnetic field is associated with noticeable flux tube expansion at a height of $\\sim$ 900~km.}{} ",
        "introduction": "Chromospheric magnetic fields are notoriously difficult to measure and their detailed structure is effectively inaccessible with modern observations. The difficulties of various line spectroscopic techniques \\citep[e.g.][]{solanki2006} occur because the magnetic field is relatively weak so the observed spectral lines are consequently broad and insensitive to the field. The computation of chromospheric vector magnetic field from spectral lines is also an ill-conditioned inverse problem \\citep[e.g.][]{metcalf1995}. In addition, current ground based vector magnetograms have spatial resolution comparable with the vertical size of the chromosphere itself $~2-3''$ \\citep{lagg2007}. Therefore, various indirect techniques are often employed to determine the magnetic field in the chromosphere: optical observations of photospheric magnetic fields combined with extrapolation into the chromosphere \\citep[e.g.][]{mcclymont1997}; radio observations of gyroresonance emission \\citep{lang1993,aschwanden1995,vourlidas1997,white1997}.  The solar chromosphere being only about 2000~km thick ($\\sim 3''$) strategically covers the layer where the solar atmosphere turns from the gas-dominated lower chromosphere/photosphere into the magnetic field dominated upper chromosphere/corona. \\citet{gabriel1976} has proposed that magnetic field in the chromosphere fans out (canopies) and \\citet{giovanelli1982} have found that the canopy height should be typically 300--400 km. However, polarisation measurements by \\citet{landi1998} suggest a very small horizontal component of the magnetic field and \\citet{schrijver2003} argue that the ``wine-glass'' shaped magnetic field should return to the photosphere near their parent flux tube. In addition, different magnetic field models predict different canopy heights \\citep{solanki1999}.  The transport of both thermal and energetic charged particles in the solar atmosphere is governed by individual magnetic flux tubes. Therefore flare accelerated electrons one-dimensionally propagating along magnetic field lines can trace individual flux tubes from the electron acceleration site in the corona down to the deep layers of chromosphere where electrons emit hard X-ray emission. The asymmetry of hard X-ray footpoint sizes of a flaring loop \\citep{melrose1979} could be a measure of the ratio of magnetic field strengths in the X-ray loop footpoints \\citep{schmahl2007}. The simple dependency of X-ray emission maximum location on the photon energy and density structure in the low corona/chromosphere due to Coulomb collisions \\citep{brown2002}, has allowed \\citet{aschwanden2002} to infer the chromospheric density structure from Reuven Ramaty High Energy Solar Spectroscopic Imager RHESSI \\citep{lin2002} X-ray observations.  In this letter, we show that X-rays can be a diagnostic tool for the analysis of not only energetic electrons in solar flares but of the magnetic flux tubes and density structure in the chromosphere. We analyse the spatial and energy distribution of hard X-ray sources using RHESSI data to infer chromospheric density and magnetic field structure. Our results show that the density distribution of neutral hydrogen in the flaring loop has a scale height of 140~km and the magnetic flux tube of the flaring loop fans out by a factor of $\\sim 3$ at the height of around $900$ km. ",
        "conclusions": "Forward fitting X-ray visibilities allow simple and reliable measurements of not only locations of the emission maxima but also the characteristic sizes of hard X-ray sources. By using this visibility analysis on a good candidate limb flare we have improved earlier chromospheric height and density measurements of \\citet{aschwanden2002}, reducing the uncertainty of emission maximum positions down to $\\sim0.2''$. Assuming collisional transport in neutral hydrogen it can be concluded that the chromospheric density is consistent with a gravitationally stratified atmosphere of density scale height $140\\pm 30$~km. We show for the first time that not only is the higher energy X-ray emission produced continuously deeper in the chromosphere but the X-ray source sizes decrease from $\\sim6.2''$ down to $\\sim2.3''$. The precise measurement of the characteristic sizes allows us to conclude  that magnetic field directing the energetic electrons converge, with the magnetic flux tube shrinking from FWHM $\\sim3.5$~Mm at $h\\sim 1$~Mm down to FWHM $\\sim2.5$~Mm at $h\\sim 0.8$~Mm above the photosphere. The magnetic scale height estimated using flux conservation $B(h)\\sim FWHM(h)^{-2}$ as the ratio of areas $B(h)\\left(dB/dh\\right)^{-1}= -FWHM(h)\\mbox{d}h/\\mbox{d}FWHM(h)/2$ is found to be of the order of 300~km. This confirms the fanning out (canopies) of the chromospheric magnetic fields \\citep{gabriel1976}.  Our deduced density structure agrees quite well with independent estimations from other methods, possibly a surprising conclusion in view of the simple treatment of electron transport, completely neglecting any pitch-angle changes. As already noted, collisional scattering will modify the range only by a factor of order unity but magnetic moment conservation might have a greater effect, unless fast electrons all start with velocity vectors parallel to $\\vec{B}$. Large pitch angles would develop via collisional scattering only as electrons reached the end of their ranges, and the mirror force would be similarly unimportant most of the time. Such a concentration at small pitch angles seems at odds with findings of a nearly isotropic electron distribution from studies of photospheric albedo \\citep{KontarBrown06mirror}. The magnetic field convergence we find here however offers a simple resolution. The flux tube implied by our HXR source FWHM's implies magnetic field lines substantially inclined to the vertical, possibly by as much as 60\\% on average at the outer edge and in addition are likely to be twisted. The HXR polar diagram of electrons populating the whole of an ensemble of field lines, with such a range of angles to the vertical, would be much closer to isotropic than expected for vertical field lines. The presence of magnetic canopies, then, appears to be a critical factor for interpreting HXR images and even spectra, but one that needs a much more substantial and detailed modelling effort than we attempt here.  While the X-ray diagnostics of the chromospheric magnetic field is an extremely attractive new technique, we emphasise that the morphology of X-ray sources should be further scrutinised. We note that the shape of the hard X-ray sources is more elliptical (Figure 3) rather than circular suggesting that the vertical extend of X-ray sources is govern by the spatially varying magnetic field (convergence and a twist of the flux tube), through the conservation of the electron magnetic moment. This complicates the interpretation of the HXR source sizes beyond the simple ideas embodied in (\\ref{F_ez}) and (\\ref{I_ez}) but brings a new diagnostic potential. This requires a more complicated electron dynamic and forward fitting elliptical sources to the X-ray visibilities, which will be the subject of future work.  The method of measuring magnetic field using X-rays could be viewed as unique tool of loop width measurements in the chromosphere comparable to using TRACE data to measure the widths of flux tubes in the corona \\citep{watko2000}. While it is often argued that the heights of line formation should not be assigned due to the fact that very distinct layers of the atmosphere could be sampled \\citep[e.g.][]{sanchez1996}, X-rays are uniquely related to the magnetic field lines connecting the photosphere and the electron injection site in the corona and weakly sensitive (via Coulomb logarithm dependency) to the temperature variations in the chromosphere. Therefore hard X-ray emission is likely to be a valuable diagnostic for mapping chromospheric magnetic field and density structures."
    },
    "0808/0808.0930_arXiv.txt": {
        "abstract": "Several lines of evidence have suggested that the the galaxy cluster Cl~0024+17, an apparently relaxed system, is actually a collision of two clusters, the interaction occurring along our line of sight. In this paper we present a high-resolution $N$-body/hydrodynamics simulation of such a collision. We have created mock X-ray observations of our simulated system using MARX, a program that simulates the on-orbit performance of the Chandra X-ray Observatory. We analyze these simulated data to generate radial profiles of the surface brightness and temperature. At later times, $t = 2.0-3.0$ Gyr after the collision, the simulated surface brightness profiles are better fit by a superposition of two $\\beta$-model profiles than a single profile, in agreement with the observations of Cl~0024+17. In general, due to projection effects, much of the post-collision density and temperature structure of the clusters is not seen in the observations. In particular, the observed temperatures from spectral fitting are much lower than the temperature of the hottest gas. We determine from our fitted profiles that if the system is modeled as a single cluster, the hydrostatic mass estimate is a factor $\\sim$2-3 less than the actual mass, but if the system is modeled as two galaxy clusters in superposition, a hydrostatic mass estimation can be made which is accurate to within $\\sim$10$\\%$. We examine some implications of these results for galaxy cluster X-ray surveys. ",
        "introduction": "Clusters of galaxies have proven to be interesting laboratories for studying the dynamics of the different kinds of matter in the universe. Early measurements of galaxy velocity dispersions in clusters demonstrated that most of the mass of galaxy clusters was in the form of a non-luminous component \\citep{zwy37}. Modern cosmological constraints (e.g. Turner 2000) suggest that this non-luminous material must be mostly non-baryonic. In addition, X-ray observations of clusters of galaxies reveal that most of their baryonic mass is in the form of a hot, diffuse, X-ray emitting gas (the intracluster medium or ICM). The presence of these different types of matter (stellar, gaseous, and dark) in these systems provides not only insights into the nature of the clusters themselves but also the nature of the different components of matter and their dynamics.  In the cold dark matter (CDM) picture of structure formation, galaxy clusters are the product of many mergers between galaxies, galaxy groups, and smaller clusters of galaxies \\citep{dav85}. Observations of galaxy clusters indicate that this process is still ongoing in many systems. A famous example is the so-called ``Bullet Cluster'' (1E~0657-56), a merging system in which X-ray and weak gravitational lensing observations demonstrate a clear separation between a collisionless dark matter component and the intracluster medium \\citep{mar02, clo06}.  Another system that has recently attracted attention is Cl~0024+17, a cluster at a redshift $z = 0.395$ exhibiting both weak and strong lensing. Early attempts at reconstructing the mass profile of this system using lensing \\citep{tys98, bro00, com06} suggested a conflict with the predictions of the standard CDM model due to the flattening of the density profile in the inner regions of the cluster. Cosmological simulations assuming CDM indicate that galaxy cluster mass profiles should exhibit a nonzero logarithmic slope in the inner regions (e.g. Navarro, Frenk, \\& White 1997). This apparent discrepancy (and others) led to suggestions that the CDM paradigm would need to be modified \\citep{spr00,hog00,moo00}, for example to include self-interaction of the dark matter.  \\citet{czo01, czo02} demonstrated that the redshift distribution of the cluster galaxies in Cl~0024+17 is bimodal, with a large, primary component and a more diffuse, secondary component, separated from the primary component in velocity space by $v \\sim$ 3000 km/s. They suggested that the system is composed of two clusters undergoing a high-velocity collision along the line of sight. They also performed a simulation demonstrating that such a scenario reproduces not only the bimodal redshift distribution but also the flattening of the central density profile. X-ray observations of the cluster using Chandra \\citep{ota04} and XMM-Newton \\citep{zhg05} revealed that the surface brightness profile is better fit by a superposition of two ICM models rather than one. They suggested on the basis of the isothermal temperature profile that the individual clusters of the system had returned to equilibrium after the collision and that consequently the encounter must have occurred several Gyr ago.  Recently, a weak lensing analysis presented by \\citet{jee07} revealed a ringlike structure in the projected matter distribution. They proposed that dark matter from the cores of the clusters had been disrupted and ejected from the systems by the collision, and they demonstrated that such features could be reproduced in a simulation of a collision of two pure dark matter halos. From this they suggested the current state of the system is $\\sim$1-2 Gyr past the pericentric passage of the cluster cores. We are addressing this issue in a separate paper (ZuHone et al. 2008, in preparation). \\citet{jee07} also demonstrated that if the system is modeled as two ICM profiles in superposition, the mass determination based on hydrostatic equilibrium agrees with the result from their lensing analysis.  If the merger scenario for Cl~0024+17 is correct, it raises several important questions: how long after the collision would the clusters appear relaxed, and correspondingly when would a mass estimate based on hydrostatic equilibrium yield an accurate measurement? What effect does viewing such a collision along the line of sight, with both cluster components in superposition, have on the observed density and temperature structure? In this paper we seek to provide answers to these questions by simulating a similar high-speed collision between two clusters of galaxies. Most importantly, we include the dynamics of the cluster gas, which were not included in the previous simulations.  Throughout this paper we assume a flat CDM cosmology with $h = 0.5$, $\\Omega_{\\rm m} = 1.0$, as in \\citet{ota04}. ",
        "conclusions": "The galaxy cluster Cl~0024+17 is thought to be a colliding system of galaxy clusters that is viewed along the collision axis. We have performed a simulation of a high-velocity head-on collision of galaxy clusters and created mock X-ray observations, viewing the clusters from the same direction. Previous investigations of this system had focused on the collisionless dynamics of the galaxies and the dark matter; we also include gas in our simulation. The mock X-ray observations of our simulation indicate that the X-ray emitting gas is still undergoing moderate evolution at times $\\sim$1-2 Gyr after the collision. However, much of the complicated structure in density and temperature is obscured due to projection effects. The X-ray surface brightness profile is better fit by a superposition of $\\beta$-models than a single $\\beta$-model at times $t =$ 1-2 Gyr after the collision, corresponding to the observed situation in Cl~0024+17. The cluster gas in the center appears colder than it really is due to the projection of denser, colder gas along the line of sight of the hotter gas. The temperature profile exhibits only moderate evolution with time, though the temperature distribution of the clusters in the simulation has significant structure. Though a mass estimate of the cluster assuming one cluster component in hydrostatic equilbrium underestimates the mass by a factor of $\\sim$2-3, a mass estimate based on the assumption of two galaxy clusters in superposition comes close ($\\sim$10$\\%$) to the actual projected mass of the system within the arc radius $r_{\\rm arc}$, in agreement with the mass measurements made of the cluster Cl~0024+17 under the same assumptions. These results may be valuable when looking at clusters discovered in X-ray surveys. With current photometric redshift errors it will not likely be possible to distinguish such line-of-sight collisions from single clusters discovered in X-ray surveys at high redshift by identifying separate galaxy concentrations. However, our results show that it may be possible to distinguish merging clusters and clusters in projection from single clusters by testing multiple model fits against single model fits, though this will become more difficult at higher redshifts."
    },
    "0808/0808.3980_arXiv.txt": {
        "abstract": "Using the 2004 Venus transit of the Sun to constrain a semi-empirical point-spread function for the \\emph{TRACE} EUV solar telescope, we have measured the effect of stray light in that telescope. We find that 43\\% of 171\\AA\\ EUV light that enters \\emph{TRACE }is scattered, either through diffraction off the entrance filter grid or through other nonspecular effects. We carry this result forward, via known-PSF deconvolution of \\emph{TRACE }images, to identify its effect on analysis of \\emph{TRACE }data. Known-PSF deconvolution by this derived PSF greatly reduces the effect of visible haze in the \\emph{TRACE }171 \\AA\\ images, enhances bright features, and reveals that the smooth background component of the corona is considerably less bright (and hence more rarefied) than might otherwise be supposed. Deconvolution reveals that some prior conclusions about the Sun appear to have been based on stray light in the images. In particular, the diffuse background {}``quiet corona'' becomes consistent with hydrostatic support of the coronal plasma; feature contrast is greatly increased, possibly affecting derived parameters such as the form of the coronal heating function; and essentially all existing differential emission measure studies of small features appear to be affected by contamination from nearby features. We speculate on further implications of stray light for interpretation of EUV images from \\emph{TRACE }and similar instruments, and advocate deconvolution as a standard tool for image analysis with future instruments such as SDO/AIA. ",
        "introduction": "All telescopes, including \\emph{TRACE } \\citep{Handy1998}, scatter light. The principal scattering mechanisms in a space-based telescope include diffraction through the aperture and any obscuration in the beam of the telescope, irregularities or dust on the mirrors themselves, and reflection or scattering in the detector at the focal plane. All these effects contribute to forming broad, shallow wings on the point-spread function (PSF) of the instrument, which describes the image produced by that instrument when viewing an ideal point source of light. These wings are typically 3-5 orders of magnitude fainter than the core of the PSF, but 3-4 orders of magnitude larger, so that a significant fraction of the light incident into the telescope is spread over a large portion of the image.  For discrete scenes such as starfields, broad PSF wings are not greatly important except that they reduce the net efficiency of the telescope by reducing the apparent brightness of point sources. For continuous or near-continuous, high-contrast scenes such as viewed by solar telescopes, broad scattering wings are quite important: scattering from large distributed bright structures can overwhelm the emission from dark regions in the same image. These effects, while well known, have received little attention from the solar data analysis community when applied to data from normal-incidence EUV telescopes such as \\emph{TRACE, }but they are present nonetheless and must be accounted for when interpreting images from \\emph{TRACE }(and all other solar telescopes). Not only quantitative analysis, but even some qualitative interpretations of \\emph{TRACE }data may be compromised if stray light effects are not accounted for.  X rays and ultraviolet light are particularly susceptible to scattering and defocus, and efforts have been made to account for point-spread function effects in previous instruments. For example, \\cite{Maute1981} attempted blind iterative deconvolution on the X-ray data from \\emph{Skylab,} \\cite{svestka1983} applied it to the \\emph{SMM/}HXIS instrument; and \\cite{Martens1995} determined a spatially variable PSF for the \\emph{Yohkoh}/SXT instrument. These studies have largely focused on iterative methods to identify and remove blurring effects caused by a broad PSF core in the subject instruments, though stray light has also been an object of study. Stray light deconvolution was commonly used on SXT data in the later years of that mission (e.g. \\citet{Foley1997,Gburek2002,Schrijver2004}).  The \\emph{TRACE} EUV PSF has been studied by several groups. \\cite{Lin2001} used compact, bright flares to study diffraction patterns on the \\emph{TRACE} focal plane and concluded that diffraction from the aluminum filter grids used in \\emph{TRACE} scatters 19\\% of incident EUV photons into a highly structured, broad diffraction pattern; they speculated that the scattering may be affecting imaging performance. \\cite{gburek2006} used that scattering pattern both to derive both a best-fit PSF core for \\emph{TRACE} and also to determine a portion of the emitted EUV spectrum from particular flare events.  In this report, we consider primarily the diffuse scattering wings of the \\emph{TRACE} EUV PSF, and particularly their implications for interpreting coronal images. The wings are not readily measured using a point source such as a flare, because the local intensity of the PSF is quite small far from the core. Outside of diffraction maxima the weak scattered signal is overwhelmed by local emission even for bright events such as flares. Deriving the PSF thus requires analysis of occulted images, using an obstructing body such as the Moon or a planet. We examined \\emph{TRACE }data collected near the times of several solar eclipses, but did not find a suitable EUV image set that contained a clear image of the lunar limb. On 2004 Jun 08, Venus passed in front of the Sun, and several 171 \\AA\\ image sequences were collected as the planet traversed the disk of the Sun and the off-limb corona. We have used those images to derive a semi-empirical scattering PSF for the \\emph{TRACE} 171 \\AA\\ channel, and have tested the PSF for correctness by using it to deconvolve several representative\\emph{ }images of interesting coronal structures.  In \\S\\ref{sec:Review-of-Deconvolution} we briefly review deconvolution and how it is performed, in \\S\\ref{sec:Constraint-of-the} we describe the forward modeling process and present our measured PSF, and in \\S\\ref{sec:Deconvolution-of-sample} we demonstrate deconvolution of some representative images. Finally, in \\S\\ref{sec:Discussion} we discuss implications for interpretation of EUV coronal images and recommend deconvolution as a standard reduction pipeline component for future telescopes. ",
        "conclusions": "Discussion}  The result that \\emph{TRACE} images contain a significant amount of scattered light is not, in itself, new. Most telescope PSFs include scattering wings, and \\emph{TRACE }is no exception. The filter grid in the front of the telescope is known to scatter $\\sim20\\%$ of incident light into a structured diffraction pattern with myriad local maxima (\\cite{Lin2001,gburek2006}). The present analysis is new in three important respects: it is (to our knowledge) the first analysis of scattered light using \\emph{TRACE }occultation data to derive a PSF; much more scattered light is found than can be accounted for merely by diffraction; and the process is taken to its natural conclusion of deconvolving the original \\emph{TRACE }images to show the effect of the stray light on scientific interpretation of the images.  The \\emph{TRACE }Venus data are not detailed enough to constrain a highly structured PSF model, but the simple empirical fit described here is sufficient to improve existing images via deconvolution, and passes the most basic of deconvolution tests, suggesting that it is not overcompensating for the scattering wings. Deconvolution with our scattering PSF has a similar effect to background subtraction on the interpretation of small features: for features small compared to the scattering wings, the effect of the surrounding bright features is approximately constant, so subtraction of a modeled or fixed background yields similar effects in particular local areas of a given \\emph{TRACE }image. The principal advantages of deconvolution are that an approximation of absolute brightness is reproduced, rather than the offset relative brightness that may be extracted from simple background-subtracted images; and that moderate-scale features are treated correctly (they are not treated correctly by simple background or pedestal subtraction).  Deconvolution not only darkens the faintest portions of the image, it both increases the relative contrast of small bright features embedded in a bright background and also affects photometric estimates of the relative density of any small bright feature seen with \\emph{TRACE}. This describes several features of interest in the \\emph{TRACE} data, including active-region threads that are a subject of current debate \\citep{WatkoKlimchuk2000,WarrenWinebarger2003,Fuentes2006,DeForest2007}.  The greater coronal contrast we find in deconvolved \\emph{TRACE} images gives indirect support to the idea that the corona is close to hydrostatic equilibrium despite the observed tallness of bright features such as active region loops. The coronal density scale height is about 50 Mm at 1 MK, so the emissivity scale height of Fe IX \\& Fe X emission line features (close to 1MK ionization temperature) is about 25 Mm (0.035 $R_{S}$) assuming local thermal equilibrium. Thus the EUV-visible corona might be expected to form a thin layer near the photosphere with no significant emission arising at altitudes higher than about 0.07 $R_{S}$. Essentially all \\emph{TRACE }EUV images show significant background brightness high in the corona; the brightness is visible above the detector {}``pedestal'', because there is contrast between dark but {}``live'' pixels that are part of the image, and corner pixels that are vignetted by the round filter mount at the back of the instrument. The current measured instrument PSF suggests that most or all of this background brightness is due to scattering within the telescope, because dark features (such as the prominences in the bottom row of Figure \\ref{fig:deconvolution-demo}) are reduced nearly to zero brightness when deconvolved. This is the general behavior to be expected from a thin hydrostatic atmosphere at a particular temperature: tall features that are more than 1-2 scale heights tall should have little or no emission above them.  Active region loops appear to have a large scale height compared to that expected for 1-2 MK plasma \\citep{SchrijverMcMullen2000,AschwandenNitta2000,WinebargerWarrenMariska2003,Fuentes2006,DeForest2007}. Three explanations that have been advanced are resonant scattering of EUV (which varies as $n_{e}$ rather than $n_{e}^{2}$), support by non-hydrostatic momentum transport mechanisms such as siphon flows or wave motion, or geometric considerations that attenuate brightness at the bases of the loop. Our result that the quiet corona appears to be consistent with the expected hydrostatic scale height seems to eliminate resonant scattering as a mechanism for tallness, because it would imply a stronger haze in the foreground at high solar altitudes. Further, it seems to limit the functional form of anomalous support mechanisms that could lengthen active region loops' scale height, because such mechanisms must act preferentially against active region loops and not quiet sun loops, to be consistent with the morphology of the deconvolved images. This can further be construed as circumstantial evidence for a geometric, rather than intrinsic, explanation for active region loops' long apparent length \\citep{DeForest2007}.  \\cite{Moore2008} have recently used image-processing techniques to separate the hazy and sharp components of active region loops viewed with \\emph{TRACE. }Such analyses rely on the sharp component of the corona as an indicator of stray light. With 43\\% scattering of the total light incident on the telescope, 57\\% is left to be focused; hence, we expect that the hazy portion of such separated image pairs derived from \\emph{TRACE }171\\AA\\  data should contain about 75\\% as much total brightness as the sharp portion does, on the basis of stray light alone.  In addition to morphological differences, corrections to the relative brightness of features such as active region threads and voids affect parameters such as the derived Alfv\\'en speed, because of the $n_{e}^{2}$ dependence of EUV emission. Structures with spiky density profiles emit more EUV per electron than do smooth structures, and the inferred Alfv\\'en speed depends both on the magnetic field and the derived electron density. Onset of some \\emph{TRACE-}observed EIT waves appears to require high Alfv\\'en speeds of up to $3\\, Mm\\, s^{-1}$ \\citep{willsdavey2007}; this high speed is difficult to explain in the presence of a diffuse background corona around the source region of the EIT wave. If in fact active regions contain nearly evacuated regions (as in the center panel of Figure \\ref{fig:deconvolution-demo}), then the variation in Alfv\\'en speed is greatly increased and the region-wide average Alfv\\'en speed may be significantly higher than would otherwise be inferred.  As a final example of the impact of stray light in EUV images\\emph{, }coronal heating properties have been derived \\citep{Schrijver2004} by examining the contrast between coronal holes and bright structures, and may be affected by scattering in \\emph{SOHO}/EIT and/or \\emph{Yohkoh}/SXT. Specifically, if coronal holes are significantly darker, and bright structures are significantly sharper and brigher, than is apparent in raw EUV and X-ray images, then the coronal heating mechanism may not be as distributed as might otherwise be inferred.  The model PSF that we have derived for \\emph{TRACE} is somewhat simplified: we have included, \\emph{a priori, }detailed structure that is known from earlier studies \\citep{Lin2001}, and parameterized an additional scattering term based on the empirical behavior of stray light on rough mirrors. We have not taken into account possible anisotropy or spatial variability of the scattering, attempted to gain physical understanding of the causes of the PSF, or modeled scattering phenomena that do not fit within the paradigm of a simple PSF. Based on measurements of the Venus transit in 2004, we have found that roughly 43\\% of incident energy is scattered by \\emph{TRACE}, so that approximately half of the scattered energy may be ascribed to the diffraction pattern found by Lin et al. and approximately half to other mechanisms. Deconvolution greatly improves contrast in \\emph{TRACE }images, raising concerns about the interpretation of those images.  More generally, deconvolution to increase contrast in images with scattering wings is strongly recommended for observation from present and future EUV and X-ray telescopes. We have shown that deconvolution can greatly affect the contrast of observed features, and discussed how this effect may affect a broad variety of science questions. Further, deconvolution to remove broad scattering wings is in general not as hazardous to the data as is deconvolution to increase sharpness in the core of the telescope PSF. That is because there are high spatial frequencies present in the core of the kernel, even if it is added to a much broader distribution, so that noise is not increased as much from deconvolution of a scattering PSF as from a broad PSF core."
    },
    "0808/0808.4044_arXiv.txt": {
        "abstract": "We compare the detection rates and redshift distributions of low-luminosity (LL) GRBs localized by {\\it Swift} with those expected to be observed by the new generation satellite detectors on {\\it GLAST} (now {\\it Fermi}) and, in future, {\\it EXIST}. Although the {\\it GLAST} burst telescope will be less sensitive than {\\it Swift}'s in the 15--150 keV band, its large field-of-view implies that it will double {\\it Swift}'s detection rate of LL bursts. We show that {\\it Swift, GLAST} and  {\\it EXIST} should detect about 1, 2 \\& 30 LL GRBs, respectively, over a 5-year operational period. The burst telescope on {\\it EXIST} should detect LL GRBs at a rate of more than an order of magnitude greater than that of {\\it Swift}'s BAT. We show that the detection horizon for LL GRBs will be extended from $z \\simeq 0.4$ for {\\it Swift} to $z \\simeq 1.1$ in the {\\it EXIST} era. Also, the contribution of LL bursts to the observed GRB redshift distribution will contribute to an identifiable feature in the distribution at  $z \\simeq 1$. ",
        "introduction": "GRBs are the brightest transient astronomical events known. Multi-wavelength observations of GRBs have confirmed that a significant fraction of long GRBs\\footnote {Hereafter, `GRB' refers to bursts classified as long.} are associated with the collapse of short-lived massive stars \\citep{Hjorth2003, Stanek2003}.  A number of satellites equipped for GRB detection have been placed in low Earth orbit, the most recent of these being the orbiting observatories {\\it Swift}$\\,$\\footnote{http://swift.gsfc.nasa.gov} (launched 2004 November; \\citealt{Gehrels2007}) and {\\it GLAST}$\\,$\\footnote{http://glast.gsfc.nasa.gov} (launched 2008 June). Up to the end of 2008 July {\\it Swift} has detected 345 bursts\\footnote{http://swift.gsfc.nasa.gov/docs/swift/archive/grb\\_table.html}  since its launch. In this period 116 GRB redshifts have been measured from this satellite's rapid localization capability and by follow-up spectroscopy of the optical/NIR afterglow and host galaxies by large ground-based telescopes.  Prior to GRB 980425 it was thought that all GRBs were from a limited range of energies \\citep{Ramirez-Ruiz2005}. This particular GRB was the first to be associated with a supernova of Type Ib/c and was also the first low-luminosity (i.e., under-luminous) GRB (LL GRB) to be identified \\citep{Ramirez-Ruiz2005}. The total (isotropic-equivalent) gamma-ray energy emission was some 5--6 orders of magnitude less than for a typical `classical' GRB (fig. 2 of \\citealt{Ghirlanda2006}), demonstrating that GRBs can occur with a wide range of energies \\citep{Woosley2004}. Initially it was thought that this GRB was anomalous but more recent detections have challenged that view. GRBs 031203 and 060218 have been identified as possible analogues of GRB 980425 (\\citealt{Soderberg2004, Ghisellini2006, Watson2006}) and these three can be considered to form a `typical' set of LL GRBs. For a brief description of these three bursts see \\citet{Guetta2007}.   Because of their sub-energetic nature and the flux limits of recent satellite detectors, the LL GRBs detected thus far are in the local universe ($z  \\la  0.1$). Apart from the three typical LL GRBs mentioned above, other subluminous GRBs have been detected --- e.g., GRB 050826 \\citep{Mirabal2007} --- some of which have been classified as X-ray flashes --- e.g., XRFs 020903, 030723  and 060218 \\citep{Pian2006} --- or X-ray rich --- e.g., GRB 040223 \\citep{McGlynn2005}.  X-ray flashes may represent an extension of the GRB population to low energies \\citep{Barraud2005}.  Although LL GRB events may be more numerous than `classical' events \\citep{Pian2006, Guetta2007, Liang2007} only a few have been observed in the local universe \\citep{Soderberg2004, Guetta2007, Daigne2007}. Small-number statistics limit the precision in estimating the local rate density of LL GRBs --- see uncertainties for values of this parameter reported by \\citet{Guetta2007} and \\citet{Liang2007}.   We show in this paper that the improved sensitivity and larger field-of-view (FoV) of these new generation satellite detectors will dramatically increase the statistics of these currently rare bursts. We also show how the improved capabilities of these detectors allows detection of LL GRBs at higher $z$. (For a perspective on studies of  high-$z$ GRBs with new generation instruments see \\citet{Salvaterra2008}.)  The aim of this study is to compare the detection capabilities of three gamma-ray detector satellites for observing LL GRBs, focusing on GRB detection rates of the primary burst detectors of {\\it Swift}, {\\it GLAST} and {\\it EXIST}$\\,$\\footnote{http://exist.gsfc.nasa.gov} (intended to be launched about 2015). We utilize only the primary detector parameters of flux sensitivity, FoV and detector bandwidth for this study. For {\\it GLAST} we assume the Gamma-ray Burst Monitor (GBM) is the detector that will trigger on and monitor LL GRBs \\citep{Steinle2006}. Similarly, for {\\it EXIST} we assume the High Energy Telescope (HET) is the primary LL GRB detector. Thus, our results for {\\it GLAST} and {\\it EXIST} define a lower bound on the detection capabilities of these new satellite detectors. ",
        "conclusions": "\\label{sec:ResultsDiscussion} \\begin{figure} \\includegraphics[scale=0.40]{Fig1.eps} \\caption{The LL GRB detection probability, $\\psi_{\\rm eff}$, for the 3 satellite detectors. The curves show the relative detection efficiencies for detecting LL GRBs. The functions are determined by $F_{\\rm lim}$, the k-correction, bandwidth scaling, FoV and LF. As the LF is identical for all detectors, differences between the curves result from instrument parameters. At $z = 0$, the probability of detecting a single isotropic LL GRB is just the detector's FoV (as a sky-fraction). The detection horizon distance for the detector, $z_{\\rm h}$, is defined by $\\psi_{\\rm eff}(z_{\\rm h}) = 0$. The detection horizon is $z \\le 1.1$ for all three detectors. The high sensitivity of HET results in a detection horizon which is about 3--4 times that of the other two detectors.} \\label{fig:DetectorScFns} \\end{figure}  We have not assumed that the $E_{\\rm p,i}$--$E_{\\rm iso}$ correlation is valid for LL GRBs. While GRB 060218 obeys this relation the other LL GRBs (980425 and 031203) do not (fig. 3 of \\citet{Amati2006}). Although \\citet{Ghisellini2006} have indicated that the two outliers may indeed obey the Amati relation if spectral evolution is taken into account there appears to be no compelling physical evidence to support this at present. Given the above, we take a conservative approach, retaining the Band spectrum and setting $E_{\\rm peak} = 100$ keV for LL GRBs (which is between the observed values for GRB 980425 and 031203 -- see fig. 3 or tables 1 \\& 2 of \\citet{Amati2006}). Calculations made with $E_{\\rm peak} = 10$ keV reduce the horizon limits (see next paragraph) to (85, 70, 90)\\% those for $E_{\\rm peak} = 100$ keV. Similarly, detection rates are (60, 35, 75)\\% those of the higher $E_{\\rm peak}$. However, these do not change the qualitative results we report here. With greater numbers of LL GRBs detected in future the observed $E_{\\rm peak}$ may be used to confirm whether LL GRBs are consistent with the Amati relation or not.  The calculated detection probabilities for the three satellites for LL GRBs are shown in Fig. \\ref{fig:DetectorScFns}. We define the `horizon distance', $z_{\\mathrm h}$, of a detector to be the greatest distance at which a LL GRB can trigger the detector ($z_{\\mathrm h} \\propto \\sqrt{F_{\\rm lim}}$).  We convolve the instrument effectiveness above with the GRB beaming factor and GRB rate evolution, $e(z)$, to produce detection rate curves (i.e., $dR_{\\rm obs}(z)/dz$) for each detector and GRB population. The predicted distributions of LL and HL GRB detections are plotted in Fig. \\ref{fig:PredObsRates_LLandHL}. The combined GRB detection  rate curves are shown in Fig. \\ref{fig:PredObsRates_Combined} for each detector. We note that the LL GRB distribution produces a noticeable `notch' at $z \\approx 1$ for HET, a result of its higher sensitivity.  \\begin{figure} \\includegraphics[scale=0.40]{Fig2a.eps} \\includegraphics[scale=0.40]{Fig2b.eps} \\caption{Predicted differential observation rates for BAT, GBM and HET for LL (upper panel) and HL GRBs (lower panel). The GRB rate density peak values for GRBs  are determined by the sensitivity and FoV of the detectors while their $z$-locations are determined by sensitivity (and, to a lesser extent, redshift).} \\label{fig:PredObsRates_LLandHL} \\end{figure}  \\begin{figure} \\includegraphics[scale=0.40]{Fig3.eps} \\caption{Predicted differential detection rates for the GRB population comprised of LL \\& HL bursts. These curves are the sum of the LL \\& HL detection rate models (Fig. \\ref{fig:PredObsRates_LLandHL}). The effect of two different beaming factors results in a LL GRB `notch' at ${z \\approx 1}$ for HET. Because of the lower sensitivities of the GBM \\& BAT this feature is not visible to those detectors.} \\label{fig:PredObsRates_Combined} \\end{figure}  \\begin{table} \\label{tbl:results} \\caption{Predicted LL GRB detection rates and horizons.} \\begin{tabular}{lcc} \\hline {\\bf ~Detector}               &   {\\bf LL GRB}          &    {\\bf LL Horizon}                    \\\\ {\\bf (Satellite)}                &   {\\bf [yr$^{-1}$]}     &    {\\bf [z$_{\\mathrm h}$] }        \\\\\\hline BAT ({\\it Swift})              &     0.2                         &      0.4                                         \\\\ GBM (${\\it GLAST}$)    &     0.4                        &       0.3                                         \\\\ HET (${\\it EXIST}$)       &     6.5                        &      1.1                                          \\\\ \\hline \\end{tabular} \\end{table}   The locations of the LL GRB peaks in redshift (upper panel of Fig.~\\ref{fig:PredObsRates_LLandHL}) is determined by the relative sensitivities (and, to a lesser extent, trigger bandwidths) of the detectors. A more sensitive detector will be triggered by a greater number of distant GRBs. The peak LL detection rates of BAT and GBM occur in the redshift regime of $[0.15, 0.3]$ while for HET they occur at the significantly higher redshift of just over 0.7. The LL GRB detection horizon distances of the above detectors are $z = (0.4, 0.3, 1.1)$. The sensitivity of HET will also enable it to detect HL GRBs at higher $z$, thus probing a larger volume of space, as shown by the lower panel of Fig.~\\ref{fig:PredObsRates_LLandHL}.  The height of each curve in Fig. \\ref{fig:PredObsRates_LLandHL} is proportional to the FoV of the detector. For a bright GRB population the FoV is the dominant parameter, but sensitivity becomes important for detecting fainter bursts, as illustrated by comparing the GBM and HET curves for both populations in Fig.~\\ref{fig:PredObsRates_LLandHL} (upper and lower panels). The moderately large FoV of HET coupled with its relatively high sensitivity implies it will detect an order of magnitude more LL GRBs than the other two satellites combined (Table~3). The large FoV of GBM enables it to capture the most bursts for $z \\la 0.3$ even though it is the least sensitive detector in this comparison study. Its large FoV also makes it the optimal detector out of the three for HL GRBs out to redshift about 4.5. Fig. \\ref{fig:PredObsRates_Combined} indicates that HET will also be able to detect GRBs in the early universe, thus potentially probing the evolutionary history of GRBs.  From the upper panel of Fig.~\\ref{fig:PredObsRates_LLandHL} we note that HET will detect large numbers of nearby LL GRBs and, thus, significantly add to the statistics of this under-luminous population. In addition, Fig.~\\ref{fig:DetectorScFns} shows that HET takes over from GBM as the most efficient LL GRB detector for $z \\ga 0.3$ in spite of it having a quarter of the GBM's FoV. This is due to HET's higher sensitivity, which is also reflected in HET having a much larger LL horizon distance than the other two detectors.  {\\it GLAST}$\\,$'s GBM, having the largest FoV, produces thebroad and high differential detection rate prominent in the lower panel of Fig.~\\ref{fig:PredObsRates_LLandHL}. The predicted $z$-integrated detection rates of LL GRBs are (0.2, 0.4, 6.5) yr$^{-1}$ for BAT ({\\it Swift}), GBM ({\\it GLAST}) and HET ({\\it EXIST}), respectively.  The notch in Fig.~\\ref{fig:PredObsRates_Combined} for the HET curve is a result of two factors: (1)  HET's high sensitivity will detect a higher proportion of LL GRBs than will the other two satellites and (2) the low LL GRB beaming factor and high local rate density (compared to HL GRBs) results in an effective differential LL rate about 250 times that of HL GRBs. However, it may take two to three decades for this feature to become apparent. Hence, it is unlikely that any single satellite observatory will detect this `hidden' LL peak during an operational career. Nonetheless, even the 3 or so LL GRBs detectable by {\\it Swift} and {\\it GLAST} in the next 5 years will be a significant addition and will provide constraints on this population.  We note two selection effects that may be significant but are not included in this study. Firstly, as \\citet{Band2006} showed, the BAT is more sensitive to high-redshift bursts because of their longer duration. Secondly, the rate density of GRBs may be evolving at high redshift, as suggested by \\citet{Daigne2006} and \\citet{Le2007}. \\citet{Firmani2004} reach a similar conclusion, but explain a higher than expected rate as evidence for an evolving GRB luminosity function. These issues remain uncertain. However, their effect does not influence the main results from this study since our focus is on the detectability of small-redshift LL bursts.  We also note that selection effects that plague the present GRB redshift distribution need to be understood to exploit fully data for studying GRB rate evolution. \\citet{Coward2008} show that selection effects are especially dominant in the optical afterglow and spectroscopic observations of {\\it Swift}-triggered bursts.  In summary, we show that HET will probe both LL GRBs and early-universe GRBs while GBM will be more sensitive to the HL population out to $z \\approx 4.5$. Surprisingly, we find that despite the lower sensitivity of GBM, its large FoV implies a LL GRB detection rate greater than that of BAT and HET out to the GBM's sensitivity horizon of $z = 0.3\\;.$ Importantly, HET will increase the LL GRB detection rate by a factor of about 30 compared to {\\it Swift}'s BAT. It will open up a new detection regime by probing the LL GRB population in a significantly larger volume out to $z \\sim 1$."
    },
    "0808/0808.3722_arXiv.txt": {
        "abstract": "{} {To understand the environment and extended structure of the host galactic gas whose molecular absorption line chemistry, we previously observed along the microscopic line of sight to the blazar/radiocontinuum source NRAO150 ($aka$ B0355+508).} {We used the IRAM 30m Telescope and Plateau de Bure Interferometer to make two series of images of the host gas: i) 22.5\\arcsec\\ resolution single-dish maps of \\twCO\\ J=1-0 and 2-1 emission over a 220\\arcsec by 220\\arcsec\\ field; ii) a hybrid (interferometer+singledish) aperture synthesis mosaic of \\twCO\\ J=1-0 emission at 5.8\\arcsec\\ resolution over a 90\\arcsec-diameter region.\\thanks{The three spectra cubes (RA, Dec, Velocity) are available in electronic form (FITS format) at the CDS.}} {At 22.5\\arcsec\\ resolution, the CO J=1-0 emission toward NRAO150 is 30-100\\% brighter at some velocities than seen previously with 1\\arcmin\\ resolution, and there are some modest systematic velocity gradients over the 220\\arcsec\\ field. Of the five CO components seen in the absorption spectra, the weakest ones are absent in emission toward NRAO150 but appear more strongly at the edges of the region mapped in emission. The overall spatial variations in the strongly emitting gas have Poisson statistics with rms fluctuations about equal to the mean emission level in the line wings and much of the line cores. The J=2-1/J=1-0 line ratios calculated pixel-by-pixel cluster around 0.7.  At 6\\arcsec\\ resolution, disparity between the absorption and emission profiles of the stronger components has been largely ameliorated. The \\twCO\\ J=1-0 emission exhibits i) remarkably bright peaks, $\\Tmb = 12-13\\K$, even as $4''$ from NRAO150; ii) smaller relative levels of spatial fluctuation in the line cores, but a very broad range of possible intensities at every velocity; and iii) striking kinematics whereby the monotonic velocity shifts and supersonically broadened lines in 22.5\\arcsec\\ spectra are decomposed into much stronger velocity gradients and abrupt velocity reversals of intense but narrow, probably subsonic, line cores.} {CO components that are observed in absorption at a moderate optical depth (0.5) and are undetected in emission at $1'$ resolution toward NRAO 150 remain undetected at $6''$ resolution. This implies that they are not a previously-hidden large-scale molecular component revealed in absorption, but they do highlight the robustness of the chemistry into regions where the density and column density are too low to produce much rotational excitation, even in CO. Bright CO lines around NRAO150 most probably reflect the variation of a chemical process, \\ie{} the C\\p-CO conversion. However, the ultimate cause of the variations of this chemical process in such a limited field of view remains uncertain.} ",
        "introduction": "Many detailed and sensitive studies of interstellar gas employ absorption spectra taken against the microscopic disks of stars or compact radio continuum sources.  We have recently used such spectra at radiofrequencies to elucidate an unexpectedly widespread molecular chemistry in diffuse clouds; for instance, see \\cite{LisLuc+06} and the many references therein).  Lamentably, observing gas in absorption along such single, isolated lines of sight does not generally allow the identification of the intervening features with recognizable physical bodies in space, and their possible association with other nearby entities (for instance dark clouds) remains hidden.  These limitations have obvious consequences for our understanding of the intervening medium and the absorption profiles themselves.  But in some rare cases it is becoming possible to observe the same feature in emission and absorption, even in the same chemical species, and so perhaps to characterize the nature of the intervening material by imaging it across the line of sight away from the background object.  This paper begins a short series of reports of studies undertaken to elucidate the nature of the gas occulting the compact extragalactic mm-wave background sources employed in our work.  It describes one case, using the J=1-0 and 2-1 mm-wave rotational transitions of carbon monoxide to image the particularly complex kinematics of relatively distant absorbing material along the low-latitude line of sight toward NRAO150 \\citep{CoxGue+88} at the highest usefully attainable resolution.  It is organized as follows.  Section 2 describes the sightline toward the blazar NRAO150 as it is understood from radiofrequency observations of CO and \\hcop\\ in emission and in approximately a dozen other atomic and molecular species in absorption.  Section 3 describes the techniques used to image and/or synthesize the pattern of foreground CO emission using singledish and interferometeric synthesis data.  Section 4 describes the results of the new observations and Section 5 is dedicated to models and Section 6 is a summary. ",
        "conclusions": "After a long series of investigations into the absorption line chemistry of galactic diffuse gas coincidentally seen toward distant radio point sources at low-moderate galactic latitude, we undertook to image the absorbing medium by mapping \\twCO\\ in emission around one source at the highest practicably obtainable spatial resolution.  There are many questions which might be asked of such data and the comparisons between relatively microscopic (sub-arcsecond pencil absorption beams and the much larger ($6-60''$) fan beams in emission.  The present work is a small step in this direction.  Some of the results are rather direct, for instance concerning the various disparities and differences between the absorption profiles and the existing 60\\arcsec-resolution KP CO data (Fig.~\\ref{fig:B0355f1}).  An atypical disparity between the strong emission and absorption profiles at $-10 \\kms$ was shown to arise because of insufficient spatial resolution in the emission and not, for instance, due to the presence of a warmer (more highly-excited) CO component seen in emission with low optical depth.  The weakness of the $-14$ and $-4.5/-3.5\\kms$ emission lines toward NRAO150, which are seen in absorption at moderate optical depth (0.5) in Fig.~\\ref{fig:B0355f1}, arises because the regions of sufficient collisional excitation and stronger emission are somewhat removed from the continuum position.  These clouds, which are common in our previous absorption surveys, are not a previously-hidden large scale molecular component revealed in absorption, but they do highlight the robustness of the chemistry into regions where the density and column density are too low to produce much rotational excitation, even in CO.  This is in fact the norm in optical absorption studies of CO.  We explored the spatial structure of the gas in several ways, \\ie{} by examining the maps at 22.5\\arcsec\\ and 6\\arcsec\\ resolution (Fig.~\\ref{fig:AverageMoments}) and by calculating the spatial statistics both as profile rms over the mean velocity profile (Fig.~\\ref{fig:ExcessRMS}) and as channel by channel probability distributions over the synthesized region (Fig.~\\ref{fig:PixelHistograms}). At the lower resolution, the velocity profile of rms fluctuations due to spatial structure, calculated as the channel by channel quadrature difference between the total rms and that seen in line-free channels, closely resembles the mean profile, except in a narrow interval near the core of the $-10\\kms$ line.  Such statistics are quasi-Poisson in character, consistent with a random macroscopically transparent ensemble of independent clumps.  However, when the gas is studied at higher resolution, where there is an even wider range of observed brightness at all velocities, the statistics in the core of the $-17.5\\kms$ noticeably change character and the rms is only 50\\% of the mean (Fig.~\\ref{fig:ExcessRMS}).  In this work, it appears that every profile seen at 22.5\\arcsec\\ with a linewidth greater than that attributable to the local sound speed (full-width at half-maximum $= 0.96\\kms$ for a pure-\\HH\\ gas with $\\TK = 40\\K$) is found to be composite when viewed at higher resolution: decomposable into multiple narrower components and/or having spatially-resolved velocity gradients in narrower-lined gas.  Whether this would continue in the broadest components seen at 5.8\\arcsec\\ resolution is an open question.  The full extent of the emission structure is realized only when the highest spatial and velocity resolution are viewed together in position--velocity tracks along portions of the gas where the medium may be supposed to be macroscopically optically thin (Fig.~\\ref{fig:Kine}).  A series of tracks parallel to but somewhat displaced from the peak of a particularly well defined filamentary or trunk structure in the gas around $-10\\kms$ was used to show the substructure in what are almost featureless position--velocity diagrams at 22.5\\arcsec\\ resolution.  The high resolution version resolves a monotonic and rather bland velocity gradient in the $-17.5\\kms$ feature into a striking series of kinks and abrupt reversals which strongly resemble simulations of turbulent flow of the sort shown by \\cite{FalLis+94}.  The flat profile core of the $-10\\kms$ component seen at 60\\arcsec\\ resolution, which is slightly split at 22.5\\arcsec\\ resolution (Fig.~\\ref{fig:B0355f1} and Fig.~\\ref{fig:ProfileEvolution}), shows an underlying triple nature (Fig.~\\ref{fig:ProfileBrightSpot} and Fig.~\\ref{fig:Kine}) whose individual components also show this striking pattern of velocity reversals when the medium is macroscopically thin and show some excursions in the line wing where it is thick.  In this connection, note that, in our data, it was not necessary to avoid the strongly-emitting optically thick line cores (at least they are optically thick toward NRAO150) in order to see such structure.  Furthermore, as noted just before, when studied at lower angular resolution, regions of the line core exhibit the same fractional level of fluctuation as the line wings for the component at $-17.5\\kms$.  At $-10\\kms$ the level of fluctuation is not very different at the two resolutions, and is still relatively large in the line core.  In the large, the CO excitation in the host gas sampled in the J=1-0 and J=2-1 lines at 22.5\\arcsec\\ resolution (Fig.~\\ref{fig:12co21-vs-12co10}) is readily explained as arising in low-moderate density n(H) $= 50 - 250 \\pccm$ diffuse gas (Fig.~\\ref{fig:ModelTBvsTB}).  The 12-13\\K{} \\twCO\\ lines seen in Fig.~\\ref{fig:ProfileBrightSpot}, at the high end of the brightness distribution (Fig.~\\ref{fig:PixelHistograms}) are probably regions of higher density (n(H) = 300 - 500 $\\pccm$) and column density N(\\HH) $> 10^{21} \\pscm$, N(CO) $> 10^{16} \\pscm$ where C\\p-CO conversion is more advanced even while the temperature is still elevated, $\\TK >> 10\\K$.  It is something of a coincidence that the \\twCO\\ brightnesses observed in diffuse and dark gas are so similar, given the profound differences in abundance and excitation: excitation in diffuse clouds is considerably sub-thermal and the temperatures are high, typically above 30\\K{}, and only a small fraction of the gas-phase carbon nuclei reside in CO -- typically 5\\% or less. This coincidence will complicate the interpretation of CO emission in general.  While absorption studies enable very precise determination of various chemical tracer column densitites, they can currently be done only on a limited number of pencil beam lines of sight. Those studies thus miss important information about the environment such as the geometry and the kinematics of the host gas, information which could help to unravel remaining mysteries like the abundance of \\hcop{} or CH\\p~\\citep[see \\eg{}][]{falgarone06,hilyblant07,hilyblant08}.  This paper is the first of a series in which we image diffuse gas around lines of sight which have previously been studied in absorption.  In this series our observations will range from the highest possible angular resolution in small fields of view (as as present) to lower angular resolution in fields of view wide enough to contain the whole structure of the host gas.  In this way we will relate such different gas and dust tracers as H I, CO, \\hcop{}, extinction, {\\it etc.} using whatever means are available to render the diffuse absorbing material both more readily visible and more completely understood."
    },
    "0808/0808.1727_arXiv.txt": {
        "abstract": "We present a simple formalism to interpret the observations of two galaxy statistics, the UV luminosity function (LF) and two-point correlation functions for star-forming galaxies at $z$$\\sim$$4$, $5$ and $6$ in the context of $\\Lambda$CDM cosmology. Both statistics are the result of how star formation takes place in dark matter halos, and thus are used to constrain how UV light   depends on halo properties, in particular halo mass. The two physical quantities we explore are the star formation duty cycle, and the range of UV luminosity that a halo of mass $M$ can have (mean and variance). The former directly addresses the typical duration of star formation activity in halos while the latter addresses the averaged star formation history and regularity of gas inflow into these systems. In the context of this formalism, we explore various physical models consistent with all the available observational data, and find the following: 1) the typical duration of star formation observed in the data is $ \\lesssim 0.4$ Gyr ($1\\sigma$), 2) the inferred scaling law between the observed $L_{UV}$ and halo mass $M$ from the observed faint-end slope of the luminosity functions is roughly linear out to $M\\approx 10^{11.5} - 10^{12}$\\hmsun\\ at all redshifts probed in this work, and 3) the observed $L_{UV}$ for a fixed halo mass $M$ decreases with time, implying that the star formation efficiency (after dust extinction) is higher at earlier times. We explore several different physical scenarios relating star formation to halo mass, but find that these scenarios are indistinguishable due to the limited range of halo mass   probed by our data. In order to discriminate between different scenarios, we discuss the possibility of using the bright-faint galaxy cross-correlation functions and more robust determination of luminosity-dependent galaxy bias for future surveys. ",
        "introduction": "In the last decade, substantial progress has been made in advancing our understanding of galaxy clustering in connection to dark matter halo clustering. Numerous surveys conducted  out to $z\\sim6$ ($t_{universe}\\approx 0.9$ Gyr) have selected large samples to measure the clustering as a function of galaxy properties such as color, luminosity, spectral type, and morphology \\citep[][]{norberg01, zehavi02, zehavi05, GD01, foucaud03, adelberger05, allen05, ouchi05, lee06, kashikawa06, coil06a, coil08, yoshida08}. These results have convincingly shown that the clustering strength of galaxies has a strong dependence on their physical properties. In general, the trend goes in a direction that more luminous (in the optical or UV) or redder galaxies are more strongly clustered in space than the less luminous or bluer ones.  The observed trends of galaxy clustering are similar to those of halos. The hierarchical theory of structure formation predicts that the halo clustering is a strong function of their masses and assembly history  \\citep*{mw96, gao05, gao07, wechsler06}.  Because galaxies formed inside dark matter halos, as baryonic matter is pulled into the gravitational potential wells of halos, cools and initiates star formation \\citep{wr78}, the astrophysical processes of galaxy formation are invariably linked to the characteristics of dark matter halos. The main halo properties include their sizes, masses, angular momentum, assembly history, and the internal distribution \\citep{nfw97, moore98}.  Recent evidence has further corroborated the halo--galaxy connection. \\citet{zehavi04} have measured the galaxy two-point correlation function (CF) of the Sloan Digital Sky Survey (SDSS) galaxies with unprecedented high precision. From these measures, they have detected a small feature in the shape of the correlation function at a physical scale of $\\approx$1 Mpc.  The observed scale of the bump coincides with the physical scale where the transition from the one-halo term ($r\\lesssim 1$ \\mpc) to the two-halo term ($r\\gtrsim 1$ \\mpc) occurs. The former arises from the spatial correlation between the parent halo and its substructure (subhalos) and between subhalos, while the latter arises from the correlation between distinct halos. Soon after, similar transitions were found at larger look-back times for galaxies selected in the rest-frame optical at $z\\sim1$ \\citep{coil06b}, $z\\sim2.5$ \\citep{quadri08}, and in the UV at $z\\sim3$ \\citep{hildebrandt07}, $z\\sim4$ and $5$ \\citep{ouchi05, lee06}. The physical scale of the transition is observed to increase with time (decreasing redshift), which is expected because halos grow in size.  The two independent lines of evidence, the observed luminosity/color dependence of galaxy clustering and the detection of a transition scale in the CFs, lend support to the close connection between halos and galaxies.  A logical next step is to constrain the scaling law of the two properties, namely, galaxy luminosity $L$ and halo mass $M$, in order to obtain a more detailed picture of the physical processes. Furthermore, by comparison of the scaling laws at different cosmic epochs, one can begin to understand the time sequence of galaxy formation in the context of the halo evolution \\citep[e.g.,][]{zhengetal07, mwhite07,conroy07,conroy08}.  Many authors have successfully modeled such scaling relations for local galaxies based on surveys such as the SDSS  \\citep*{yang03, vandenbosch03b}. They have used joint constraints of the observed luminosity function (LF) and the clustering measures for the same galaxies. Using the 2MASS data, \\citet{vale04, vale06}  modeled the scaling relation by directly mapping the shape of the halo mass function \\citep[the number density of halos as a function of halo mass:][]{ps74, st99, sheth01} to the galaxy LF (the number density of galaxies as a function of luminosity), assuming that there is a unique one-to-one relation between the halo mass and galaxy light. A similar abundance-matching method was used to constrain the relation of stellar masses to halo masses for galaxies at intermediate redshift \\citep{conroy08}. These models assume that each halo hosts a visible galaxy only above the mass threshold of halos (given by the integrated LF constraint).  The assumption is a reasonable one in the local universe, because the wavelength ranges probed by these surveys trace the general stellar population, in other words, the integrated star formation history over the course of the galaxy's entire history, and thus is insensitive to the details of a galaxy's recent star formation history. Hence, the halo mass, as a robust indicator of the area's local density contrast, is well correlated with the stellar masses of the galaxy therein.  High-redshift galaxy samples, however, are often selected in the rest-UV, which traces the instantaneous formation of massive stars \\citep[e.g.,][]{madau96}. Moreover, the intrinsic UV luminosity is obscured and reddened by dust in the interstellar medium to add an additional uncertainty to the halo--galaxy association \\citep[][]{conroy08b}.  Hence, at high redshift, the modeling of such a relation requires extra caution as there is no reason to believe that every halo hosts a currently star-forming galaxy. The advantage, however, is that the same uncertainty that we face in modeling this connection will give us important clues to understanding various star formation processes, such as their typical duration and the dependence of star formation rate on the host halo mass.  There are reasonable prospects of achieving such a goal, as the observed clustering properties indicate that the observed UV luminosity (after dust extinction) still correlates strongly with their clustering strength \\citep[and thus, with halo mass; e.g.,][]{GD01, adelberger05, allen05, ouchi04b, ouchi05, lee06,   overzier06, kashikawa06, yoshida08}. By using these constraints, combined with the UV LF measured at the same cosmic epochs \\citep[][]{gabasch04, ouchi04a, sawicki06, yoshida06, bouwens07, iwata07, reddy08, mclure08}, we can constrain the typical duration of star formation as well as the physical scaling law between galaxy UV luminosity and halo mass.  Our effort is motivated by the dramatic improvement in our understanding of dark matter substructures from high-resolution dark matter (DM) simulations and analytic calculations \\citep{kravtsov04, gao04, delucia04, zentner05, wechsler06}, made more solid given the recent tight constraints on cosmological parameters from the WMAP and other recent studies \\citep{spergel03,spergel07,komatsu08}. \\citet{kravtsov04} and \\citet*{conroy06} have demonstrated that halos and subhalos identified in DM simulations provide an excellent match to the observed correlation functions at all scales ($\\theta \\geq 2-3$\\arcsec\\ at $z\\sim4$, for example).  The tidal stripping and mass losses of small halos occurring as they enter into the potential well of a larger halo is better understood with high-resolution simulations \\citep{gao04, delucia04}.  These dynamical processes may play an important role in shaping the observed galaxy statistics. The strong dynamic evolution experienced by subhalos may not be felt as strongly for the embedded galaxies, because they are more tightly bound at the core of the system gravitationally \\citep[e.g.,][]{moore96, klypin99, hayashi03, kravtsov04, nagai05}.  In this paper, we attempt to take advantage of the recent progress seen in both the observed galaxy statistics, namely, the UV LF and correlation functions of high-redshift galaxies, and the analogous quantities for the DM halos, to understand their statistical association, and thus, the star formation physics of high-redshift galaxies in relation to their local environments. Our approach is similar in character to halo occupation distribution (HOD) models which assume that all galaxies are harbored in DM halos \\citep[e.g.,][]{berlind02,bullock02, zheng05}, but is generalized to accommodate galaxy luminosity as a joint variable similar to the conditional LF formalism \\citep*[][]{yang03, vandenbosch03b}. We discuss the details of how our approach differs from the standard HOD formalism in the next section. We also note that our methodology is complementary to ab initio calculations of semi-analytic models or hydrodynamic simulations, for which many detailed physical processes need to be modeled to produce the observable properties of galaxies \\citep[e.g.,][]{kauffmann93, cole00, somerville01, wechsler01, bower06, croton06, delucia07, nagamine07}. Furthermore, our empirical approach will help provide insight into the physical recipes implemented in these simulations.  The main goals of this paper are 1) to build a realistic and empirical model well suited to analyze high-redshift data,  2) to interpret the observed galaxy statistics  simultaneously in the light of the properties of $\\Lambda$CDM halos,  and 3) to draw general conclusions about the physics of star formation when the universe was less than 2 billion years old. We will take advantage of the newly available observational measures made for the high-redshift star-forming galaxies.   This paper is organized as follows. In Section 2, we describe the methodology providing the detailed calculations for the model predictions of the galaxy LF and correlation functions. Readers who are not interested in the details of the methodology  may skip to Section 3, where we describe the data sets used for the measurements, and in Section 4, the improved measures of the auto-correlation functions as well as the LF and cross--correlation function. In Section 5 we report our main results, and the results and their physical implications are discussed in Section 6. Finally, in  Section 7, we present the model predictions for the galaxy cross-correlation function, which may help overcome the current limitations of the data to discriminate different physical scenarios of star formation. All magnitudes in this work are in the AB scale \\citep{oke83}. We use a cosmology with $\\Omega_m = 0.3$, $\\Omega_\\Lambda=0.7$, $\\sigma_8=0.9$, $\\Gamma = 0.21$, $H_0= 100 h$ km s$^{-1}$ Mpc$^{-1}$ with $h=0.7$ and the baryonic density $\\Omega_b=0.04$. ",
        "conclusions": "We have used the observed UV LF and correlation function measures for star-forming galaxies at $z\\sim4$, $5$, and $6$ to infer the nature of star formation and its dependence on halo mass, in particular for the sub-$L^*$ galaxies.  The main conclusions from this work are as follows:  1. The star formation duty cycle of Lyman-break galaxies should be less than $<0.35$ Gyr at both $z\\sim4$ and 5. The best-fit duty cycle value for $z\\sim4$ is 15\\%-60\\% ($1\\sigma$), and $<70$\\% for $z\\sim5$. The relatively short time scale during which galaxies are visible in the UV implies that the galaxies observed at $z\\sim4$ are unlikely to be the direct descendants of those at $z\\sim5$ unless the star formation is recurrent after a long intermission. \\\\ 2. The observed UV luminosity scales approximately linearly with the halo mass in order to reproduce the faint-end slope of the UV LF $\\alpha\\approx$-1.7 observed at $z\\sim4-6$, for galaxies less luminous than the characteristic value $L^*$. In this interpretation, the constant faint-end slope with redshift is a direct result of, 1) the low-mass slope of the total halo mass function remains constant with redshift, and 2) the observed UV luminosity scales with the halo  mass with a power-law slope close to unity ($\\alpha=$ 0.9-1.2) at $z$=4-6.\\\\ 3. While the slope of the $L$-$M$ scaling law does not change with redshift, the amplitude of the relation decreases with cosmic time, such that for a fixed halo mass, $z\\sim5$ galaxies appear brighter by $\\approx$ 0.3 mag than $z\\sim4$ galaxies. If the dust properties do not change significantly at those redshifts, this implies that star formation efficiency per halo mass was higher at earlier times consistent with \\citet{lee06} results. \\\\ 4. We interpret the nonevolution of the normalization parameter $\\phi^*$ with redshift observed at $z\\sim$4-6 as a result of the two competing processes canceling each other: the number density of halos for a fixed halo mass increases with time, while the average UV luminosity in halos of a fixed halo mass decreases with time.\\\\ 5. The star formation in massive halos ($M>10^{10.8}$\\hmsun) should be relatively quiescent, and thus can be described by a slowly varying star formation history. The degree of burst can  be a mildly varying function of halo mass at this regime, and it may be attributed to the merger-induced star formation in massive halos (as the halo merger rate is also a mildly increasing function of mass). Data from  wide-field surveys are crucially needed to quantify the contribution from bursty star formation in further detail. \\\\ 6. The average star formation histories in  low-mass halos ($M<10^{10.5}$\\hmsun) is not as well constrained from the current data mainly due to the uncertainties in the true large-scale bias. The main mode of star formation at this regime is crucial to understand the formation histories of the majority of galaxies detected in the rest-UV surveys:  whether they are forming stars as quiescently as their brighter counterparts (Scenario 2), or they represent a small fraction of low-mass halos undergoing ``bursty'' star formation (Scenarios 1 and 3)."
    },
    "0808/0808.0517_arXiv.txt": {
        "abstract": "We present \\sauron\\ integral-field stellar velocity and velocity dispersion maps for four double-barred early-type galaxies: NGC\\,2859, NGC\\,3941, NGC\\,4725 and NGC\\,5850. The presence of the inner bar does not produce major changes in the line-of-sight velocity, but it appears to have an important effect in the stellar velocity dispersion maps: we find two \\emph{$\\sigma$}-hollows of amplitudes between 10 and 40 km~s$^{-1}$ on either side of the center, at the ends of the inner bars. We have performed numerical simulations to explain these features. Ruling out other possibilities, we conclude that the \\emph{$\\sigma$}-hollows are an effect of the contrast between two kinematically different components: the high velocity dispersion of the bulge and the more ordered motion (low velocity dispersion) of the inner bar. ",
        "introduction": "Double-barred galaxies are rather common systems in the Universe, representing $\\sim$1/3 of all barred galaxies \\citep{2002AJ....124...65E,2002ApJ...567...97L}. These structures are mainly seen by surface brightness profile decomposition, which requires not only a bulge and a disk but also two additional components with non-axial symmetry. The presence of inner bars does not seem to correlate with the morphological type of the host galaxy and there is also no preferred angle between the two structures \\citep{1993A&A...277...27F,1995A&AS..111..115W}. This property suggests that both bars rotate independently, as seen in observations and consistent with numerical simulations \\citep[e.g.][]{2003ApJ...599L..29C,2004ApJ...617L.115E,1989Natur.338...45S,2002ApJ...565..921S}. Inner bars play an important role in the evolution of galaxies as they might transport gas to the central regions, where they may trigger the formation of stars that, in turn, lead to the appearance of new structures such as disk-like components \\citep{1995FCPh...15..341M}. Theoretically, this phenomenon cannot be achieved with a single bar, since the flow of material would stop before reaching the galactic center (e.g. at the Inner Lindblad resonance). The presence of two embedded bars makes it possible and may even help to feed the central AGN \\citep{1989Natur.338...45S,1990Natur.345..679S}.  The study of the orbital make-up of different morphological components in double-barred galaxies is one of the best ways to understand their internal structure and dynamics. Nevertheless, this is a rather complex issue to address as there are three fundamental frequencies for the regular orbits \\citep[i.e. two for the bars and one for the free oscillations;][]{MA07} instead of the two frequencies as in the case of a single bar. As a result, they do not have closed and periodic orbits, so \\citet{2000MNRAS.313..745M} introduced the \\emph{loop} concept to represent the stable orbits in a double-barred potential.  A more popular approach is the use of numerical simulations to describe the structures seen in real observations, such as in the recent works by \\citet{2007ApJ...657L..65H} and \\citet{2007ApJ...654L.127D}. In their simulations, these authors generate long-lived double bars with and without a dissipative component, respectively. However, most existing observations have relied on long-slit spectroscopy along only a few position angles \\citep[e.g.][]{2001A&A...368...52E}, which makes the comparison with the models somewhat difficult. Better suited integral-field spectroscopy has generally not been used to observe these systems, with the notable exceptions of \\citet{2001BSAO...51..140M} and \\citet{2004A&A...421..433M}.  Here we present \\sauron\\ integral-field spectroscopy results on the stellar kinematics of four double-barred early-type galaxies. We will focus on the detailed analysis of the stellar velocity dispersion maps, that reveal a peculiar feature, not seen before, along the extremes of the inner bar. We present the observational evidence and discuss possible interpretations using numerical simulations. A more complete account of the ionised gas, stellar population properties and supporting numerical simulations will be the subject of a forthcoming paper. ",
        "conclusions": "We have presented 2D kinematical maps of double-barred early-type galaxies covering the whole nuclear region. We have found $\\sigma$-hollows appearing at the ends of the inner bar for our four objects. The main result presented in this paper is that these hollows are signatures of inner bars, which indicates that it they must have significantly lower velocity dispersions than the surrounding bulges. This means that inner bars are cold systems, and are not related to triaxial bulges \\citep{2004ARA&A..42..603K}. Moreover, the observations presented here put constraints on the degree of rotational over pressure support for these structural components. Finally, since the $\\sigma$-hollows are not due to other structural or kinematical components (e.g. inner disks), they may be used to identify inner bars from a purely stellar kinematic analysis.\\looseness-2  It is likely that previous numerical simulations were unable to predict the observed $\\sigma$-hollows because they did not include a sufficiently dynamically hot component (i.e. classical bulge). However, we have been able to reproduce this feature with simulations including a hot bulge, a single small bar and a disk. The $\\sigma$-hollows can be seen due to a contrast between the high velocity dispersion of the bulge and the low velocity dispersion of the bar. The amplitude of the hollows depends on the relative contributions to the total flux and the relative sizes of the bulge and the inner bar. While young stellar populations can help to lower the stellar velocity dispersion, no clear signatures of them are found in our sample. Based on the pieces of evidence presented in this paper, we think that the $\\sigma$-hollows are simply a matter of contrast. \\\\"
    },
    "0808/0808.0721_arXiv.txt": {
        "abstract": "Highly magnified lensed galaxies allow us to probe the morphological and spectroscopic properties of high-redshift stellar systems in great detail. However, such objects are rare, and there are only a handful of lensed galaxies which are bright enough for a high-resolution spectroscopic study with current instrumentation. We report the discovery of a new massive lensing cluster, SDSS\\,J120923.7$+$264047, at $z=0.558$. Present around the cluster core, at angular distances of up to $\\sim40''$, are many arcs and arc candidates, presumably due to lensing of background galaxies by the cluster gravitational potential. One of the arcs, $21''$ long, has an $r$-band magnitude of 20, making it one of the brightest known lensed galaxies. We obtained a low-resolution spectrum of this galaxy, using the Keck-I telescope, and found it is at redshift of $z=1.018$. ",
        "introduction": "\\label{Introduction}  High redshift galaxies ($z\\gtorder1$) are out of reach for current high-resolution spectrographs. Fortunately, such galaxies may be highly magnified, when lensed by galaxies or cluster of galaxies, making them bright enough for high-resolution spectroscopic studies. However, high magnification lensed galaxies are rare, and only small number of targets suitable for high-resolution spectroscopic studies are known. Among these are: MS1512$-$cB58 (Yee et al. 1996), with $r=20.4$\\,mag at $z=2.72$; an $r=20.3$\\,mag Lyman break galaxy at $z=3.07$ (Smail et al. 2007); and finally a highly magnified galaxy, at $z=2.73$, with $r=19.2$\\,mag (Allam et al. 2007). Studies of these objects (e.g., Teplitz et al. 2000; Pettini et al. 2000; 2002; Baker et al. 2004; Schaerer \\& Verhamme 2008) provided valuable information on their metallicity and inter-stellar medium.   In this paper we report on the discovery of a previously unknown rich cluster of galaxies at $z=0.558$. The cluster containing a bright ``giant arc'', which is an appropriate target for high-resolution spectroscopic studies. In addition, the cluster, which has a considerable Einstein radius of probably $\\sim20-40''$, contains a large number of additional arcs and lensed galaxies. This cluster is probably as rich as the Abell 1689 lensing cluster (Broadhurst et al. 2005), and therefore it is an excellent target for detailed strong lensing and weak lensing modeling (e.g., Brada\\u{c} et al. 2006; Halkola et al. 2006).  The outline of this paper is as follows. We describe the observations in \\S\\ref{Obs} and in \\S\\ref{Arcs} we discuss the cluster and arcs along with their measured properties. Finally, we briefly discuss the results in \\S\\ref{Sum}. ",
        "conclusions": "\\label{Sum}   Strong and weak lensing are being used to trace the mass distribution of galaxy clusters. However, detailed modeling of the mass distribution of the cores of galaxy clusters requires the identification of multiple images of lensed sources. Such an identification needs good multi-band color information, or/and good spatial resolution (e.g., Broadhusrt et al. 2005; Sharon et al. 2006). Indeed we are planing additional multi-band observations of this cluster, which will enable a detailed modeling of the cluster gravitational potential.   To confirm the presence of the cluster, in Fig.~\\ref{fig:SDSS1209_RedSequence} we show the color magnitude diagram for SDSS galaxies found with $2'$ from the cluster bright galaxy (black circles). For comparison, we also show all the SDSS galaxies found between $5'$ and $1^{\\circ}$ from the cluster bright galaxy (gray dots). The solid (dashed) red-lines display the locus of Elliptical (Sb) galaxies as a function of redshift. These lines take into account the Galactic extinction ($E_{B-V}=0.019$\\,mag) in the direction of the cluster (Schlegel et al. 1998; Cardelli et al. 1989). It is evident, that in the cluster direction there is an excess of high-redshift galaxies, relative to the field. \\begin{figure} \\centerline{\\epsfxsize=8cm\\epsfbox{f5.eps}} \\caption{The color-magnitude diagram of SDSS galaxies in the direction of the cluster. The black circles represent galaxies found with $2'$ from the cluster bright galaxy. For comparison, we also show all the SDSS galaxies found between $5'$ and $1^{\\circ}$ from the cluster bright galaxy (gray dots). The solid (dashed) red-lines display the locus of Elliptical (Sb) galaxies as a function of redshift. These lines take into account the Galactic extinction ($E_{B-V}=0.019$\\,mag) in the direction of the cluster (Schlegel et al. 1998; Cardelli et al. 1989). The galaxies colors as a function of redshift are based on synthetic photometry of galaxy templates from Kinney et al. (1996). The synthetic photometry was performed using the code of Poznanski et al. (2002).} \\label{fig:SDSS1209_RedSequence} \\end{figure}   At the approximate position of the bright cluster galaxy, we find a known ROSAT source (Voges et al. 2000). Using the PIMMS webtool\\footnote{http://cxc.harvard.edu/toolkit/pimms.jsp}, we converted the ROSAT/PSPC count rate to an unabsorbed flux of $6.9\\times10^{-13}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$, assuming a Galactic neutral Hydrogen column density of $1.8\\times10^{20}$\\,cm$^{-2}$ (Dickey \\& Lockman 1990; Kalberla et al. 2005) and a power-law spectrum with a photon index of $\\Gamma=1$ (assuming thermal optically thin Bremsstrahlung). Given the luminosity distance to the cluster, this flux implies a luminosity of $\\sim8\\times10^{44}$\\,erg\\,s$^{-1}$ in the 0.1--2.4\\,keV band. Using the X-ray luminosity-mass relation given by Reiprich \\& B\\\"{o}hringer (2002), we estimate that the $M_{200}$ mass\\footnote{$M_{200}$ is the mass enclosed within $r_{200}$, which is the radius in which the mean density equals 200 times the mean density of the Universe (e.g., Navarro, Frenk, \\& White 1997).} of the cluster is approximately $10^{15}$\\,M$_{\\odot}$.  We can also estimate the mass of the cluster from its Einstein radius: \\begin{equation} \\theta_{E}=\\Big[ \\frac{4GM(\\theta_{E})}{c^{2}}\\frac{D_{ls}}{D_{l}D_{s}} \\Big]^{1/2}. \\label{ER} \\end{equation} The Einstein radius is a robust measure of the lens mass within this radius. Although we did~not identify the counter-image(s) of arc-A we can estimate the Einstein radius based on the distance of the arc from the cluster bright galaxy ($\\sim11''$). Using Eq.~\\ref{ER} we estimate that the total mass enclosed within 90~kpc (i.e., the Einstein radius, $11''$, at $z\\approx1$) is $\\sim6\\times10^{13}$\\,M$_{\\odot}$.               To summarize, we report on the discovery of SDSS\\,120923.7$+$264047 -- a massive cluster of galaxies at $z=0.558$. The cluster hosts many prominent arcs, including a giant arc with $R$-band magnitude of 20.0, at $z=1.018$. Moreover, presumably the cluster has a large Einstein radius, suggesting that it is one of the richest strong lensing clusters known."
    },
    "0808/0808.2986_arXiv.txt": {
        "abstract": "{We report moderate magnetic flux of $450$\\,G$\\,< Bf < 750$\\,G (3$\\sigma$) on the nearby M5.5 flare star Proxima Centauri. A high resolution UVES spectrum was used to measure magnetic flux from Zeeman broadening in absorption lines of molecular FeH around 1\\,$\\mu$m.  The magnetic flux we find is relatively weak compared with classical strong flare stars, but so are Proxima\u2019s flaring rates and actual emission levels. We compare what is known about the rotation rate, Rossby number, and activity levels in this star to relations between these quantities that have been recently being developed more generally for M dwarfs. We conclude that the magnetic flux is higher than the best estimates of the Rossby number from period measurements. On the other hand, the activity levels on Proxima Centauri are at the high end of what could be expected based on the measured field, but not so high as to exceed the natural scatter in these relations (other stars lie along this high envelope as well).  } ",
        "introduction": "\\label{sect:Introduction}  Proxima Centauri ($\\alpha$ Cen C, GJ~551\\,C, in the following Prox~Cen), our closest stellar neighbor, has been known to be a flare star for more than 50 years \\citep{Thackeray50}. Prox~Cen is of particular importance for our understanding of very cool stars, because we can observe it with very high precision. On Prox~Cen, activity can be detected at very low level, and great effort has been put into monitoring campaigns to investigate it.  At a spectral type of M5.5, Prox~Cen is probably completely convective.  It shows persistent H$\\alpha$ emission and flaring, which may lead to the expectation of rapid rotation ($\\ga 3$\\,km\\,s$^{-1}$). However, the strength of quiescent activity in Prox~Cen is relatively low compared with other active mid-M stars \\citep{Mathioudakis91, Patten94, Hawley96, Bochanski07}, and the flare rate is as well \\citep{Walker81}.  The relation between rotation and magnetic activity is well known in solar-type stars \\citep[e.g.][]{Pizzolato03}. It is observed in very slowly rotating early-M dwarfs \\citep[$v\\,\\sin{i} \\ga 1$\\,km\\,s$^{-1}$,][]{Reiners07}, too. \\citet{Mohanty03} found that the rotation-acitivity relation is valid in low-mass stars down to spectral class M9, at least to the extent that rapid rotators show activity saturation. In solar-type stars, chromospheric and coronal activity is believed to be due to magnetic fields generated through a magnetic dynamo, and \\citet{RB07} found that magnetic flux scales with H$\\alpha$ activity also in the late-M dwarfs. This all leads to the expectation that activity on a flare star like Prox~Cen is due to strong magnetic fields generated through a magnetic dynamo that may be driven by rotation.  There has been considerable effort to determine the rotation period of Prox~Cen. \\citet{Benedict98} searched for a rotation period in HST FGS data claiming a rotation period of $P = 83.5$\\,d, and \\citet{Guinan96} report a period of 31.5\\,d from IUE data. From the intensity of Mg\\,II emission lines in the same IUE data, \\citet{Doyle87} predicted a rotation period around 51\\,d according to the rotation-activity relation. However, \\citet{Doyle87} also predicted a rotation period of 27~d for UV~Cet. UV~Cet has a projected rotational velocity of $v\\,\\sin{i} = 32.5$\\,km\\,s$^{-1}$ \\citep{Mohanty03}, i.e. a period of less than about half a day, which means that an estimate for the rotational period from Mg\\,II data is probably not very useful in mid-M dwarfs. Recently, \\citet{Kiraga07} reported a rotation period of 82.5\\,d for Prox~Cen, very close to the 83.5\\,d period reported by \\citet{Benedict98}. Altogether, there is evidence that $\\sim 83$\\,d is the true rotational period of Prox~Cen.  The rotation velocity of Prox~Cen is probably rather low, but still it shows persistent flaring activity. Although its flare rate is below those of the very active flare stars like CN~Leo, it is still an active star and strong magnetic flux might be expected.  In this paper we present the first direct measurement of the magnetic flux on Prox~Cen. ",
        "conclusions": "\\subsection{Activity and Magnetic Flux}  Active M dwarfs typically exhibit magnetic flux of a few kG \\citep{RB07}. However, the magnetic flux on Prox~Cen is substantially lower than 1\\,kG although it clearly is an ``active'' star.  Does Prox~Cen have particularly weak magnetism given its level of magnetic activity?  The flare rate of Prox~Cen is comparably low; \\citet{Walker81} report 35 Flares during 25.25\\,h, which means it is less actively flaring than UV~Cet, CN~Leo, or EQ~Peg. In particular, the flare rate of $\\sim 1.4$\\,h$^{-1}$ found by \\citet{Walker81} is only about 60\\,\\% the flare rate of CN~Leo found by \\citet{Kunkel73} at comparable U-band sensitivity. This indicates that Prox~Cen may be a flare star with relatively little activity.  Several measurements of H$\\alpha$ emission on Prox~Cen are reported in the literature. \\citet{Mathioudakis91} measured an H$\\alpha$ equivalent width of 1.67\\,\\AA. \\citet{Patten94} report a minimum of 3\\,\\AA, and 27\\,\\AA\\ during flare state, and \\citet{Hawley96} found an H$\\alpha$ emission line of 3.01\\,\\AA\\ equivalent width. The mean H$\\alpha$ equivalent width in active M5 dwarfs is $5.9 \\pm 1.1$\\,\\AA\\, \\citep[][although one has to keep in mind that this is from data of lower resolution]{Bochanski07}. The normalized H$\\alpha$ luminosities corresponding to H$\\alpha$ equivalent widths of 1.67\\,\\AA\\ and 3.0\\,\\AA\\ are $\\log{L_{{\\rm H}\\alpha}/L_{\\rm bol}} = -4.2$ and $-4.0$, respectively \\citep[see][]{RB08a}.\\footnote{\\citet{Hawley96} calculate a value of $\\log{L_{{\\rm H}\\alpha}/L_{\\rm bol}} = -4.2$ for an equivalent width of 3.0\\,\\AA, i.e., 0.2 dex lower than our calibration.}  According to \\citet{Mohanty03}, $\\log{L_{{\\rm H}\\alpha}/L_{\\rm bol}} \\approx -4.2$ is the level where activity saturates in stars of spectral type M5.5--M8.5, i.e., activity in Prox~Cen is barely in the regime of saturation. In fact, the lowest value of H$\\alpha$ emission (1.67\\,\\AA) is comparable to H$\\alpha$ measured in GJ\\,1286 (1.50\\,\\AA) by \\citet{RB07}.  This M5.5 star exhibits magnetic flux of $0.4 \\pm 0.2$\\,kG (1$\\sigma$), which is very similar to Prox~Cen (and GJ~1286 is a slow rotator).  Another proxy of magnetic activity are X-rays. X-ray emission of Prox~Cen was measured with ROSAT and XMM: The X-ray flux measured with ROSAT is $\\log{L_{\\rm X}} = 27.2$\\,erg\\,s$^{-1}$ \\citep{Hunsch99}.  At $M_{\\rm bol} = 11.98$ \\citep{Hawley96} this means that the normalized X-ray luminosity is $\\log{L_{\\rm X}/L_{\\rm bol}} = -3.5$.  Newer observations from XMM are consistent with this result \\citep[see][]{NEXXUS}. \\citet{Kiraga07} report a normalized X-ray luminosity of $\\log{L_{\\rm X}/L_{\\rm bol}} = -3.8$. This value probably comes from the two longest ROSAT PSPC observations only, for which $\\log{L_{\\rm X}} \\approx 26.9$ is provided by \\citet{NEXXUS} (while $\\log{L_{\\rm X}} \\approx 27.4$ is given for the two XMM exposures).  For a mid-M dwarf, a normalized X-ray luminosity between $\\log{L_{\\rm X}/L_{\\rm bol}} = -3.5$ and $-3.8$ is consistent with a normalized H$\\alpha$ luminosity between $-3.7$ and $-4.2$ \\citep[although the scatter is fairly large, see][]{RB07}. This is consistent with H$\\alpha$ measurements (see above).  Both, X-ray- and H$\\alpha$ emission put Prox~Cen on the low end of active M dwarfs.  \\citet{RB07} investigated the relation between magnetic flux and H$\\alpha$ emission among M dwarfs. From their sample, a star of spectral class M5.5 with a normalized H$\\alpha$ luminosity of $\\log{L_{{\\rm H}\\alpha}/L_{\\rm bol}} \\approx -4$ can be expected to have a magnetic flux level on the order of $Bf \\approx 1$\\,kG.  For $Bf \\approx 600$\\,G, i.e. the magnetic flux found on Prox~Cen, H$\\alpha$ emission on the order of $\\log{L_{{\\rm H}\\alpha}/L_{\\rm bol}} = -4.2$ can be expected. Thus, the magnetic flux estimated from the lowest X-ray and H$\\alpha$ detections is consistent with the magnetic flux measured here, while the emission seen in all X-ray and H$\\alpha$ data is scattering around a somewhat higher level. This leads to the conclusion that Prox~Cen has somewhat higher activity than would be predicted from its magnetic field, but not an unreasonably higher amount, given that the scatter in all these relations is substantial.  \\subsection{Rotation and Rossby number}  In M5.5--M8.5 stars, the rotation-activity relation saturates at a surface velocity of about 10\\,km\\,s$^{-1}$ \\citep{Mohanty03}. However, there are also slowly rotating mid-M stars that exhibit strong magnetic activity. While rapid rotators always seem to be active, the opposite is not true. Prox~Cen shows no detectable Doppler broadening due to rotation, its projected rotation velocity is smaller than $v\\,sin{i} \\approx 3$\\,km\\,s$^{-1}$.  \\citet{RB07} found that all M dwarfs with magnetic flux $Bf < 1$\\,kG are slow rotators ($v\\,\\sin{i} < 3$\\,km\\,s$^{-1}$). The low magnetic flux found on Prox~Cen is consistent with slow rotation.  At a radius of $R=0.145$\\,R$_{\\sun}$ \\citep{Segransan03}, the spectroscopic limit on the rotation velocity translates into a projected rotation period of $P/\\sin{i} > 1$\\,d.  There is growing evidence that the rotation period of Prox~Cen is on the order of 80\\,d \\citep{Benedict98, Kiraga07}, which is consistent with the lack of Doppler broadening (although it is almost two orders of magnitude below the detection limit). That long a rotation period would imply an extremely low surface rotation velocity of $v = 90$\\,m\\,s$^{-1}$.  The turnover time of Prox~Cen is in the range 70--100\\,d \\citep{Gilliland86}. If the rotation period of Prox~Cen is on the order of 80\\,d, its Rossby number is on the order of $Ro = 1$ (or $\\log{Ro} = 0$). This high a Rossby number would put Prox~Cen among the stars occupying the very beginning of the linear part of the rotation-activity relation, i.e. among the stars with the lowest measurable activity. From the relation between Rossby number and magnetic flux \\citep{Saar01, RB08b} one would expect magnetic flux on the order of 100\\,G for such a small Rossby number. Normalized X-ray activity would be on the order of $\\log{L_{\\rm X}/L_{\\rm bol}} = -5.0$ \\citep{Pizzolato03, Kiraga07}, and H$\\alpha$ activity also would be at about $\\log{L_{{\\rm H}\\alpha}/L_{\\rm bol}} \\approx -5$ implying that H$\\alpha$ would probably not be detectable as an emission line. The conclusion from the various rotation-activity relations is that from H$\\alpha$ emission, X-ray emission, and from magnetic flux, one would estimate a Rossby number on the order of $\\log{Ro} \\approx -0.6$. Together with a convective overturn time of $\\tau_{\\rm conv} = 70..100$\\,d, this leads to a rotation period estimate between 17\\,d and 25\\,d rather than 80\\,d. To reach $\\log{Ro} \\approx -0.6$ at a rotation period of 80\\,d, the turnover time needs to be on the order of 300\\,d rather than 100\\,d.  We conclude that Prox~Cen is a slow rotator with a rotation period longer than about 15\\,days, but its activity and magnetic flux are indicative of a rotation period substantially shorter than 80\\,d. It is less active than more rapid rotators like CN~Leo or UV~Cet. Its level of chromospheric and coronal activity is relatively high for its moderate magnetic flux or Rossby number, but given the substantial scatter in the relations between these quantities, it merely lies with other stars on the high activity side of the linear relation between rotation, activity, and magnetic flux in M dwarfs."
    },
    "0808/0808.0384_arXiv.txt": {
        "abstract": "We study the progenitor dependence of the black hole formation and its associated neutrino signals from the gravitational collapse of non-rotating massive stars, following the preceding study on the single progenitor model in \\citet{sum07}. We aim to clarify whether the dynamical evolution toward the black hole formation occurs in the same manner for different progenitors and to examine whether the characteristic of neutrino bursts is general having the short duration and the rapidly increasing average energies. We perform the numerical simulations by general relativistic $\\nu$-radiation hydrodynamics to follow the dynamical evolution from the collapse of pre-supernova models of 40M$_{\\odot}$ and 50M$_{\\odot}$ toward the black hole formation via contracting proto-neutron stars.  % For the three progenitor models studied in this paper, we found that the black hole formation occurs in $\\sim$0.4--1.5 s after core bounce through the increase of proto-neutron star mass together with the short and energetic neutrino burst. We found that density profile of progenitor is important to determine the accretion rate onto the proto-neutron star and, therefore, the duration of neutrino burst.  % We compare the neutrino bursts of black hole forming events from different progenitors and discuss whether we can probe clearly the progenitor and/or the dense matter. ",
        "introduction": "\\label{sec:intro}  The final fate of massive stars is one of key issues in stellar physics as well as astrophysics. The massive stars beyond $\\sim$8M$_{\\odot}$ are the origin of compact objects, i.e. neutron stars and black holes, as the outcome of gravitational collapse at the end of their lives. They contribute also to the production of heavy elements through the explosive nucleosynthesis and to the high energy phenomena with the bursts of electro-magnetic radiations and neutrinos, which in turn play important roles in the evolution of galaxies. The massive stars in the wide range between $\\sim$8M$_{\\odot}$ and $\\sim$100M$_{\\odot}$ lead to various explosive phenomena such as core-collapse supernovae, hypernovae (or gamma ray bursts) and failed supernovae \\citep{heg03}. Such energetic phenomena are dependent on the properties of progenitors such as masses, rotations and profiles, for example, as in the mass limit of ordinary supernovae and the branches of hypernovae and failed supernovae as a function of mass \\citep{mae03}. We focus here on the fate of non-rotating massive stars, as an extreme of failed supernovae, in the mass range of $\\sim$40--50M$_{\\odot}$.  Such a death of massive stars attracts attention recently since it is one of possible channels of the black hole formation and is characterized by the unique neutrino signals \\citep{lie04,sum06,sum07}: the short duration and increase of average energies and luminosities in time. This is in contrast to the ordinary supernova neutrinos, whose emissions from a nascent proto-neutron star last for $\\sim$20 s with a gradual decrease of average energies and luminosities \\citep{bur88,suz94}. The black hole formation will be inevitable for massive stars beyond $\\sim$40M$_{\\odot}$, which was estimated with large uncertainties to be the mass threshold for black hole formation by \\citet{fry99}. The dynamical collapse to a black hole is triggered by the intense accretion of matter onto the proto-neutron star. Note that this is different from the so-called delayed black hole formation of the proto-neutron star with a fixed mass \\citep{kei95,pon99,bau96b} owing to the termination of accretion after a successful launch of shock wave. It should be also emphasized that the neutrino signal may be used as a probe into the equation of state of nuclear matter \\citep{sum06} and exotic matter \\citep{nak08} in a different way from the delayed collapse of proto-neutron stars \\citep{pon01a,pon01b}.  This article is the second paper, following the first paper \\citep{sum07}, of serial studies on the black hole formation associated with the short burst of energetic neutrinos. In the first paper, the evolution toward the black hole formation from the gravitational collapse of the 40M$_{\\odot}$ star have been examined by taking account of the dependence of equation of state.  % By solving general relativistic $\\nu$-radiation hydrodynamics, they have followed the hydrodynamics up to the moment of black hole formation and the neutrino distributions during the whole evolution to obtain the detailed features of neutrino emission. They have shown that the black hole formation occurs in $\\sim$1 s after the core bounce via proto-neutron star evolution, but it depends crucially on the equation of state of dense matter. The equation of state (EOS) controls the timing of the re-collapse from the proto-neutron star to the black hole since the EOS determines the maximum mass. The mass of proto-neutron star increases due to the accretion of material in the situation of failed shock wave and the dynamical collapse is triggered when the mass reaches the critical mass. We remark that the progenitor model of 40M$_{\\odot}$ by \\citet{woo95} has been uniquely adopted in their studies and other studies on such massive stars \\citep{lie04}.  %  In the first paper \\citep{sum07}, associated neutrino burst has been shown to be unique to identify the black hole formation and to probe the properties of dense matter.  % The neutrino burst is terminated by the black hole formation and the average energies of neutrinos increase toward the end because of high temperature realized in the collapsing proto-neutron stars. These features are again sensitive to the properties of equation of state, which determines the duration of neutrino burst and the properties of compressed matter at neutrino emitting region. The neutrino burst toward the black hole formation can, in principle, be detectable at the terrestrial neutrino detector facilities % \\citep{ike07} % and can be used as a signature of black hole formation and a probe of dense matter \\citep{sum06}.  % Moreover, if a collapse of massive star occurs in a nearby galaxy, the event may be observed, for example, in the proposed survey of disappearance of massive stars \\citep{koc08} and it will be possible to obtain the information on the progenitor.  % However, one has to provide the detailed features of neutrino burst by taking care of uncertainties of progenitor models in addition to dense matter if one wants to use them as templates for the neutrino search.  In this study, we % make a first attempt to see % whether the features of black hole formation reported in the first paper \\citep{sum07} are taken over to % different progenitors. We adopt % three % stellar models as the initial condition of the numerical simulations to make the comparison with the evolutions by the numerical modeling we have adopted in the previous study \\citep{sum07}. In addition to the stellar model of 40M$_{\\odot}$ star by \\citet{woo95}, we adopt the two stellar models by \\citet{has95} and by \\citet{tom07}. As a case of different mass, we adopt a 50M$_{\\odot}$ star by \\citet{tom07}, who studied the evolution of metal-poor massive stars. We have chosen the model with zero metallicity among their series of different metallicities. This model is the limiting case of metal-poor massive stars without mass loss. As a model for the same mass (40M$_{\\odot}$) but by a different method of stellar evolutionary calculation, we adopt a 40M$_{\\odot}$ star by \\citet{has95}. The comparison between the different models enables us to explore the dependence of numerical results on the uncertainty in the stellar evolution.  The investigation on the progenitor dependence is important for the utilization of neutrino bursts as a probe of equation of state. If the neutrino burst is quite different depending on the progenitors, it becomes difficult to distinguish the differences found by the equation of state. Using the two sets of EOS (LS-EOS and Shen-EOS), it has been shown that the time duration of neutrino bursts differ by a factor of $\\sim$2 depending on the softness of EOS. If the progenitor models give large differences, they may smear out the EOS difference, which we want to use as a probe.  Our aim of the paper is, therefore, to % make a first exploration to % the dependence on the progenitors by adopting % a small number of % stellar models and to show that the dynamics toward the black hole formation is common in the massive stars having a large central iron core. We aim also to demonstrate that the neutrino signal has similar features once we set the equation of state.  % We analyze how the evolution is related with the density profile of progenitor, which determines the accretion rate. We discuss the characteristics of neutrino bursts obtained by different progenitors with the sets of equation of state for possible detection in future.  We plan this paper as follows. We briefly explain the method of numerical simulations in \\S \\ref{sec:simulations}. We compare the progenitors adopted for initial models and discuss the differences in \\S \\ref{sec:initial}. We report the numerical results on the dynamics in \\S\\S \\ref{sec:50M} and \\ref{sec:40M} for two models through the comparison with the previous result. We describe the accretion rates and the profiles during the evolution in \\S\\S \\ref{sec:accretion} and \\ref{sec:profile} to explain the differences among the models. We discuss the characteristics of neutrino signals in the models to distinguish the progenitor and the equation of state in \\S \\ref{sec:nu-signal}. Implications of this work and summary will be made in \\S\\S \\ref{sec:discussions} and \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussions}  In this section, we discuss the perspective of the current work regarding the progenitor and the equation of state. We also remark on the observational chances of neutrino burst and implications to astrophysical phenomena with the black hole.  The present study on the black hole formation and the associated neutrino burst relies on the presupernova configurations of massive stars. Further studies with updated models of progenitors are definitely warranted. Necessary condition to have this type of phenomena is the appropriate density distribution to cause the intense accretion. In the three models (T50, H40 and W40) adopted in the current study, the size of iron core is large ($\\sim$2M$_{\\odot}$) and the dense region extends up to M$_b$ $\\sim$3M$_{\\odot}$. % This is different from the case of a 15M$_{\\odot}$ star by \\citet{woo95}, in which the iron core is smaller and the density slope is steep already at M$_b$ $\\sim$1.5M$_{\\odot}$, for example. It is necessary to study further the density profiles of various progenitors and its relation to the accretion in order to examine whether the size of iron core is essential. It may be enough to have the flat profile of density so that the density remains high ($\\sim$10$^{6}$ g/cm$^{3}$) at M$_b$=2--3M$_{\\odot}$, which corresponds to the critical mass of proto-neutron stars. If the density profile is steep and the density is too low at the corresponding region in 40--50M$_{\\odot}$ stars, the black hole formation may take much longer ($\\sim$10--100 s) than the current case. Such events with longer neutrino bursts may look similar to the ordinary supernovae of $\\sim$15M$_{\\odot}$ stars in terms of neutrino signals.  Regarding the uncertainty of progenitors, one should pay attention to the mass loss as well as the rotation. The progenitor models in \\citet{woo95} and \\citet{has95} are obtained by the presupernova evolution without the mass loss. The inclusion of the mass loss leads to the smaller mass at the precollapse stage even from the large initial mass (ex. 50M$_{\\odot}$) at the main sequence. The iron core mass and the density distribution may end up to be similar to those in the progenitors of about 20M$_{\\odot}$.  % In this regard, metal poor massive stars are more likely to be the progenitors for the black hole formation as studied in the current work. In fact, T50 adopted in the current study is the one with zero metallicity that forbids the mass loss. Its large central core results in the necessary intense accretion, which in turn leads to the quick black hole formation with short neutrino bursts. Whether the change of evolution toward the solar metallicity may cause the suppression of the current scenario remains to be seen along with the systematic studies on the metallicity dependence of progenitors and the careful studies on the mass loss.  We study the black hole formation under the spherical symmetry. This geometrical assumption limits ourselves to the non-rotating massive stars among varieties in rotation rates. One should note, however, that the spherical symmetry is a good approximation if the rotation is not significantly fast.  % It is challenging to study the collapse of rapidly rotating massive stars, which are possible progenitors for hypernovae and gamma ray bursts. One has to perform the fully general relativistic calculations in multi-dimensions to follow the evolution toward the black hole formation. Such challenges are made in the numerical simulations of general relativistic hydrodynamics but with simple treatment of neutrino transfer, microphysics (ex. EOS) and progenitors \\citep{shi05,sek05,sek07}. Multi-dimensional simulations with proper treatment of neutrino transfer are awaited for to predict the features of neutrino emission.  The rotating collapse to the black hole and the following evolution is also interesting from the aspect of nucleosynthesis. In the context of collapsar scenario, the ejection of neutron-rich material in the jet from the surroundings of black hole and the neutrino-driven winds from the accretion disk around the black hole are proposed to be the possible site for the heavy elements production including r-process elements \\citep{sur06,fuj07}. Although we performed the numerical simulations in the limit of no rotation, one may use the current results to guess the time duration till the black hole formation and the accretion rate of material, which are influential in the scenario of nucleosynthesis. The time duration obtained in the current work is a lower limit since the maximum mass of rotating proto-neutron stars is larger than the non-rotating ones \\citep{sum99} and the rotation supports accreting material away due to the centrifugal force. The accretion rate in the current study is an upper limit, accordingly. Large accretion rates in the disk formed around black hole, which are assumed in the studies of nucleosynthesis of r-process \\citep{sur06}, may be limited from the density profile of progenitors as we have pointed out.  The observational possibility to detect the neutrino bursts toward the black hole formation from massive stars deserves to be mentioned. Occurrence rate of the phenomena (non-rotating black hole formation and its short neutrino burst) as in the current scenario is proportional to the fraction of massive stars having large masses for intense accretion without significant rotation. This depends on the initial mass function and rotation rates with large uncertainties, but such stars comprise a certain fraction % ($\\sim$30$\\%$) of massive stars, which can be estimated from the Salpeter's initial mass function with an assumption on the mass threshold ($\\sim 25$M$_{\\odot}$) for the black hole formation. \\citet{koc08} estimated from the observations done so far that the rate of black hole formation could be comparable to the rate of normal core-collapse supernovae.  % Observational analyses of explosion energies and Ni productions for various supernovae indicate that there % have been observed at least several % faint supernovae (SN~1997D and SN~1999br, for example) from massive stars beyond 20M$_{\\odot}$ without significant rotations, which may be akin to the event considered in this paper \\citep{nom07}.  %  The planned survey of disappearance of massive stars within a distance of 10 Mpc \\citep{koc08} will find both dim supernovae and collapses with no optical display that lead to the black hole formation. If they occur in the Galaxy or nearby galaxies, the information on progenitor will be obtained by the optical records before it disappeared. In addition, thousands of neutrinos will be detected at the terrestrial facilities. Neutrino data combined with optical ones will provide us with the information on the dynamics of black hole formation and constrain the properties of dense matter and the progenitor.  % In the search of supernova neutrinos by the Super-Kamiokande detector, neither neutrinos associated with black hole formation nor ordinary supernova neutrinos are detected during the period of their survey unfortunately \\citep{ike07}.  %  The numerical result on the massive star with zero metallicity (T50) suggests that the black hole formation in the current scenario may occur for Pop III stars in the early Universe (see also \\citet{nak06,nak07}). Since the massive stars become more popular in the metal poor epoch, the neutrinos from black hole formation in the early stage of galaxies may contribute to the relic signal as back ground neutrinos. Although the signal is short, higher energies and luminosities may help to enhance this contribution and they may not be negligible. This argument again depends on the fraction of massive stars resulting the current scenario through the evolution of galaxies.  Before we close the discussion, we comment on the choice of equation of state. Although density and temperature become extremely high enough to have new degrees of freedom such as hyperons and quarks in the evolution toward the black hole formation, we adopt the two sets of EOS within the nucleonic degree of freedom in the current study to assess the progenitor dependence. Studies on the modifications of the evolution of black hole formation and the neutrino signal due to hyperon mixtures and the appearance of quark phase are now under way and will be reported in separate papers % \\citep{nak08}.  %    We study the black hole formation and the associated neutrino emission from the gravitational collapse of massive stars to clarify the dependence on the progenitor models following the basic study of the phenomena in the first paper \\citep{sum07} of the series. We perform the sets of numerical simulations of general relativistic $\\nu$-radiation hydrodynamics under the spherical symmetry to follow the collapse of massive stars of 40M$_{\\odot}$ and 50M$_{\\odot}$ adopting different models of the presupernova evolution. We clarify that the characteristics of evolution of proto-neutron stars toward the black hole formation and % the associated short burst of neutrinos are found similar in the numerical simulations for the three progenitor models.  %  We find that the different progenitor models may result in different rates of matter accretion onto the proto-neutron star born in the collapse of massive stars. The different accretion rates can lead to different histories of increasing mass of proto-neutron stars toward the critical mass and to different durations till the black hole formation.  % In order to extract the properties of equation of state from the neutrino burst toward the black hole formation from massive stars, as proposed in the previous studies \\citep{sum06,sum07}, it is necessary to constrain the uncertainty from progenitor models. In the comparison of the three progenitors we studied, the change of time duration is found not so large as the difference originating from the two sets of equation of state. The change of time profiles of average energies and luminosities of neutrinos due to progenitors is found to be minor as well. Note, however, that the present study is a first trial exploration and more systematic investigations are obviously necessary to see if our conclusion is generic or not.  %  We demonstrate that there is a relation between the density profile of progenitor and the accretion history to cause the difference in black hole formation. A higher density at M$_{b}$=2.0--3.0M$_{\\odot}$ in a progenitor results in a shorter free fall time and a higher accretion rate. A flat density profile of core in massive stars is a key factor to cause the rapid formation of black hole formation in $\\sim$1 s together with the characteristic neutrino burst.  % When the density gradients are different as in the comparison between W40L and H40L, the density profiles may be distinguished from the detailed information of neutrino detections.  % This relation of the density profile of progenitors to the accretion rates is helpful to discuss the evolution after the collapse of massive stars in other stellar models and may be useful to infer the accretion dynamics of rotating collapse to the black hole formation and associated explosive phenomena such as collapsars. Further studies on rotating collapse of massive stars with intense accretion are called for to clarify whether the outcome may lead to possible success of modeling gamma ray burst and/or nucleosynthesis of heavy elements."
    },
    "0808/0808.0673_arXiv.txt": {
        "abstract": "\\noindent We show that models of new particle physics containing massless pseudoscalar fields super-weakly coupled to photons can be very efficiently probed with CMB polarization anisotropies. The stochastic pseudoscalar fluctuations generated during inflation provide a mechanism for converting $E$-mode polarization to $B$-mode during photon propagation from the surface of last scattering. The efficiency of this conversion process is controlled by the dimensionless ratio $H/(2\\pi f_a)$, where $H$ is the Hubble scale during inflation, and $f_a^{-1}$ is the strength of the pseudoscalar coupling to photons. The current observational limits on the $B$-mode constrain this ratio to be less than 0.07, which in many models of inflation translates to a sensitivity to values of  $f_a$ in excess of $10^{14}$ GeV, surpassing the sensitivity of other tests. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0808/0808.0735_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\noindent The neutrino \\cite{hh00} is an elementary particle that scatters only through the weak interaction, and consequently rarely interacts in matter.  Neutrinos are neutral, carry spin-1/2, and are members of the family of elementary particles called leptons.  Thus they differ from the quarks of the standard model (spin-1/2 particles which participate in strong and electromagnetic interactions) and from the other leptons, which are charged and thus interact electromagnetically.   Neutrinos and their antiparticles come in three types $-$ or flavors $-$ labeled according to the charged partners, the electron, muon, or tauon, that accompany neutrino production in charge-changing weak interactions.  The most familiar such reaction is $\\beta$ decay \\[ (N,Z) \\rightarrow (N-1,Z+1) +  e^- + \\bar{\\nu}_e \\] in which a nucleus containing N neutrons and Z protons decays to a lighter nucleus by converting a neutron to a proton, with the emission of an electron and an electron antineutrino.  Indeed, it was the apparent absence of energy conservation in nuclear $\\beta$ decay that first lead Wolfgang Pauli, nearly 80 years ago, to suggest that some undetected neutral particle (the $\\nu_e$) must be escaping from nuclear $\\beta$ decay experiments.  Neutrinos play a very special role in astrophysics \\cite{bahcallbook}.  First, they are the direct byproducts of the nuclear reaction chains by which stars generate energy: each solar conversion of four protons into helium produces two neutrinos, for a total of $\\sim$ 2 $\\times$ 10$^{38}$ neutrinos each second.  The resulting flux is observable on Earth.  These neutrinos carry information about conditions deep in the solar core, as they typically leave the Sun without further interacting.   They also provide experimentalists with opportunities for testing the properties of neutrinos over long distances.  Second, they are produced in nature's most violent explosions, including the Big Bang, core-collapse supernovae, and the accretion disks encircling supermassive black holes.  Recent discoveries of neutrino mass show that neutrinos comprise a small portion of the ``dark matter\" that influences how large-scale structure -- the pattern of voids and galaxies mapped by astronomers -- formed over cosmological times.  Third, they dominate the cooling of many astrophysical objects, including young neutron stars and the degenerate helium cores of red giants.  Neutrinos can be radiated from deep within such bodies, in contrast to photons, which are trapped within stars, diffusing outward only slowly.  Finally, neutrinos are produced in our atmosphere and elsewhere as secondary byproducts of cosmic-ray collisions.  Diagnosis of these neutrinos can help constrain properties of the primary cosmic ray spectrum.  Neutrinos also mediate important astrophysical processes.  The supernova ``neutrino wind,\" which blows off the surface of the proto-neutron star just seconds after core collapse, is thought to produce conditions favorable for the synthesis of about half of the neutron-rich nuclear species heavier than iron, through rapid neutron capture (the r-process).  Neutrinos also directly synthesize certain rare nuclei, such as ${}^{19}$F.  Nuclear and particle physicists are exploiting astrophysical neutrino fluxes to do important tests of the standard model of particle physics.  These tests include neutrino oscillations (the process by which a massive neutrino can be produced in one flavor state but detected later as a neutrino with a different flavor), neutrino decay, neutrino mass effects, and searches for nonzero neutrino electromagnetic moments. ",
        "conclusions": ""
    },
    "0808/0808.2506_arXiv.txt": {
        "abstract": "The discovery that the short-lived radionucleide \\fe60\\ was present in the oldest meteorites suggests that the formation of the Earth closely followed the death of a massive star. I discuss three astrophysical origins: winds from an AGB star, injection of supernova ejecta into circumstellar disks, and induced star formation on the boundaries of HII regions. I show that the first two fail to match the solar system \\fe60\\ abundance in the vast majority of star forming systems. The cores and pillars on the edges of HII regions are spectacular but rare sites of star formation and larger clumps with masses $10^{3-4}$\\,\\Msun\\ at tens of parsec from a supernova are a more likely birth environment for our Sun. I also examine $\\gamma$-ray observations of \\fe60\\ decay and show that the Galactic background could account for the low end of the range of meteoritic measurements if the massive star formation rate was at least a factor of 2 higher 4.6\\,Gyr ago. ",
        "introduction": "The study of planet formation can be approached from two sides: the large scale encompassing the molecular core that collapses to a star and surrounding planetary disk, and the small scale in which the planets and the remnants of their formation, meteorites, are examined in detail to learn about the conditions of the early solar system (ESS).  Astronomers have identified each of the main stages by which the interstellar medium (ISM) becomes molecular, fragments, and individual cores collapse to protostars. We have imaged the disks of planet forming material around young stars and witnessed the debris from planetesimal collisions. Although many details remain to be worked out, the basic properties of a proto-planetary disk, such as mass, size, composition, and lifetime, are well characterized. We have also identified over 200 extrasolar planetary systems around nearby stars. The surprising diversity of these systems, however, raises the vexing issue of how typical is our solar system.  Cosmochemists have identified the oldest rocks in the solar system and determined their age, 4.567\\,Gyr, to astonishing precision. Careful study of their mineralogy shows the detailed conditions, including, for example, thermal history, radiation field, and transport processes, in the ESS. As with cosmologists studying the Universe, however, there is one and only one system at hand which raises the vexing issue of how typical is our solar system.  The areas where these two different lines of inquiry, telescopic and microscopic, overlap provide interesting comparisons. Short lived radionucleides (SLR), with half-lives less than 3\\,Myr (see below), are of particular interest for understanding planet formation and the astrophysical environment of the ESS. In this chapter, I focus on the puzzle presented by the presence of \\fe60\\ in the most primitive meteorites. As I show, it may be hard to reconcile the astronomical and cosmochemical pictures but it is important to try since it may show the limitations in our knowledge or understanding and it may also help show whether our solar system is indeed typical, or not.  I begin by reviewing the astronomical and cosmochemical background to this subject. These sections are brief as they are covered in other chapters in this book. I discuss three different scenarios which have been proposed for the delivery of SLR into the ESS and assess the probability for incorporation of \\fe60. I conclude that the only viable mechanism that might apply to a large number of planetary systems is the rapid collapse of large cluster forming clumps neighboring a massive star forming region. Finally I return to a basic assumption that the \\fe60\\ abundance is greater than the Galactic background and show a discrepancy between the predicted level and recent $\\gamma$-ray observations. ",
        "conclusions": "I have critically evaluated several proposed mechanisms for the delivery of \\fe60\\ to the ESS and found that most are highly improbable. AGB stars have moved away from their birthsite and any cross-pollination of a molecular core or circumstellar disk would be serendipitous. Further, their output is too small to affect a significant fraction of the ISM. No more than 0.1\\%, and probably far less, of stars could receive \\fe60\\ in this way.  Supernova may inject SLR directly into disks but the range of distances over which disks both survive the blast and are enriched to ESS levels is very small. Further, few disks remain by the time of the supernova for all but the most massive, shortest-lived progenitors. In this case, I estimate that no more than about 1\\% of all protoplanetary systems could receive \\fe60\\ in this way.  I was unable to assign probabilities to induced star formation scenarios since the statistics are not well known. The population of young stars on the boundary of HII regions appears to be only a small fraction of that within the central cluster and the number that form deeper in the surrounding molecular cloud. HII regions expand rapidly and are tens of parsecs in size before any supernova explosion. Matching the required high levels of \\fe60\\ places individual star forming cores well within the boundary of mature HII regions and they would be shredded by stellar winds. Even if a core were to collapse and form a star, the surrounding envelope would be photoablated or swept away by the time of a supernova which would impact, at most, a circumstellar disk. This scenario is effectively the same as the direct disk injection hypthesis since that used a prescription for disk evolution that included observations of clusters with massive stars.  Larger, cluster forming clumps in the molecular cloud beyond the HII region can capture enough supernova material and collapse fast enough to deliver ESS levels of \\fe60\\ to planetary systems. The large scale ``collect-and-collapse'' scenario for molecular clumps neighboring a massive star forming region may be the most promising solution to the \\fe60\\ puzzle. Molecular clouds are found around many supernova remnants and, although they may be are strongly disrupted, star formation can occur in the surviving gas (Huang \\& Thaddeus \\cite{Huang86}; Reach \\& Rho \\cite{Reach99}; Reynoso \\& Mangum \\cite{Reynoso01}). Perhaps the Orion subgroups are an example where this happened in the past and the Rosette cloud an example where this will happen in the future. However, only through detailed surveys of the molecular and protostellar content in young cloud-supernova interactions that quantify the amount and efficiency of star formation can we assess the likelihood that the Sun was born in such an environment.  The discrepancy between the $\\gamma$-ray measurements of the \\al26\\ and \\fe60\\ background with model calculations needs to be understood. Also, more sensitive cosmochemical measurements will reduce the uncertainty on the initial abundance of \\fe60\\ and constrain the timing of its injection (Bizzarro et al. \\cite{Bizzarro07}, Dauphas et al. \\cite{Dauphas08}). If the high end of of the current range is confirmed, the background can be ruled out as the source although it may be a significant component that reduces the amount required from a discrete source.  Finally, if there really is no common scenario that can deliver large amounts of \\fe60\\ from a massive stellar core to a protoplanetary system, we are faced with the prospect that our solar system may be a one-in-a-hundred rarity. Astronomers have a long history of rejecting the idea that we are special but what are the implications in this case?  It seems reasonable to assume that however \\fe60\\ was delivered to the ESS, many other SLR were incorporated in the same manner. But if this was an unusual occurrence, then most planetary systems have lower levels not just of \\fe60\\ but of other SLR, particularly \\al26. The latter is the dominant heating source for planetesimals at early times (Hevey \\& Sanders \\cite{Hevey06}) and a reduced abundance would imply a different thermal history and, potentially, result in a higher water content of terrestrial planets (Desch \\& Leshin \\cite{Desch04}; Gaidos, Raymond, \\& Williams \\cite{Gaidos08})."
    },
    "0808/0808.4099_arXiv.txt": {
        "abstract": "Using a hydrodynamics plus $N$-body simulation of galaxy cluster formation within a large volume and mock Chandra X-ray observations, we study the form and evolution of the intrinsic scatter about the best-fit X-ray temperature-mass relation for clusters. We investigate the physical origin of the scatter by correlating it with quantities that are closely related to clusters' formation and merging histories. We also examine the distribution of the scatter for merging and nonmerging populations, identified  using halo merger trees derived from the simulation as well as X-ray substructure measures. We find a strong correlation between the scatter in the $M-T_X$ relation and the halo concentration, in the sense that more concentrated clusters tend to be cooler than clusters with similar masses. No bias is found between the merging and relaxed clusters, but merging clusters generally have greater scatter, which is related to the properties of the distribution of halo concentrations. We also detect a signature of non-lognormality in the distribution of scatter for our simulated clusters both at $z=0$ and at $z=1$. A detailed comparison of merging clusters identified by substructure measures and by halo merger trees is given in the discussion. We conclude that, when cooling-related effects are neglected, the variation in halo concentrations is a more important factor for driving the intrinsic scatter in the $M-T_X$ relation, while departures from hydrostatic equilibrium due to cluster mergers have a minor effect. ",
        "introduction": "Galaxy clusters are potentially valuable cosmological probes because of their unique position in the hierarchy of structure formation. They are the largest gravitationally bound objects, having just separated from the cosmic expansion and collapsed from density fluctuations in the past few billion years. Therefore, statistical measures of clusters, such as their mass distribution as a function of redshift, are sensitive to the cosmic matter density parameter $\\Omega_m$, the dark energy density parameter $\\Omega_{de}$, the normalization of the primordial fluctuation spectrum $\\sigma_8$, and the dark energy equation of state parameter $w$ \\citep{2005astro.ph..7013H}. Future cluster surveys will yield cosmological parameter constraints that are complementary to those from upcoming microwave background probes (e.g.\\ Planck) and Type~Ia supernova observations.  However, clusters present us with a dilemma: their masses are well-predicted by numerical simulations, but in the real world, 80--85\\% of their mass is in the form of invisible dark matter. In order to make contact with observations, observational mass proxies are needed. Fortunately, cluster masses correlate with many observable quantities, such as X-ray temperature $T_X$, X-ray luminosity $L_X$, optical richness, infrared luminosity, and the Sunyaev-Zel'dovich effect \\citep{1999ApJ...517..627M,2003ApJ...591..749L,2005A&A...433..431P,2005ApJ...633..122H,2006ApJ...648..956S}. These mass-observable relations indicate that clusters are fairly regular and close to equilibrium, despite having dynamical timescales of order 1/10 the age of the universe.  For clusters to provide meaningful constraints on the cosmological parameters, the systematic errors in mass estimates based on these relations must be well-understood. For example, masses determined under the assumption of hydrostatic equilibrium in general underestimate the spherical overdensity mass $M_{200}$ by $\\sim20\\%$ \\citep{2007ApJ...655...98N}. The X-ray temperature determined through spectral fitting is also known to bias low with respect to the emission-weighted temperature commonly used in numerical simulations \\citep{2005ApJ...618L...1R}. The discrepancy in the normalization of the $M$--$T_X$ relation between observations and numerical simulations has been alleviated only recently by taking these two effects into account. This emphasizes the importance of using mock-observation tools to make direct comparisons between results from numerical simulations and observational data.  The origin and distribution of scatter in these relations are also important. Due to the exponential shape of the mass function, scatter in the $M$--$X$ relation (for observable $X$) boosts the number density of clusters observed in logarithmic bins of $X$, as the overall number of lower-mass clusters scattering to higher values of $X$ far exceeds the number of high-mass clusters scattering in the opposite direction. Underestimating this scatter can lead to an overestimate of $\\sigma_8$, for instance \\citep{2002ApJ...577..579R}. Attempts have been made to reduce the observed scatter to get better constraints on cluster masses. Such attempts include the use of core-excised quantities to reduce the large scatter in the $L_X$--$T_X$ relation due to the effects of cool cores \\citep{1998MNRAS.297L..57A,2006ApJ...639...64O} and the invention of the X-ray counterpart of the Compton y-parameter, $Y_X \\equiv M_{\\rm gas}T_X$, to obtain a very tight $M$--$Y_X$ correlation \\citep{2006ApJ...650..128K}. These successful examples illustrate the possibility of obtaining better mass estimates if our knowledge of the physical origin of scatter is improved. It is also possible to self-calibrate cluster surveys \\citep{2002ApJ...577..569L,2004ApJ...613...41M,2004PhRvD..70d3504L}, fitting the mass-observable relation as an unknown together with the cosmological parameters. However, this technique requires assumptions about the functional form and the mass and redshift dependence of the scatter, and errors in these assumptions can lead to misinterpretation of the obtained constraints \\citep{2005PhRvD..72d3006L}.  In this paper we begin a systematic study of the influence of internal cluster physics on the scatter in cluster mass-observable relations. We focus initially on X-ray observables, as they are less directly dependent on the uncertain details of galaxy formation than, for example, optical observables. Our aim in this paper is to examine the effects of mergers and dynamical state in isolation from effects due to radiative cooling, feedback due to stars and black holes, and diffusive transport. Future papers will examine these other effects separately. Accordingly, the simulation described here includes only dark matter and gasdynamics. While we examine substructure-based measures of dynamical state to make contact with previous work, we extend this work by considering direct measurements of dynamical state that may not be observable but that can give us physical insight into the origin and form of the scatter. For example, we use halo merger histories derived from halo catalogs created every $100 h^{-1}$\\ Myr to identify which clusters are merging at any given epoch. \\S~\\ref{Sec:simulations} describes our numerical methods and simulation parameters. In \\S~\\ref{Sec:analysis} we describe our merger tree and virial analysis procedures, our method for generating simulated X-ray observations, and the substructure measures we employ. We present our results in \\S~\\ref{Sec:results} and discuss them in \\S~\\ref{Sec:discussion}. Finally, we summarize our conclusions in \\S~\\ref{Sec:conclusion}.  Throughout this paper we have taken the Hubble constant to be $H_0 = 100h$~km~s$^{-1}$~Mpc$^{-1}$, with $h = 0.708$.  When quoting masses or radii defined using an overdensity criterion (e.g., $M_{200}$, $R_{200}$), we refer to overdensities relative to the critical density at the relevant epoch. ",
        "conclusions": ""
    },
    "0808/0808.3393_arXiv.txt": {
        "abstract": "We have observed a large sample of spectroscopic binary stars in the Hyades Cluster, using high resolution infrared spectroscopy to detect low mass companions.  We combine our double-lined infrared measurements with well constrained orbital parameters from visible light single-lined observations to derive dynamical mass ratios. Using these results, along with photometry and theoretical mass-luminosity relationships, we estimate the masses of the individual components in our binaries.  In this paper we present double-lined solutions for 25 binaries in our sample, with mass ratios from $\\sim0.1-0.8$.  This corresponds to secondary masses as small as $\\sim0.15\\msun$.  We include here our preliminary detection of the companion to vB~142, with a very small mass ratio of $\\q=0.06\\pm0.04$; this indicates that the companion may be a brown dwarf.  This paper is an initial step in a program to produce distributions of mass ratio and secondary mass for Hyades cluster binaries with a wide range of periods, in order to better understand binary star formation.  As such, our emphasis is on measuring these distributions, not on measuring precise orbital parameters for individual binaries. ",
        "introduction": "}  Observations of spectroscopic binary stars provide dynamical measurements of stellar mass and binary mass ratio that are important inputs to theories of binary star formation \\citep{bonnell2003, tohline2002, clarke2001}.  Measuring the radial velocity versus orbital phase for a single-lined spectroscopic binary (SB1) yields orbital parameters and the mass function, \\fmass, in which the mass ratio, $\\q=\\Mtwo/\\Mone$, is inseparable from the orbital inclination. For an ensemble of SB1s, a statistical approach can be used to derive the distribution of mass ratios \\citep[e.g.,][]{mazeh1992}.  If, however, a binary is observed as a double-lined system (SB2), then the dynamical mass ratio follows directly from the radial velocity measurements \\citep{mazeh2002}.  With an estimate of the primary mass, say based on the spectral type, the mass ratio gives the secondary mass, and in the case of an ensemble of SB2s, the secondary mass distribution.  Obtaining a purely dynamical mass ratio distribution using SB2s is challenging: a binary with a mass ratio of much less than one has a small flux ratio, $\\alpha=\\mathit{F}_2/\\mathit{F}_1$, so the secondary component is difficult to detect.  This flux ratio problem is particularly inhibitive at visible wavelengths where much of the long term monitoring of spectroscopic binaries has occurred.  For this reason, most identified SBs are SB1s, and there is a strong selection effect favoring the detection of SB2s with \\q{} near 1. A binary composed of main sequence stars with unequal masses will have an $\\alpha$ that increases towards longer wavelengths, making the companion easier to detect using infrared spectroscopy.  Such observations have allowed the measurement of SB2s with \\q's as small as $\\sim0.1-0.2$ \\citep[e.g.,][]{mazeh2002,prato2002}.  \\citet{duquennoy1991} examined a sample of 164 nearby SB1s, SB2s, visual binaries, and common proper motion pairs, with F- and G-type primary stars, and found that the mass ratio distribution for medium to long period binaries increases towards small mass ratios. \\citet{goldberg2003} used a well defined sample of 129 SB1s and SB2s from the Carney-Latham sample of high proper motion stars \\citep{carney1994} to derive a bimodal mass ratio distribution with peaks at $\\q\\sim0.2$ and $\\q\\sim0.8$.  They found somewhat different distributions for halo versus disk stars, and for binaries with primary masses greater than and less than 0.67\\msun.  By working in the infrared, \\citet{mazeh2003} observed 32 binaries from the Carney-Latham sample as SB2s, with mass ratios as small as $\\sim$0.20. When combined with SB2s from visible light spectroscopy, they found a mass ratio distribution that is approximately flat from $\\q\\sim0.3 - 1.0$.  The samples studied by \\citet{goldberg2003}, \\citet{mazeh2003}, and, for \\q's near 1 by \\citet{lucy2006}, represent averages of star formation outcomes in the solar neighborhood over billions of years. Most star formation occurs in localized regions within molecular clouds.  Studies of the binary population in diverse star forming regions can isolate the physical parameters that determine the properties of the binaries formed.  For example, the competing effects of fragmentation, dynamical interactions, and accretion are not fully understood \\citep{goodwin2007, ballesteros2007, whitworth2007}. Observations of binaries in open clusters can clarify the important processes because they share a common star formation origin and history.  Ideally, such observations would be carried out on young clusters that have an easily identifiable membership and still retain their most massive members.  However, it is difficult to measure the orbital velocities of the youngest stars because they tend to have high rotation velocities.  Stars in older open clusters still retain much of their cluster identity yet are rotationally spun down.  The Hyades is one such nearby open cluster: it is well studied and presents an excellent laboratory for measuring a well defined sample of binary stars in order to investigate binary star formation.  The cluster has an age of $\\sim650$\\,Myr \\citep{lebreton2001}, old enough that stars later than spectral type F3 have slowed to $v\\sin i \\la 25\\kps$ \\citep{kraft1965}, and G-type and later to $v\\sin i\\la10\\kps$ \\citep{stauffer1987}.  Most mid K-type and earlier cluster members have individual \\emph{Hipparcos} distances, and with a mean cluster distance of $\\sim46$\\,pc \\citep{perryman1998} the known binaries are bright.  Additionally, with metallicity [Fe/H]=0.14 \\citep{lebreton2001} the spectral lines of Hyades members are deep and well suited to spectroscopic analysis.  Radial velocity surveys of the Hyades \\citep[e.g.,][]{wilson1948, kraft1965, detweiler1984, stefanik1985, griffin1978,griffin1981,griffin1985, griffin1988} have identified a sufficient number of SBs to facilitate a statistical analysis of their physical properties. \\citet{patience1998} carried out a study of Hyades ``visual binaries'' using speckle interferometry.  These wide binaries have periods longer than most SBs and several are the wide components of hierarchical triples with known SBs.  From their observations, \\citet{patience1998} derived a photometric mass ratio distribution for long period systems that appears to increase towards small \\q.  At the Harvard-Smithsonian Center for Astrophysics (CfA), Robert P. Stefanik (RPS) and David W. Latham (DWL) have monitored Hyades members for SBs by visible light spectroscopy since 1979 \\citep[e.g.,][]{latham1982,stefanik1985}.  Their observations have refined the cluster membership and yielded precise parameters for many SB1s and SB2s.  In 2003, the current authors began a collaboration with the CfA group to extend, through infrared observations, the detection of binary companions to smaller masses than possible by visible light spectroscopy alone.  To this end, RPS and DWL made available to us their unpublished parameters for many Hyades SB1s.  In this paper we present infrared SB2 detections that span the range of mass ratios from $\\q\\sim0.1$ to $\\q\\sim0.9$.  We verified these results using the sample of Hyades SB1s available in the literature \\citep[e.g.,][]{sanford1921,griffin1978,griffin1981,griffin1985}.  When combined with available surveys (both present and future) of Hyades binaries, our SB2 detections will enable better determinations of the distribution of secondary masses.  In \\S\\ref{sample}, we present our Hyades binary sample; \\S\\ref{infrared} describes our infrared observations and velocity measurement techniques; \\S\\ref{sb2solutions} includes the SB1 and SB2 solutions for our sample binaries; \\S\\ref{discussion} discusses the quality of our SB2 solutions and derives the secondary masses; and \\S\\ref{summary} presents a brief summary and comments on our future papers. ",
        "conclusions": "}  \\subsection{\\sc Validation of SB2 Solutions\\label{sb2validation}}  Four of the SB2 solutions shown in Figures~\\ref{fig:sb2plots1} and \\ref{fig:sb2plots2} depend on a single measurement of the secondary velocity, and another nine on only two measurements.  Consequently, many of the derived \\Ktwo's, and the resulting \\q's, have large uncertainties.  Figure~\\ref{fig:qvsqmin} plots \\qmin{} against the measured \\q's to evaluate the plausibility of the measured values.  As expected, \\q{} is greater than or within $1\\sigma$ of \\qmin{} for all of the binaries, except vB~59 and H382 which are consistent with \\qmin to better than $2\\sigma$.  Both vB~59 and H382 have long periods and were observed at orbital phases where the velocity separation was small, resulting in low precision measurements of \\Ktwo.  However, we have two observations of each system at slightly different phases and the measured secondary velocities are consistent.  Therefore, we are confident that we are detecting the secondary in both of these systems.  Observations of vB~59 in $\\sim2012$ and of H382 as soon as $\\sim2009$, when the orbital phases will have changed significantly, would further improve the SB2 solutions for these systems.  The flux ratio, $\\alpha$, measured by our cross-correlation routine offers an additional validation of each SB2 solution. Figure~\\ref{fig:qvsalpha} plots the $\\alpha$ measured for each binary, averaged over all observations, against the measured mass ratio.  Also shown are theoretical \\emph{H}-band curves calculated from BCAH for binaries with primary masses from 0.6\\msun{} to 1.2\\msun.  Two factors complicate a direct comparison between our measured flux ratios and the theoretical values.  First, the small wavelength range of CSHELL spectra severely limits our ability to measure precise flux ratios because individual spectral lines vary only a small amount relative to each other with changing spectral type.  Considerably more accurate flux ratios can be measured with spectra covering a larger wavelength range, and thereby having many more spectral lines with differing dependencies on spectral type.  Second, the theoretical curves represent the integrated flux over the entire \\emph{H}-band, of which our CSHELL spectra only sample $\\sim1.5\\%$.  Nonetheless, our measured values are well grouped along the curves.  The two obvious outliers are L79 and vB~142.  We have only a single, low S/N observation of the long period binary L79.  Our measured flux ratio is poorly constrained, varying significantly depending on the exact pair of templates used; the uncertainty shown in the figure probably underestimates the actual uncertainty.  However, our measured secondary velocity is well constrained, independent of the template pair, and we are confident that the secondary detection is real. vB~142 has a very small mass ratio, making the measurement of the companion particularly difficult; \\S\\ref{vb142} addresses this system in detail.   \\subsection{\\sc Infrared Non-Detections\\label{sb2nondetections}}  Our infrared observations failed to detect the secondary in seven systems: vB~8, vB~30, vB~39, H411, L77, L90, and vB~114. Figure~\\ref{fig:qminhist} shows the distribution of \\qmin{} for the infrared sample binaries; the systems not detected as SB2s are indicated by the hashed region.  We expected that our sensitivity to binary companions would be incomplete for systems with the smallest mass ratios because these systems also have small flux ratios. However, several of the systems for which we did not detect the secondary have large \\qmin{}.  For three additional systems, vB~43, L57, and H509, we did not detect the secondary in a subset of our observations.  We address each of these ten systems below.  vB~8 and vB~30 have rapidly rotating F-type primary stars and both have small \\qmin{}, 0.15 and 0.19, respectively.  We observed both spectra in the infrared at multiple epochs, with orbital phases where the predicted velocity separation was large.  We propose two possible explanations for our failure to detected these companions.  The true value of \\q{} may actually be close to \\qmin.  For a given mass ratio, the \\emph{H}-band flux ratio of a binary decreases as the mass of the primary increases.  Because vB~8 and vB~30 have primaries more massive than those of a typical binary in our sample, their flux ratio's may be too small to detect the secondaries at the S/N of our spectra. Alternatively, if the companions are also rotating rapidly, their spectral lines would be too broad to measure with CSHELL's limited wavelength coverage, even with high S/N observations.  vB~114 and vB~39 have orbital periods of 4578 days and 5083 days, respectively.  Our observations of both systems occurred at orbital phases such that the primary velocities were near \\g, and the predicted velocity separations were only a few \\kps.  Velocity measurements under such conditions are inherently difficult, and even had we detected the secondary components, the resulting \\Ktwo's would be poorly constrained.  vB~114 will have a more favorable orbital phase in $\\sim$2010; the phase of vB~39 will not improve until $\\sim$2013.  Both have moderate \\qmin{}, 0.44 for vB~39 and 0.29 for vB~114, so the flux ratios should not impede in detecting these secondaries.  H411, L77, and L90 fall towards the faint edge of the distribution shown in Fig.~\\ref{fig:hmag}, and our observations of these targets have S/N insufficient to detect their secondaries.  H411 has a small \\qmin, $\\sim0.18$, which corresponds to a minimum H-band flux ratio of only a few percent, and is at the limit of our best observation of this binary, with S/N$\\sim$50.  Because H411 is faint, obtaining a spectrum with CSHELL that has better S/N would require many hours of integration; such an observation could be carried out with a high resolution spectrometer at an $8-10$ m telescope in a relatively small amount of time.  L77 has a large \\qmin{} of $\\sim0.61$ and we expected to be sensitive to its secondary.  However, our single observation of L77 has S/N$\\sim$25, and occurred at an orbital phase when the velocity separation was small.  The orbit of this system is such that even had we detected the secondary, the resulting SB2 orbit would be largely unconstrained.  L77 will be at a more favorable orbital phase in 2009.  L90 has $\\qmin\\sim0.32$; we observed it on two occasions, each with S/N$\\sim$50.  However, its long period and current orbital phase indicate a small velocity separation.  This system will not be at a more favorable orbital phase until $\\sim$2011.  We did not detect a companion in our observations of vB~43 on 2005 Sep 30, L57 on 2005 Nov 25 and 2006 Feb 2, and H509 on 2005 Oct 2.  Each of these observations were carried out at orbital phases corresponding to primary velocities near \\g. The observation of vB~43 was compromised further by a low S/N, relative to our other observations of it.  L57 has small \\Kone{} and \\Ktwo{}, which when combined with it's K2 primary spectral type make it a difficult target to observe with CSHELL: future observations would benefit from higher spectral resolution or a larger free spectral range.  \\subsection{\\sc vB~142\\label{vb142}}  Our four infrared observations of vB~142 provided good phase coverage, and included measurements near the maximum velocity separation and on both sides of the \\g-velocity.  Despite this, the SB2 solution that we derive is poorly determined and the residuals, $V_2(fit)-V_2(measured)$, are large. We do not fully understand the reasons behind this poor solution, but we include it here for the following reasons.  Our observations, except for that on 2005 November 28, return plausible velocities for the companion, albeit with large uncertainties.  The 2005 November 28 observation occurred at phase $\\sim0.40$, and with small velocity separation, so our inability to accurately measure the secondary in this spectrum is not surprising. The average $\\alpha$ that we measure for all of the observations, $\\sim0.04$, is, however, much larger than that predicted by BCAH (Figure~\\ref{fig:qvsalpha}).  To investigate this discrepancy we used our M-type template LHS2351 to introduce an additional spectrum into our observed vB~142 spectra from 2004 October 1 and 2005 February 22. We added in this component with a ``true'' flux ratio, $\\alpha_{true}$, ranging from 0.05 to 0.0001, and, to avoid confusion with the actual vB~142 companion, at a radial velocity of $-20\\kps$. We then attempted to recover this signal with our correlation procedure and the set of templates used in the original vB~142 analyses, excluding LHS2351.  We recovered the LHS2351 spectrum at the proper velocity with $\\alpha_{true}$ as small as 0.0005.  However, the uncertainty of the measured velocity increased as $\\alpha_{true}$ decreased, and the measured flux ratio, $\\alpha_{meas.}$ became unreliable for $\\alpha_{true}\\la0.01$.  Consider, for example, the case of LHS2351 introduced into the 2005 February 22 spectrum with a radial velocity of -20\\kps{} and $\\alpha_{true}=0.005$.  Our correlation routine recovered this signal with a velocity of $-25.1\\pm7.0\\kps$ and $\\alpha_{meas.}=0.019\\pm0.005$.  The larger than expected $\\alpha_{meas.}$ is consistent with the results obtained for the vB~142 companion, and may arise because our primary templates do not precisely match the vB~142 primary spectrum.  For small flux ratios, the correlation routine tries to correct for this mismatch by scaling the primary using $\\alpha$; we have previously reported this behavior \\citep{bender2005}.  Our modeling with LHS2351 gives us confidence that we can detect a companion with a very small flux ratio.  The \\q{} that we derive for vB~142, $0.06\\pm0.04$, is consistent with $\\qmin\\sim0.04$ from the SB1 solution.  Whatever the true value of \\q{} may be, our $3\\sigma$ upper limit of $\\q\\la0.18$ is very small.  Our primary goal in this endeavor is to measure the binary mass ratio distribution for the Hyades (\\S1): \\q{} for vB~142 is sufficiently well determined for this purpose.  Of additional interest, the primary of vB~142 is a G5 star, and so a companion with $\\q=0.06\\pm0.04$ could be a brown dwarf, which would be an important discovery in the Hyades \\citep{guenther2005}.  \\subsection{\\sc Component Masses\\label{componentmasses}}  While spectroscopic observations of an SB2 yield its dynamical mass ratio, they do not measure its orbital inclination and so alone they cannot provide a dynamical measurement of the individual component masses.  Observations of the visual orbit measure the inclination and the total mass, and when combined with the mass ratio result in the individual masses. All of the binaries in our sample have components with a small angular separation and their visual orbits are not currently available.  Those with periods longer than a few hundred days are resolvable with adaptive optics imaging at a large aperture telescope.  In the absence of visual orbits, we can still obtain good estimates of the individual masses if the distance is known.  \\emph{Hipparcos} measured the parallax of most of our sample binaries \\citep{perryman1998}, and all have precise photometric measurements of their total flux from 2MASS at J, H, and K.  We combined these with our measured mass ratios and a theoretical mass-luminosity isochrone from BCAH to calculate the individual component masses.  We chose the BCAH models for several reasons: they show a good, albeit not perfect, agreement with measured dynamical masses \\citep{hillenbrand2004}; they include the effects of atmospheres, which are important in the low mass regime that applies to most of our secondaries; and lastly, they are provided in a convenient form that specifies magnitudes in the standard photometric bands used by observers.  We used the 625 Myr isochrone, while noting that at such an old age the mass-luminosity relationship is mostly insensitive to age. Table~\\ref{table:componentmasses} lists the calculated component masses.  The uncertainties given in Table~\\ref{table:componentmasses} include contributions from the parallax, photometry, and mass ratio; they do not include any uncertainties from the BCAH models.  All of the 2MASS J, H, and K photometric uncertainties are small, $\\sim0.02-0.03$ mag. Because our binaries have small flux ratios, the precision with which we determine the primary masses is mostly dependent on the precision of the parallax measurements.  The uncertainty for the secondaries strongly depends on the precision of the mass ratios.  Most of the primaries have masses determined to better than 10\\%, while some of the secondaries approach this level.  \\emph{Hipparcos} did not measure the parallax of vB~59 or L57, so for these systems we used values reported in the Tycho catalog, and the resulting uncertainties on both the primary and secondary masses are large.  The parallaxes of H69, H441, and L79 have not been measured, so we estimated their primary masses directly from their spectral type and assumed an uncertainty of $\\sim0.1\\msun$.  The secondary masses then follow directly from our measured mass ratios.  Finally, the primary of vB~68 is more massive than the range covered by BCAH, so for this system only we used the empirical isochrone determined by \\citet{pinsonneault2004} and note that the resulting masses are consistent with measured spectral types."
    },
    "0808/0808.2867_arXiv.txt": {
        "abstract": "We re-examine the evidence of hemispherical power asymmetry, detected in the cosmic microwave background (CMB) WMAP (Wilkinson Microwave Anisotropy Probe) data using a new method. We use a data filtering, preprocessing , and a statistical approach different from those used previously, and pursue an independent method of parameter estimation. First, we analyze the hemispherical variance ratios and compare these with  simulated distributions. Secondly, working within a previously proposed CMB bipolar modulation model, we constrain the model parameters: the amplitude and the orientation of the modulation field as a function of various multipole bins. Finally, we select three ranges of multipoles leading to the most anomalous signals, and we process corresponding 100 Gaussian, random field (GRF) simulations, treated as observational data, to further test the statistical significance and robustness of the hemispherical power asymmetry. For our analysis we use the Internally-Linearly-Coadded (ILC) full sky map, and the KQ75 cut sky V channel foregrounds reduced map of the WMAP five year data (V5). We constrain the modulation parameters using a generic maximum a \\emph{posteriori} method.  In particular, we find differences in hemispherical power distribution, which when described in terms of a model with bipolar modulation field, exclude the field amplitude value of the isotropic model $A=0$ at confidence level of $\\sim 99.5\\%$ ( $\\sim 99.4\\%$) in the multipole range $\\ell\\in[7,19]$ ($\\ell\\in[7,79]$) in the V5 data, and at the confidence level $\\sim 99.9\\%$ in the multipole range $\\ell\\in[7,39]$ in the ILC5 data, with the best fit (modal PDF) values in these particular multipole ranges of $A=0.21$ ($A=0.21$) and $A=0.15$ respectively.  However, we also point out that similar or larger significances (in terms of rejecting the isotropic model), and large best-fit modulation amplitudes are obtained in GRF simulations as well, which reduces the overall significance of the CMB power asymmetry down to only about $94\\%$ ($95\\%$) in the V5 data, in the range $\\ell\\in[7,19]$ ($\\ell\\in[7,79]$). ",
        "introduction": "The Gaussianity and the statistical isotropy of the fluctuations in the Cosmic Microwave Backgrounds Radiation (CMBR) are two generic features of current standard cosmological model and are compatible with the simplest inflationary scenarios. These predictions have been extensively studied in number of works. An incomplete list includes: \\cite{komatsu-2003-148,2008arXiv0803.0547K,2006MNRAS.371L..50M,2004ApJ...609...22V,wiaux-2006-96,2004ApJ...613...51M, 2004MNRAS.349..973S,2007arXiv0712.1118N,2004PhRvD..69f3007C,2006MNRAS.tmp..497C,2005MNRAS.358..684C,2006astro.ph..6394C, 2008arXiv0804.0136C,2007A&A...474...23C,2007NewAR..51..250D,2008arXiv0807.2687A,2006astro.ph..7577S,2006PhRvD..74l3521H,2006PhRvD..74l3521H,2008MNRAS.385.1718S,2004MNRAS.354..641H,2004ApJ...607L..67H,PASHtestbyBernui,2007IJMPD..16..411B,2005PhRvD..72f3512N,2003ApJ...590L..65C,2006ApJ...647L..87C,2003MNRAS.346...47G,WMAPmultipol,2006MNRAS.367...79C,2006astro.ph..5135C,2005ApJ...635..750B,2006PhRvD..74f3506A,2006PhRvD..74b3005D,2005PhRvL..95g1301L,2007MNRAS.378..153L,2005ApJ...629L...1J,EfstNoProb03a,2004ApJ...605...14E,2007ApJ...660L..81E,2008ApJ...672L..87E,2005PhRvL..95g1301L,2006astro.ph..8129P,2002MNRAS.331..865S,2001PhRvL..87y1303W,2004MNRAS.349..313P,LR08dodec,Lew2008,2005PhRvD..71d3002D,2008PhRvL.100r1301Y} and references therein. Within the theory of inflation the primordial fluctuations are expected to form a Gaussian Random Field (GRF) at the leading order in perturbation theory. Their statistical properties are imprinted in the CMB fluctuations, providing an interesting window on the processes of the early Universe. Although the instrumental effects, like non-Gaussian and non-isotropic noise, or eccentric beams, and astrophysical foregrounds effects, like Galactic, and extra-Galactic point sources, and extended sources of emission, are either well controlled, or corrected for, or masked out, a set of an unexpected anomalies of various magnitudes and at various scales have been detected in the current CMB data \\citep{2007ApJ...655...11C,2007IJMPD..16..411B,2005MNRAS.357..994L,WMAPmultipol,2006MNRAS.367...79C,2006PhRvD..74f3506A,2007MNRAS.378..153L,2004ApJ...605...14E,2005PhRvL..95g1301L,2005PhRvD..72j1302L,2006PhRvD..74h3509C,2008arXiv0804.2387D} (see also \\cite{2006NewAR..50..868H,2008arXiv0805.4157M,2008arXiv0807.1816C} for recent reviews and references therein). These anomalies call for plausible theoretical explanations since, if robustly detected, these can be used as valuable observables of the physics of the early Universe, or a new window on some of the late time effects \\citep{2007ApJ...664..650I,2008arXiv0806.0377E,2008MNRAS.tmp..929B,2007arXiv0710.5897A,2008arXiv0804.2387D,2005PhRvD..72f3002B,2008PhRvD..77f3008D}.   In this paper we re-investigate the well-known, hemispherical power asymmetry observed in the CMB maps. We revisit the properties of the asymmetry, constrain parameters of the previously proposed bipolar modulation field model \\citep{2005PhRvD..72j3002G} responsible for generation of the asymmetry, and we estimate the statistical significance using a generic maximum likelihood method, and realistic Monte-Carlo simulations. Given that we introduce and utilize a different method, from those previously used, and rely on different assumptions, while relax some other, our results can serve as a separate cross consistency and stability check. We provide a through tests of the method so as to validate the presented results.  For the first time we estimate the parameters of the hemispherical modulation for selected ranges of multipoles. Finally, we assess the significance of the hemispherical power asymmetry via direct comparison to the GRF CMB simulations.  The main differences from, and extensions to the previous analyses are: (i) we rely on the local  real-space measurements of the variance as an estimator for the power asymmetry, (ii) we do not assume any priors on the probability distribution function (PDF) for any of the modulation parameters, and explore the likelihood function in the full (albeit sparse) parameter space and apply interpolation, (iii) we fully include the effects of the cut-sky, cross-multipoles power leakage and (iv) we analyze the power asymmetry in the various slices through the spherical harmonic space, pre-filtering the data prior to the analysis, rather than considering all scales scrambled together.  Since the full exploration of the parameter space is CPU expensive, our analysis is based on a few assumptions that greatly simplify, and speed-up the parameter estimation process. Using this different approach, while providing tests and justification for the assumptions made, we give a new estimates on the significance of the hemispherical power asymmetry anomaly.  We also discuss limitations of usage of the method with regard to the extent to which the assumptions of the method remain acceptably valid.  The organization of the paper is as follows: In Section ~\\ref{sec:data_and_sims} we describe our datasets and CMB simulations. In Section ~\\ref{sec:power_ratio} we present results of a statistics that measure the hemispherical power ratios. In Section ~\\ref{sec:modulation_parameter_estimation} we focus on the properties of the power modulation model, our assumptions and tests of the assumptions, and then detail on our method for modulation parameter estimation. In Section ~\\ref{sec:tests} we present the results of various tests of the method. Results of the application of the method to the real CMB data are presented in section ~\\ref{sec:results}. Discussion and conclusions are given in sections~\\ref{sec:discussion} and~\\ref{sec:conclusions} respectively. ",
        "conclusions": "\\label{sec:conclusions}  We performed tests of the hemispherical power asymmetry found in the CMB WMAP data for different bins of multipoles in two ways.  First, we introduced a statistic that searches for the orientation of two opposing, hemispherical regions that maximize or minimize the hemispherical variance ratio, and compared these with the expectations from the GRF simulations. We found that the maximal asymmetry revealed this way is found within a multipole range $\\ell \\in [8,15]$, with the southern hemisphere having larger variance than the northern hemisphere. When these results are compared to the GRF simulations, the northern hemisphere appears to be suppressed below the average expectation.  Secondly, we have introduced and tested a new method for measuring the power asymmetry in the CMB data, as quantified within a bipolar modulation model \\citep{2005PhRvD..72j3002G}. For the first time we constrained the modulation parameters as a function of various multipole  bins. For each multipole range, we obtained the constraints on the the modulation amplitude and orientation. On the basis of the data sets, analyzed up to the maximal multipole $\\lmax=80$, i.e. the WMAP five-year inverse noise co-added, KQ75 sky cut map from the V channel (V5), and the five-year, full-sky, foregrounds cut ILC map (ILC5) we found that: \\\\ (i) generally the modulation amplitude decays as higher multipole bins are cumulatively added or independently analyzed,\\\\ (ii) the best fit modulation amplitude is small $A<0.03$ and insignificant for multipoles beyond $\\ell\\approx 40$\\\\ (iii) the most anomalous signals in terms of the modulation amplitude and its significance come from multipole range $\\ell\\in[7,19]$, and  $\\ell\\in[7,39]$ in the V5 and ILC5 data respectively. For these ranges the significances of rejecting the isotropic cosmological model are $99.5\\%$ and $99.9\\%$ respectively and the constraints on the best fit, (PDF modal) modulation values are: $(0.07) 0.14 < 0.21 < 0.26 (0.031)$ and $(0.06) 0.10 < 0.13 < 0.17 (0.20)$ at $68\\%$ ($95\\%$) CL respectively.  Focusing on the two selected multipole ranges we performed an additional tests of the significance using GRF simulations processed as data, and found that similar or stronger and more significant (in terms of rejecting the isotropic model) modulation values are obtained in 6 (1) cases in 100 simulations, which decreases the overall significance of the power asymmetry in the CMB down to 94\\% (99\\%) in V5 (ILC5) data respectively. To complement the results in the limit of high multipoles as well, we additionally tested the range $\\ell\\in[7,79]$ of the V5 data that also yields a strong and significant (99.4\\%) best-fit modulation value - $(0.02) 0.05 < 0.08 < 0.11 (0.13)$ at $68\\%$ ($95\\%$) CL - but when this result was compared with the GRF simulations the effective significance is again decreased down to about 95\\%.  Although the significance in case of the ILC5 data is still rather high, we warn that the results in this case were obtained without any sky cut, and therefore the asymmetry significance can be overestimated due to residual foregrounds.  Finally we note that a further analysis of the significance in terms of comparison with GRF simulations of other multipole ranges would be interesting, as well as analysis of the  power asymmetry in the ILC data as a function of different sky cuts."
    },
    "0808/0808.0609_arXiv.txt": {
        "abstract": "Large Scale Shocks are responsible for the heating of the ICM and can be important sources of Cosmic Rays (CR) in the Universe. However the occurrence and properties of these shocks are still poorly constrained from both a theoretical and an observational side.  In this work we analyze the properties of Large Scale Shocks in a $(103 {\\rm Mpc/h})^{3}$ cosmological volume simulated with the public 1.0.1 release of the ENZO code. Different methods to identify and characterize shocks in post processing are discussed together with their uncertainties. Re-ionization affects the properties of shocks in simulations, and we propose a fitting procedure to model accurately the effect of re-ionization in non--radiative simulations, with a post--processing procedure. We investigate the properties of shocks in our simulations by means of a procedure which uses jumps in the velocity variables across the cells in the simulations and this allows us to have a viable description of shocks also in relatively under-dense cosmic regions. In particular we derive the distributions of the number of shocks and of the energy dissipated at these shocks as a function of their Mach number, and discuss the evolution of these distributions with cosmological time and across different cosmic environments (clusters, outskirts, filaments, voids).  In line with previous numerical studies (e.g. Ryu et al.2003, Pfrommer et al.2006) relatively weak shocks are found to dominate the process of energy dissipation in the simulated cosmic volume, although we find a larger ratio between weak and strong shocks with respect to previous studies. The bulk of energy is dissipated at shocks with Mach number $M \\approx 2$ and the fraction of strong shocks decreases with increasing the density of the cosmic environments, in agreement with semi-analytical studies in the case of galaxy clusters.  We estimate the rate of injection of CR at Large Scale Shocks by adopting injection efficiencies taken from previous numerical calculations. The bulk of the energy is dissipated in galaxy clusters and in filaments and the flux dissipated in the form of CR within the whole simulated volume at the present epoch is $\\approx 0.2$ of the thermal energy dissipated at shocks; this fraction is smaller within galaxy clusters.  Finally we discuss the properties of shocks in galaxy clusters in relation with their dynamical state. In these regions the bulk of the energy is dissipated at weak shocks, with Mach number $\\approx 1.5$, although slightly stronger shocks are found in the external regions of merging clusters. ",
        "introduction": "\\label{sec:intro}  \\bigskip  Galaxy clusters store up to several $10^{63}$ ergs in the form of hot baryonic matter, with typical temperatures of several keV. This thermal energy is mostly the by-product of shock-heating processes occurred during the formation of cosmic structures. However detecting shocks in Large Scale Structures (LSS) is still observationally challenging since they usually develop in the external regions of galaxy clusters, where the X--ray emission is faint. In a few cases, however, internal shocks driven by the merging events have been discovered with typical Mach numbers $\\approx 1.5 - 3$ (e.g. Markevitch et al. 1999; Markevitch et al. 2002; Belsole et al.2004; Markevitch \\& Vikhlinin 2007). Studies concerning the convolution of simulated cosmological shocks into simulated X--ray images (e.g. Gardini et al. 2004; Rasia et al. 2006) have clearly shown that the apparent lack of shocks in the ICM may be due to the concourse of projections and mass--weighting effects along the line of sight, which tend to bias towards lower values of the overall temperature of the ICM and to smooth temperature gradients.  Shocks are important not only to understand the heating of the ICM but also because they may be efficient accelerators of supra--thermal particles (e.g. Sarazin 1999; Takizawa \\& Naito 2000; Blasi 2001). Non thermal activity in galaxy clusters is proved by radio observations which show synchrotron emission in a fraction of merging clusters: Radio Halos , at the cluster center, and Radio Relics, elongated sources at the cluster periphery (e.g., Feretti 2005). Several mechanisms related to cluster mergers and to the accretion of matter can act as sources of non thermal components in galaxy clusters. Large scale radio emission in the form of giant Radio Halos may be powered by particle re-acceleration by MHD turbulence injected in the ICM during energetic merging events (Brunetti et al. 2001; Petrosian 2001; Brunetti et al. 2004,08). Strong shocks may accelerate supra-- thermal electrons from the thermal pool and explain the origin of Radio Relics (Ensslin et al.1998), while high energy electrons accelerated at these shocks can produce X-rays and gamma-rays via Inverse Compton scattering off CMB photons (e.g. Sarazin 1999; Loeb \\& Waxmann 2000; Blasi 2001; Miniati 2003). Relativistic hadrons accelerated at shocks can be advected in galaxy clusters and efficiently accumulated (V{\\\"o}lk, Aharonian, \\& Breitschwerdt 1996; Berezinsky, Blasi, \\& Ptuskin 1997) producing an important non-thermal component which could be directly sampled by future gamma ray observations (e.g., Blasi, Gabici \\& Brunetti 2007). Secondary particles injected in the ICM via proton--proton collisions may also produce detectable synchrotron radiation (e.g. Blasi \\& Colafrancesco 1999; Dolag \\& Enssil 2000) and may be eventually re-accelerated by MHD turbulence yielding an efficient picture to explain Radio Haloes (Brunetti \\& Blasi 2005).  The energetics associated with the population of cosmic ray particles (CR) accelerated at shocks depends on the Mach number of these shocks (e.g. Kang \\& Jones 2002). The Mach number distribution of cosmological shocks is thus important to understand CR in galaxy clusters. Semi--analytical studies pointed out that shocks that form during cluster mergers are weak, $M \\sim 1.5$, being driven by sub-clumps crossing the main clusters at the free-fall velocity (Gabici \\& Blasi 2003, Berrington \\& Dermer 2003). These approaches however are limited as they treat cluster mergers as binary encounters between ideally virialised spherical systems. Therefore cosmological numerical simulations represent a necessary avenue to address this issue in more detail. First attempts to characterize shock waves in cosmological numerical simulations were produced by Miniati et al.(2000), by employing a set of eulerian simulations and a shock detecting scheme based on jumps in the temperature variable. Later works adopted more refined shock-detecting schemes and were more focused onto the distribution of energy dissipated at shocks (Keshet et al.2003; Ryu et al.2003, Hallman et al.2003, Pfrommer et al.2006{\\footnote{After this paper was submitted, also Skillman et al.2008 provided results on shock waves in ENZO AMR simulations, which are broadly consistent with the above works in literature.}}). Ryu et al. and Pfrommer et al. basically found that the bulk of shocks in the universe is made of relatively weak shocks, but also allow to constrain the population of stronger shocks forming in the external regions of galaxy clusters, were structures are not completely virialised. In these environments, strong shocks are frequent and may provide the bulk of the acceleration of CR in large scale structures (Ryu et al. 2003; Pfrommer et al. 2006). For this reason numerical simulations are important to study the non-thermal components in galaxy clusters (e.g. Miniati et al.2001, Miniati 2003; Keshet et al.2003; Pfrommer 2008). However, the identification and characterization of shocks, as well as the calculation of the energy injected in the form of CR at these shocks, is difficult because of the complex dynamics of large scale structures and because of the severe limitations in terms of physics and numerical resolution that affect present cosmological simulations.  In this paper we follow the approach of the seminal paper by Ryu et al. (2003), studying the shock wave patterns in LSS and characterizing them in a post--processing phase. This allows us to briefly discuss the effect of different shock detection techniques and their dependence on the variation of underlying physical processes (e.g. re-ionization) in the simulation. We use the cosmological code ENZO (Bryan \\& Norman 1997), which treats gas dynamics with an Eulerian scheme and allows us to follow the assembly of LSS with sufficiently good spatial resolution.  The outline of the paper is the following. In Sect.\\ref{sec:enzo} we provide a brief introduction to the ENZO Code, in Sect.\\ref{sec:sim} we present our cluster sample and the main properties of cosmological structures in our simulations, and in Sect.\\ref{sec:reionization} we discuss the effect of re-ionization on the thermal properties of simulated cosmic structures. In Sect.\\ref{sec:algo} we provide the different methods to characterize shocks in post processing and in Sect.\\ref{sec:comparison} we discuss their main source of uncertainties in the cosmological framework. In Sect.\\ref{sec:results} we present our results about the main shocks properties and about the injection of CR. The main conclusions of this work are given in Sect.\\ref{sec:conclusions}. In the Appendix we the effect of the resolution and of $\\sigma_{8}$ on our results.  \\begin{table} \\label{tab:tab0} \\caption{Main characteristics of the simulations.} \\centering \\tabcolsep 4pt \\begin{tabular}{c|c|c|c} Volume &  Resolution &  physics & ID\\\\ \\hline $(145Mpc)^{3}$ & $125kpc$ & adiab. & {\\bf AD125}\\\\ $(145Mpc)^{3}$ & $250kpc$ & adiab. & {\\bf AD250}\\\\ $(145Mpc)^{3}$ & $500kpc$ & adiab. & {\\bf AD500}\\\\ $(145Mpc)^{3}$ & $800kpc$ & adiab. & {\\bf AD800}\\\\ $(80Mpc)^{3}$ & $125kpc$  & cool. + reion. & {\\bf CO125}\\\\ $(80Mpc)^{3}$ & $250kpc$ & cool., reion. and $\\sigma8=0.74$ & {\\bf S8250}\\\\ $(80Mpc)^{3}$ & $250kpc$ & cool.+reion. & {\\bf CO250}\\\\ \\end{tabular} \\end{table} ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have investigated the properties of shocks in cosmological numerical simulations. This subject is particularly intriguing as shock waves propagating trough LSS are the responsible for the heating of the ICM and may be important sources of CR in the Universe. This subject has been already investigated in several papers under different numerical approaches (Miniati et al.2000; Miniati et al.2001; Keshet et al.2003; Ryu et al.2003; Pfrommer et al.2006; Kang et al.2007; Pfrommer et al.2007{\\footnote{See also Skillman et al.2008, which appeared after this paper was submitted.}}).   We study shock waves by means of a post processing procedure. Although this is similar to Ryu et al.(2003) and Kang et al.(2007), our approach differs from previous ones in several points. First we use a different numerical codes, the public version of ENZO (e.g. Bryan \\& Norman 1997), to simulate LSS (Sec.\\ref{sec:enzo}). Second we adopt a more appropriate treatment of the re-ionization in our simulations (Sec.\\ref{sec:reionization}) and third we use a different approach to catch shocks in our simulations and to measure their strength (VJ method, Sect.\\ref{subsec:vj}).   \\subsection{Results}  We simulated a large cosmic volume, $(145 {\\rm Mpc})^{3} \\approx (103 {\\rm Mpc/h})^{3}$, with a fixed grid resolution of 125 kpc. Additional simulations were designed and used to investigate the effect of spatial resolutions and of the $\\sigma_{8}$ parameter (see Appendix).  \\noindent In the following we summarize the main results obtained :  \\begin{itemize}  \\item {\\it Re-ionization}: in Sect.\\ref{subsec:reion} we have shown that a correct treatment of the re-ionization is crucial to have a viable description of shocks with a post processing procedure. In particular using a constant temperature floor in the simulated data at a fixed redshift may cause a somewhat artificial flattening (or pile up) of the differential distribution of the shocks with Mach number. This is because the velocity of the sound is artificially boosted up in low density and cold regions. To overcome this point we derive formulas giving the typical temperature of the gas as a function of the local density by fitting data obtained from simulations which include a specific modeling of the re-ionization in run--time. These formulas are found to be consistent with Katz et al.(1996) and Valageas et al.(2002) and can be used to model the temperature background of adiabatic simulations in a post processing procedure. In Section 4 (Fig.\\ref{fig:madau-vs-post}) we have shown that our post processing procedure is indeed consistent with the results from simulations with run--time re-ionization.   \\item {\\it Methods to derive the Mach number of shocks}: in Sect.\\ref{subsec:Ryu0} and \\ref{subsec:vj}  we have discussed in some detail two different methods to catch shocks in simulated data and to estimate their Mach number. These methods are the temperature jump (TJ) and the velocity jump (VJ) and rely with the jumps in temperature and velocity across shocked cells, respectively.  The shock discontinuity is typically spread over a few cells and the risk in measuring the Mach number of shocks trough cell-to-cell velocity or temperature jumps (or jumps across a few cells) is to underestimate the Mach number of shocks. To study this point in the context of our numerical simulations we perform several shock--tube tests with ENZO and calculate maps of shocks in our simulations under different approaches (Sect.~6.1). We conclude that shocks in our simulations are best characterized from velocity (VJ) or temperature (TJ) jumps measured across three cells.   Both the VJ and TJ schemes use ideal conditions across non shocked cells (i.e. no velocity and no temperature gradients) and this may cause major uncertainties in the characterization of the shocks in a post processing procedure (Sect.\\ref{subsec:tjun}--\\ref{subsec:vjun}). This is because the velocity field and temperature distributions in the cosmological data sets are very complex and the passage of a shock establishes thermodynamical gradients which are superimposed to already existing ones. We discuss the strength of the uncertainties on the value of the Mach number from the two schemes by means of Monte Carlo extractions of temperature and velocity variations across non shocked cells in our data sets, and show that the VJ method may be more reliable, at least in the case of weak shocks and especially in low density environments (Figs.8--9).  \\noindent Besides these uncertainties we find that the two methods yield statistically similar Mach number distributions of shocked cells in our simulations (Fig.~12) suggesting that the characterization of shocks in our simulations is fairly solid; in Sect.7 we adopt here the VJ method.  \\item {\\it Morphology of LS shocks}: in Sect.\\ref{subsec:machs} we discuss the morphology of the shock--patterns detected in our simulated data sets. About $15$ per cent of the cells in our simulations are found to host shocks at the present epoch, and this fraction slightly decreases with look back time for post--reionization epochs. In qualitative agreement with previous studies (Ryu et al. 2003; Pfrommer et al. 2006) we find that these shocked cells form spectacular and complex patterns associated with the cosmic web. Filamentary or sheet--like shocks are found outside the virial regions of clusters and around filaments, while more regular spherical structures surround galaxy clusters.   \\item {\\it Number Distributions of LS shocks}: we study the number distribution of shocked cells as a function of the Mach number of the shocks. An important point here is that thanks to the Eulerian scheme of the ENZO code we were able to follow the hydrodynamics of the LS shocks also in very low density regions, whose exploration is challenging for present Lagrangian schemes.  We find that the bulk of cosmological shocks is essentially made by weak $M \\leq 2$ shocks and that the number distribution of shocks can be grossly described by a steep power law $M dN/dM \\propto M^{\\alpha}$. When considering the Mach number distribution of shocked cells in the total simulated volume we find an overall steep distribution $\\alpha \\approx -1.6 $ which is dominated by the contribution from voids and filaments. This distribution steepens with increasing the cosmic over-density and becomes very steep ($\\alpha \\approx -3$ to $-4$) in the case of galaxy clusters demonstrating the increasing rarity of strong shocks in these denser (and hotter) regions, where basically $M<3$.   \\item {\\it Energy dissipated at LS shocks} : the energy dissipation at LS shocks is the main focus of the previous studies on this topic (e.g., Miniati et al.2001; Keshet et al.2003; Ryu et al.2003; Pfrommer et al. 2006; Kang et al.2007; Pfrommer et al.2007). Following Ryu et al.(2003) we calculate the energy rate dissipated in the form of thermal energy at shocks by means of hydrodynamical jump conditions (Eq.\\ref{eq:eq_delta}). In agreement with these previous studies we find that about $\\sim 4 \\cdot 10^{47} erg/s$ are dissipated at shocks in a $(103 Mpc/h)^3$ volume in the simulations at the present epoch. The bulk of the energy in our simulations is dissipated in galaxy clusters which indeed contribute to about $75$ per cent of the total energy dissipation (about $80$ per cent if we also include the contribute from cluster outskirts), while filaments contribute to a $15$ per cent of the total energy dissipation. We calculate the distribution of the energy flux dissipated at LS shocks with shock-Mach number. When considering shocked cells in the total volume we find that the distribution is steep, $\\alpha_{th} \\approx -2.7$ ($f_{th}(M)M \\propto M^{\\alpha_{th}}$) and peaks at $M \\approx 2$. Although in qualitative agreement with previous studies, our distribution is steeper than that obtained by Ryu et al.(2003) that also used cosmological simulations based on a Eulerian scheme. This difference is mostly due to our more complex treatment of the re-ionization background and to the use of the VJ scheme to measure the Mach number of shocks.  Following Ryu et al.(2003) we calculate the efficiency of the injection of CR at LS shocks according to Kang \\& Jones (2002). We obtained Mach number distributions of the energy flux dissipated into CR acceleration in our simulated data sets. Our results are in line with previous findings, although our distributions are steeper than those in Ryu et al. and slightly steeper than those in Pfrommer et al.(2006 \\& 2007).  In agreement with Pfrommer et al. we find that the bulk of the energy dissipated in the form of CR at shocks is shared between clusters and filaments so that CR--acceleration happens in regions broader than those where thermal energy is dissipated at shocks and the ratio between CR energy and thermal energy increases in lower density regions (Fig.\\ref{fig:flux1}). When considering all the shocked cells in our simulations we find that the ratio between the energy dissipated in the form of CR-acceleration and of thermal energy at present epoch is $f_{CR}/f_{th} \\approx 0.2$ and this ratio becomes smaller in galaxy clusters. These fractions are a factor $\\approx$2 smaller than those found by Ryu et al. and this is likely due to the steeper distributions that we obtained in our analysis.  \\item {\\it Galaxy Clusters}: in Sect.\\ref{subsec:clusters} we briefly discuss the case of shocks propagating in galaxy clusters. We find very steep distributions of the number of shocks and of the energy dissipated at shocks as a function of the shock-Mach number. The typical Mach number within the virial radius is $M \\approx 1.5$, in nice agreement with semi--analytical studies which are indeed appropriate for virialised systems (Gabici \\& Blasi 2003). At larger distance from the cluster center stronger shocks are found and their presence is strongly correlated with the dynamical status of the clusters. Remarkably the rarity of moderate--strong shocks in the cluster central regions (within $\\approx$ Mpc distance from cluster center) makes the ratio $f_{CR}/f_{th}$ very small, expecially when the Kang \\& Jones (2007) model for the injection of CR at shocks is adopted (Fig.~25). \\end{itemize}  \\subsection{Comment on the injection of CR}  The use of numerical simulations is mandatory to understand the properties of LS shocks and to investigate their role in the injection of CR in LSS. Although we use a different approach with respect to other studies, our findings for the energy dissipated in the form of CR at these shocks are grossly consistent with previous studies (Ryu et al. 2003; Pfrommer et al. 2006).  However, the astrophysical problem is extremely complex and several {\\it hidden} ingredients in the adopted procedures are potentially sources of large uncertainties. As discussed in Sect.\\ref{subsec:CR} the efficiency of CR acceleration at shocks is investigated following several approaches. In the present paper we have used the acceleration efficiency resulting from numerical calculations of modified shocks (following Ryu et al. 2003 and references therein). On the other hand Pfrommer et al.(2006) use a linear theory with the efficiency modified to account for saturation effects at large values of the Mach numbers (actually to limit the CR efficiency at $\\approx 50$ per cent). These two approaches are formally {\\it radically} different, but nevertheless they provide an overall estimate of the CR injection efficiency which is not dramatically different in the case of typical shocks with $M \\approx 2-4$. The main {\\it hidden} ingredient in the efficiency of CR acceleration comes from the commonly adopted {\\it thermal leakage injection} scenario which essentially adopts as minimum momentum of the particles that take part in the acceleration process, $p_{inj}$, a multiple of the momentum of the thermal particles, $p_{inj} = x_{i} p_{th}$. The choice of $x_{i}$ is a {\\it guess}, since this depends on physical details which are still poorly known (e.g., Blasi 2004). In Ryu et al.(2003) (and thus in our paper) the fraction of protons injected into the CR population at shocks is taken of the order of $\\approx 10^{-3}$ which is not far from (even if larger than) the resulting efficiency from the assumption of $p_{inj} \\approx 3.5 p_{th}$ adopted in Pfrommer et al.(2006). Although this parameter is somewhat constrained by the theory (e.g., Malkov 1998), it should be stressed that having a slightly different value of $x_{i}$ (e.g. $x_{i}$=3.8 instead of 3.5) would have the net effect to reduce the acceleration efficiency by nearly one order of magnitude.   \\subsection{Constraints from observations}   Theoretical arguments suggest that the bulk of CR in galaxy clusters should be in the form of supra--thermal protons (e.g. Blasi, Gabici, Brunetti 2007 for a recent review). Gamma ray observations of a few nearby galaxy clusters limit the energy density of CR protons in these clusters to $\\approx$ 20 per cent of the thermal energy (Pfrommer \\& Ensslin 2004; Reimer 2004). Although a detailed comparison with these limits would require to follow the advection and accumulation process of CR in clusters during the simulations (e.g. Miniati 2003; Pfrommer et al. 2006), the level of injection rate of CR that we found in clusters and cluster outskirts (Sect.\\ref{subsec:clusters}) allows us to reasonably conclude that our results are barely consistent with these limits.  On the other hand more stringent upper limits can be obtained from present radio observations. The bulk of galaxy clusters does not show evidence of extended Mpc--scale synchrotron radio emission and this can be used to constrain the population of secondary electrons and thus that of the primary CR protons from which these secondaries would be injected (Brunetti et al. 2007). These limits are very stringent and actually represent a challenge for simulations: in the case that the ICM is magnetized at $\\approx \\mu$G level (consistent with Rotation Measures, e.g. Govoni \\& Feretti 2004) the energy of CR should be at $\\leq$ few percent of the thermal energy (when the spectrum of CR is fixed at that expected from simulations, i.e. $s\\approx 2 - 2.5$, $N(p)\\propto p^{-s}$ for $M\\geq 3$).  \\noindent A comparison between our simulations and present limits clearly requires a more detailed study. However, a simple estimate of the spectral shape of CR injected in our simulations (Fig.\\ref{fig:elec_distrib}) suggests that the bulk of the CR energy in clusters and cluster outskirts is associated with CR populations with relatively steep spectra, in which case the limits from radio observations become less stringent (Brunetti et al. 2007). Obviously, an alternative possibility to reconcile radio data and simulations would be that relatively flat components of CR store a non negligible fraction of the energy of the ICM, and that the magnetic field strengths in the ICM are much smaller than that claimed from the analysis of present RM data; future gamma ray observations with the FERMI Gamma Ray Telescope and with Cherenkov arrays will clarify this possibility.   \\begin{figure*} \\includegraphics[width=0.95\\textwidth,height=0.8\\textwidth]{imag/slopes_cr3.ps} \\caption{Energy injection rate of CR for the 4 galaxy clusters presented in the text, as a function of the spectral slope of the energy spectrum of injected protons, $N(p) \\propto p^{-s}$. Estimates are shown for the KJ02 ({\\it left}) and for the KJ07 model ({\\it right}). Here for simplicity we assume DSA at non--modified shocks (linear theory), in which case it is $s = 2 (M^{2}+1)/(M^{2}-1)$ (e.g. Gabici \\& Blasi 2003).} \\label{fig:elec_distrib} \\end{figure*}   \\bigskip    We thanks K. Dolag, D. Ryu, H. Kang, G. Tormen, L. Moscardini, B.O'Shea, M.Norman, R.Wagner, D.Collins, S. Skory, S. Giacintucci, R. Brunino and A.Bonadede for useful discussions and helps. We acknowledge partial support through grants ASI-INAF I/088/06/0 and PRIN-INAF 2007, and the usage of computational resources under the CINECA-INAF agreement."
    }
}